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0710 | 0710.0402.txt | {}{}{}{}{} % 5 {} token are mandatory \abstract % context heading (optional) % {} leave it empty if necessary % {The cosmic X-ray background is the summed radiation from the growth of the massive black holes that reside in the centres % of galaxies. By studying the population of sources that produce the X-ray background through X-ray surveys at various depths, % we can determine when, where and how this growth occurs. } % aims heading (mandatory) {} {X-ray sources at intermediate fluxes (a few $\times 10^{-14}\, {\rm erg\, cm^{-2}\, s^{-1}}$) with sky density of $\sim 100\, {\rm deg}^{-2}$, are responsible for a significant fraction of the cosmic X-ray background at various energies below 10 keV. The aim of this paper is to provide an unbiased and quantitative description of the X-ray source population at these fluxes and in various X-ray energy bands.} % methods heading (mandatory) {We present the XMM-Newton Medium sensitivity Survey (XMS), including a total of 318 X-ray sources found among the serendipitous content of 25 XMM-Newton target fields. The XMS comprises four largely overlapping source samples selected at soft (0.5-2 keV), intermediate (0.5-4.5 keV), hard (2-10 keV) and ultra-hard (4.5-7.5 keV) bands, the first three of them being flux-limited. } % results heading (mandatory) {We report on the optical identification of the XMS samples, complete to 85-95\%. At the flux levels sampled by the XMS we find that the X-ray sky is largely dominated by Active Galactic Nuclei. The fraction of stars in soft X-ray selected samples is below 10\%, and only a few per cent for hard selected samples. We find that the fraction of optically obscured objects in the AGN population stays constant at around 15-20\% for soft and intermediate band selected X-ray sources, over 2 decades of flux. The fraction of obscured objects amongst the AGN population is larger ($\sim 35-45\%$) in the hard or ultra-hard selected samples, and constant across a similarly wide flux range. The distribution in X-ray-to-optical flux ratio is a strong function of the selection band, with a larger fraction of sources with high values in hard selected samples. Sources with X-ray-to-optical flux ratios in excess of 10 are dominated by obscured AGN, but with a significant contribution from unobscured AGN. } % conclusions heading (optional), leave it empty if necessary {} | Supermassive black holes (SMBHs, i.e., with masses $\sim 10^6-10^9\, {\rm M}_{\odot}$) have been detected in the centers of virtually all nearby galaxies \citep{Merrit01,Tremaine02}. In many of these galaxies -including our own-, the SMBH is largely dormant, i.e., the luminosity is many orders of magnitude below the Eddington limit. Only $\sim 10\%$ of today's galaxies (at most) host active galactic nuclei (AGN), and a very large fraction of them are in fact inconspicuous at most wavelengths because of obscuration \citep{Fabian99}. %SMBH %masses correlate well with host galaxy properties (stellar or gas %central velocity dispersion, bulge luminosity), and are estimated %to contain around 0.4-0.6\% of the bulge masses %\citep{Magorrian98}. It is generally believed that the seeds of these SMBHs were the remnants of the first generation of massive stars in the history of the Universe. These early black holes may have had masses of tens of ${\rm M}_{\odot}$ at most. The growth of these relic black holes to their current sizes is very likely dominated by accretion, with additional contributions by other phenomena like black hole mergers and tidal capture of stars \citep{Marconi04}. According to current synthesis models, the integrated X-ray emission produced by the growth of SMBHs by accretion over the history of the Universe is recorded in the X-ray background (XRB). Thus the XRB can be used to constrain the epochs and environments in which SMBHs developed. There are currently a number of existing or on-going surveys in various X-ray energy bands (see \citet{Brandt05} for a recent compilation). In the pre-Chandra and pre-{\it XMM-Newton} era the Einstein Extended Medium Sensitivity Survey \citep{Maccacaro82,Gioia90,Stocke91} pioneered the procedure of determining typical X-ray to optical flux ratios for different classes of X-ray sources to facilitate the identification processes and has set the standards for serendipitous X-ray surveys. $ROSAT$ produced a number of surveys in the soft 0.5-2 keV X-ray band at various depths, e.g., the $ROSAT$ Bright Survey \citep{Schwope00}, the intermediate flux RIXOS survey \citep{Mason00} and the $ROSAT$ deep surveys \citep{McHardy98,Georgantopoulos96,Hasinger98,Lehmann01} among others. These surveys show that AGN dominate the high Galactic latitude soft X-ray sky at virtually all relevant fluxes. The majority of these AGN are of spectroscopic type 1, which means that we are witnessing the growth of SMBH through unobscured lines of sight. In a moderate fraction of the sources identified, however, there is evidence for obscuration as their optical spectra lack broad emission lines (type 2 AGN). \citet{Ueda03} discuss the results from a large area X-ray survey in the 2-10 keV band with ASCA and those from HEAO-1 and $Chandra$, where a larger fraction of the sources identified correspond to type 2 AGN. With {\it Chandra} and {\it XMM-Newton} coming into operation X-ray surveys, particularly at energies above a few keV, have been significantly boosted. Thanks to the high sensitivity and large field of view of the EPIC cameras \citep{Turner01,Struder01} on board {\it XMM-Newton} \citep{Jansen01}, X-ray surveys requiring large solid angles have been dominated by this instrument. The Bright Source Survey-BSS \citep{Dellaceca04} contains 400 sources brighter than $\sim 7\times 10^{-14}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ either in 0.5-4.5 keV or 4.5-7.5 keV. The BSS samples\footnote{{\tt http://www.brera.mi.astro.it/\~{}xmm/}}, which have been identified to $\sim 90\%$ \citep{Caccianiga07}, show an X-ray sky dominated by AGN, where the fraction of obscured objects varies with the selection band (sample selection at harder energies reveals a higher fraction of obscured objects as expected). Deep surveys have also been conducted by {\it XMM-Newton}, for example in the Lockman Hole down to $\sim 10^{-15} \, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ \citep{Hasinger01,Mateos05b}. However, thanks to its much better angular resolution, the {\it Chandra} deep surveys are photon counting limited and far from confusion and are consequently much more competitive at fainter fluxes \citep{Alexander03,Tozzi06}. Optical identification of these deep surveys is largely incomplete, a fact that is driven by the intrinsic faintness and red colour of most of the counterparts to the faintest X-ray sources. In the intermediate flux regime, however, the identified fractions are large and nearing completion. It is interesting to note that deep surveys start to find a population of galaxies not necessarily hosting active nuclei as an important ingredient. In addition, the AGN population is found to contain an important fraction of obscured objects. The wide range of intermediate X-ray fluxes, between say $10^{-15}\, \, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ and $10^{-13}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ have also been the subject of a number of on-going surveys. Besides bridging the gap between wide and deep surveys, intermediate fluxes sample the region around the break in the X-ray source counts \citep{Carrera07}, and therefore their sources are responsible for a large fraction of the X-ray background. Among these, we highlight the {\it XMM-Newton} survey in the well-studied (at many bands) COSMOS field, which covers $2\, \deg^2$ to fluxes $\sim 10^{-15}\, {\rm erg\, cm^{-2}\, s^{-1}}$ \citep{Hasinger07}. The optical identification is still on-going, reaching 40\% \citep{Brusa07}. At fluxes around $10^{-14}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$, the HELLAS2XMM survey \citep{Baldi02,Fiore03}, now extended to cover $1.4\, \deg^2$, contains over 220 X-ray sources, optically identified to 70\% completeness \citep{Cocchia07}. Other surveys in this flux range include the {\it XMM-Newton} survey in the Marano field \citep{Krumpe07}, which is 65\% identified over a modest solid angle of $0.28\, \deg^2$. Also the XMM-2dF survey (Tedds et al., in preparation), which contains almost 1000 X-ray sources optically identified in the Southern Hemisphere, is an important contributor in this regime. {\it Chandra} has also triggered surveys at intermediate fluxes, most notably the {\it Chandra} Multiwavelength Survey \citep{Kim04a,Kim04b,Green04}, covering $1.7\, \deg^2$ and identified to $\sim 40\%$ completeness \citep{Silverman05}. In the realm of this variety of X-ray surveys that yield a qualitative picture of the X-ray sky, the {\it XMM-Newton} Medium sensitivity Survey (XMS) discussed in this paper, finds its role in three important ways: a) it deals with very large samples, selected at various X-ray bands where {\it XMM-Newton} is sensitive, from 0.5 to 10 keV; b) the samples that we consider have been identified almost in full, from 85\% to 95\% completeness and c) three out of the four samples that we explore are strictly flux limited in three energy bands (0.5-2 keV, 0.5-4.5 keV and 2-10 keV). Armed with these unique features, the XMS is a very powerful tool to derive a {\it quantitative} characterization of the population of X-ray sources selected in various bands, and also to study and characterize minority populations, all of it at specific intermediate X-ray fluxes where a substantial fraction of the X-ray background below 10 keV is generated. The power of the XMS is enhanced by the fact that to some extent it is a representative sub-sample of the {\it XMM-Newton} X-ray source catalogue 2XMM\footnote{Pre-release under {\tt http://xmm.vilspa.esa.es/xsa}}, containing 150,000 entries. Specific goals that have driven the construction of the XMS whose results are presented in this paper include: a) quantify the fraction of stars versus extragalactic sources at intermediate X-ray fluxes and at different X-ray energy bands; b) quantify the fraction of AGN that are classified as obscured by optical spectroscopy at intermediate X-ray fluxes and for samples selected in different energy bands; c) find the redshift distribution for the various classes of extragalactic sources and compare soft and hard X-ray selected samples; d) study the distribution of the X-ray-to-optical flux ratio for the various classes of X-ray sources, also as a function of X-ray selection band. The X-ray spectral properties of the sources of the XMS were already discussed in \citet{Mateos05a}. Further goals that we will achieve with the XMS in forthcoming papers include: e) determine the fraction of ``red QSOs'' at intermediate X-ray fluxes and as a function of X-ray selection band; f) relate X-ray spectral properties (like photoelectric absorption) to optical colours of the counterpart; g) quantify the fraction of radio-loud AGN in the samples selected at various X-ray energies; h) construct Spectral Energy Distributions for the various classes of sources in the XMS. Results on these further aspects will be presented in a forthcoming paper (Bussons-Gordo et al., in preparation). The paper is organized as follows: in Section~\ref{sec:XMS} we define the XMS along with the 4 samples that constitute it, including the X-ray source list; in Section~\ref{sec:imaging} we discuss the multi-band optical imaging conducted on the {\it XMM-Newton} target fields and the process for selecting candidate counterparts; this is continued in Section~\ref{sec:identification} where we discuss the identification of the XMS sources in terms of optical spectroscopy, and list photometric and spectroscopic information on each XMS source. Section~\ref{sec:XMSpopulations} presents the first scientific results from the XMS, specifically a description of the overall source populations, the fraction of stars in the various samples, the fraction of optically obscured AGN, and the X-ray to optical flux ratio of the different source populations. Section~\ref{sec:conclusions} summarizes our main results. To clarify the terminology used in this paper, an AGN not displaying broad emission lines in its optical spectrum is termed as type 2 or obscured, and type 1 or unobscured otherwise. The property of being absorbed or unabsorbed refers only to the detection or not of photoelectric X-ray absorption. Throughout this paper, we used a single power law X-ray spectrum to convert from X-ray source count rate to flux in physical units, with a photon spectral index $\Gamma=1.8$ for the XMS-S and XMS-X samples and $\Gamma=1.7$ for the XMS-H and XMS-U samples. These are the average values obtained by \citet{Carrera07}, which -as opposed to what we do here- used the specific value of $\Gamma$ for each individual source and energy range. When computing luminosities, we also use the above spectra for K-correction and the concordance cosmology parameter values: $H_0=70\, {\rm km}\, {\rm s}^{-1}\, {\rm Mpc}^{-1}$, $\Omega_m=0.3$ and $\Omega_\Lambda=0.7$. All quoted uncertainties in parameter estimates are shown at 90\% confidence level for one interesting parameter. | \label{sec:conclusions} In this paper we have presented the {\it XMM-Newton} Medium sensitivity Survey XMS, and extracted a number of robust quantitative conclusions about the population of high Galactic latitude X-ray sources at intermediate flux levels. We have argued that given the completeness of our identifications and the relatively large size of the XMS samples, these conclusions can be safely exported to a much larger X-ray source catalogue like 2XMM. Our conclusions can be summarized as follows: \begin{enumerate} \item The high galactic latitude X-ray sky at intermediate flux levels is dominated by AGN, which includes type-1 and type-2 AGN as well as the so-called XBONG which are likely to host a low luminosity or obscured nucleus (or both). The stellar content is less than 10\% in soft X-ray selected samples, and drops to below 5\% at around soft X-ray fluxes $\sim 10^{-14}\, {\rm erg}\, {\rm cm}\, {\rm s}^{-1}$. The stellar content in hard X-ray selected samples does not exceed a few per cent at most. Selection in 0.5-4.5 keV produces intermediate results. \item Given the limited sensitivity of {\it XMM-Newton} above a few keV -which is due to the roll over of effective area- current surveys conducted in the so-called ultra-hard band (4.5-7.5 keV) do not bring any new source population or any significant difference with respect to 2-10 keV selected surveys. Much longer exposure times would be needed to unveil any new heavily obscured population with {\it XMM-Newton}. \item Obscured AGN represent $\sim 20\%$ of the soft X-ray selected population of AGN, all the way from $\sim 10^{-13}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$ down to $\sim 10^{-15}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$, with no compelling evidence for an increase of this fraction towards fainter fluxes within this range. \item Likewise, obscured AGN represent $\sim 35\%$ ($45\%$ if all unidentified sources are obscured AGN) of the hard X-ray selected population of AGN, with no hint of an increase down to a hard X-ray flux $\sim 10^{-14}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm s}^{-1}$. \item The fraction of X-ray sources with X-ray to optical flux ratio $>10$ (or $X/O>1$ using the notation of this paper) is a mere 3\% in soft X-ray selected samples, but grows to 20\% in hard X-ray selected samples. \item Those sources with $X/O>1$ are mostly obscured AGN, but a fraction of around 20\% of them in the hard band are unobscured type-1 AGN. This means that $X/O>1$ alone cannot be used as a proxy for obscured X-ray sources. \end{enumerate} | 7 | 10 | 0710.0402 |
0710 | 0710.3659_arXiv.txt | }{} \textheight 240mm \textwidth 170mm \hoffset= 0mm \voffset=0cm \topmargin -2cm \oddsidemargin 0mm \evensidemargin 0mm \begin{document} \fontsize{12}{12} \selectfont \title{Luminosity Dependence of the Quasar Clustering from SDSS DR5} \author{Ganna Ivashchenko} \date{} 43024 objects, which were primarily identified as quasars in SDSS DR5 and have spectroscopic redshifts were used to study the luminosity dependence of the quasar clustering with the help of two different techniques. The obtained results reveal that brighter quasars are more clustered, but this dependence is weak, which is in agreement with the results by Porciani \& Norberg \cite{porciani2006} and theoretical predictions by Lidz et al. \cite{lidz}. \textbf{Key words:} quasars, large-scale structure, clustering, luminosity. PACS numbers: 98.65.Aj, 98.54.-h, 98.62.Ve. | \indent \indent Determining the distribution of extragalactic objects is one of the most important problems of modern cosmology, because they are the only tracers of the dark matter, except anisotropy of the CMB, that helps us to realize the matter distribution on the redshifts of about 1000. Unfortunately even in the local Universe it is difficult to see the most faint galaxies, and when we go farther to larger redshifts, the most part of objects we can observe there with the modern ground-based telescopes used for large surveys are quasars as the most luminous objects. The matter distribution that we could reconstruct with the help of quasars is just like a light sketch, but it is all we have for today. The quasars are nonuniform objects: they have various luminosities and are on the different stages of their evolution. Thus the reasonable question arises: how their clustering could depend on their physical properties, for example on their luminosity? According to different numerical simulations of galaxy mergers that incorporate black hole growth, this dependence has to exist because more luminous quasars are considered to be born in denser environment. In the main part of such models, in which the host halo mass correlates with the instantaneous luminosity of the quasars (see e.g. \cite{kaufman}), there should be a strong luminosity dependence of the quasar clustering. In the other type of such models, in which the host halo mass correlates with the peak luminosity of quasars, the luminosity dependence of the quasar clustering should be weaker because all of the quasars we see now, are considered to be similar objects but on different stages of their evolution (see \cite{lidz} and references therein). It is worth to note, that the redshift distribution of the quasars is not the same for different luminosities due to possible evolution effects. The largest quasar surveys are 2dF and SDSS. In contrast to 2dF, which is finished for today and has 2QZ catalogue as a result \cite{2dF-XII}, SDSS \cite{SDSS5} is in progress and the area covered by it is increasing. Only a part of objects primarily classified as quasars were justified by spectra analysis and were included into SDSS Quasar Catalogue IV \cite{dr4}. However even photometric classification of quasars with color diagrams \cite{phot} is sufficient for using these objects for statistical purposes \cite{myers2, myers1}. Some attempts to find luminosity dependence of the quasar clustering have been made with different samples. E.g. Adelberg \& Steidel \cite{adelberg}, who worked with their own survey, pointed to luminosity independent quasar clustering. 2dF-team, that studied 2QZ survey and found the little redshift evolution in the amplitude of the power spectrum \cite{2dF-IV}, \cite{2dF-XI} and significant increase in clustering amplitude at high redshifts \cite{2dF-XIV, 2dF-II}, detected only marginal evidence for quasars with brighter apparent magnitudes to have a stronger clustering amplitude \cite{2dF-IX}. Porciani and Norberg \cite{porciani2006} also found weak luminosity dependence of the clustering in 2QZ survey. Furthermore they noted that samples with different redshifts show different trends in luminosity dependence: in the redshift bin $1.7<z<2.1$ the brightest quasars seem to be more clustered (have larger bias parameter); in the redshift bin $1.3<z<1.7$ the bias parameter seems to follow a U-shape; and the low redshift bin ($0.8<z<1.3$) does not show any particular trend. Myers et al. \cite{myers2}, working with $\sim$300,000 photometrically classified quasars from the 4th Data Release of the SDSS, detected no significant luminosity dependence and pointed out that a 3$\sigma$ detection of such effect requires a sample several times larger. The present work deals with the study of the luminosity dependence of the quasar clustering for quasars from the Fifth Data Release of SDSS. The sample and its peculiarities are described in Section 2. The results of estimation of the correlation length as a function of luminosity are presented in Section 3. As we do not have such large sample, which we need, according to Myers et al. \cite{myers2}, for precise measurements of the correlation function, we used another technique, which is described in Section 4. Finally, Section 5 is dedicated to discussion of obtained results. | \indent \indent As we can see from the Table 2 to obtain any reliable results with the first method one need much larger samples, which are unavailable for today. Moreover the correlation length is the spatial scale of clustering on the one hand and a measure of the clustering amplitude on the other hand. Thus it is not easy to interpret the first technique results unambiguously. That is why the second method seems to be better for this purpose because it is more direct and does not require such large samples. From the results of the second method we can say that brighter quasars reside in closer pairs than faint ones. Note, that such splitting of f(r) curves on larges scales could also be the consequence of the different redshift distributions of the quasars with different luminosities. But we know, that the bright quasars represent higher redshifts and the quasar density decreases with the redshift, thus the inverse effect would be present. Anyway the difference between the mean redshift for subsamples of quasars with different luminosities (see the last column of Table 1) within each redshift interval is negligibly small and cannot affect the results. Anyway this problem requires further investigations and improvement of the methods. It would be interesting to compare for example, the same $f(r)$ function for the fifth nearest neighbour or cross-correlation with galaxies. The last possibility could increase the sample, but this could be applied only for the lower redshift intervals, than those used in the present work. Summing up the obtained results one can speal about luminosity dependence of the quasar clustering, but this dependence is not strong, which is in agreement with the results by Porciani \& Norberg \cite{porciani2006} and theoretical predictions by Lidz et al. \cite{lidz}. But even if this dependence is weak, we do not have to neglect it. Note that when we estimate e.g. the correlation length of the quasars on the low redshift we obtain the mean correlation length averaged over all the quasars with different luminosities. But on the high redshifts the obtained results correspond only to the correlation length for bright quasars and do not reflect the whole picture. That is why possible effects of luminosity dependence of quasar clustering should be taken into account when studying the redshift evolution of it. But even so one should keep in mind that the quasars (as another AGNs) do not reflect the whole distribution of extragalactic objects because quasars are considered to reside in the strongest peaks of the matter density. This fact could explain larger values of the correlation function for quasars than for galaxies (see \cite{blake}, \cite{gonzalez} for comparision). | 7 | 10 | 0710.3659 |
0710 | 0710.1140_arXiv.txt | We use narrowband imaging (\fwhm{} $=70$\,\AA{}) to select a sample of emission line galaxies between $0.20 \lesssim z \lesssim 1.22$ in two fields covering 0.5 sq.~deg. We use spectroscopic follow-up to select a sub-sample of \ha{} emitting galaxies at $z\sim0.24$ and determine the \ha{} luminosity function and star formation density at $z\sim0.24$ for both of our fields. Corrections are made for imaging and spectroscopic incompleteness, extinction and interloper contamination on the basis of the spectroscopic data. When compared to each other, we find the field samples differ by $\Delta \alpha = 0.2$ in faint end slope and $\Delta \log [ L^* (\mathrm{\ergs}) ] = 0.2$ in luminosity. In the context of other recent surveys, our sample has comparable faint end slope, but a fainter $L^*$ turn-over. We conclude that systematic uncertainties and differences in selection criteria remain the dominant sources of uncertainty between \ha{} luminosity functions at this redshift. We also investigate average star formation rates as a function of local environment and find typical values consistent with the field densities that we probe, in agreement with previous results. However, we find tentative evidence for an increase in star formation rate with respect to the local density of star forming galaxies, consistent with the scenario that galaxy-galaxy interactions are triggers for bursts of star formation. | It is now widely accepted that the amount of star formation in Universe as a whole has increased since the formation of the first galaxies, peaking around redshifts $z\sim2-3$ and subsequently declining by a factor of ten \citep[e.g.][and references therein]{Hopkins04}. Cosmic star formation history provides strong constraints on models of galaxy formation and evolution \citep{Pei99,Somerville01}, because it directly traces the accumulation of stellar mass and metal fraction \citep{Pei95,Madau96} to their present-day values \citep{Cole01,Panter03}. Its rapid decline over the past 8 Gyr is consistent with ``downsizing'' scenarios \citep{Cowie96} in which the more massive galaxies have produced their stellar mass at earlier times than the less massive galaxies \citep{Heavens04,Juneau05,Thomas05,Fardal06}. The star formation history of the universe has also been used to constrain allowable stellar initial mass functions \citep{Baldry03,Hopkins06} and cosmic supernova rates \citep{GalYam04,Daigne06}. Star forming galaxies exhibit a strong UV continuum courtesy of newly formed OB stars in sites of star formation. This newborn population can be inferred from the UV directly \citep[e.g.][]{Treyer98,Lilly96} or through a host of indirect calibrators spread across the electromagnetic spectrum \citep{RosaGonzalez02,Condon92,Schaerer00}. At low redshifts the most direct calibrator -- and of the optical calibrators the least affected by internal extinction -- is the \ha{} recombination line, which emits when stimulated by ionising UV radiation \citep[e.g.][]{Kennicutt98}. Narrowband surveys at optical wavelengths have long been recognised as a powerful way of yielding large samples of emission line galaxies, including those selected by \ha{} at redshifts $z\lesssim0.4$ \citep{Ly07,Pascual07,Jones01}. They are advantageous in that they select galaxies in exactly the same quantity that they seek to measure, and are optimised for the detection of the faint emission line signatures indicative of star formation. Narrowband surveys also have the advantage of a simplified selection function, with filters that probe only a very narrow redshift slice, thereby yielding a volume limited sample at a common distance. Many recent emission line surveys have targeted \lya{} at high redshift \citep{Ajiki03,Hu04,Rhoads04,Gawiser06}, as well as \ha{}, \hb{}, \oiii{} and \oii{} at lower redshifts \citep{Fujita03,Hippelein03,Ly07}. Here we describe a survey for \ha{} emission line galaxies at $z\sim0.24$, found as a by-product of the Wide Field Lyman Alpha Search \citep[WFILAS;][]{Westra05,Westra06}. The resulting sample has been utilised to determine the \ha{} luminosity function at $z\sim0.24$ and its associated co-moving star formation density. In Section~\ref{sec:candsel} we describe the selection of candidates using narrow- and broadband imaging. In Section~\ref{sec:spectroscopy} we detail follow-up spectroscopy used to identify the nature of the emission and test completeness of the sample. In Section~\ref{sec:lumfieSFD} we derive the \ha{} luminosity function for galaxies at $z\sim0.24$ and explore its variation with the local environment in Section~\ref{sec:environment}. A summary and concluding remarks are made in Section~\ref{sec:conclusionsLowz}. Throughout this paper we assume a flat Universe with $(\Omega_{\rm m}, \Omega_{\Lambda}) = (0.3,0.7)$ and a Hubble constant $H_0 = 70$\,\kmsMpc. All quoted magnitudes are in the {\it AB} system \citep{Oke83}\footnote{$m_{AB} = -2.5 \log f_\nu - 48.590$, where $m_{AB}$ is the {\it AB} magnitude and $f_\nu$ is the flux density in ergs\,s$^{-1}$\,cm$^{-2}$\,Hz$^{-1}$}. | \label{sec:conclusionsLowz} In this paper we report the results of a survey for \ha{} emitting galaxies at $z\sim0.24$. We used two fields from the Wide Field Imager Lyman Alpha Search (WFILAS). It consists of imaging in three narrowband filters (\fwhm{} $=70$\,\AA{}), an encompassing intermediate band filter (\fwhm{} $=220$\,\AA{}), supplemented with broadband $B$ and $R$. The narrowband filters cover a redshift range of $0.23 \lesssim z \lesssim 0.26$ for \ha{} galaxies. These galaxies were selected by having an excess flux in one of the narrowband over to the other two, while also being detected in the intermediate and broadband $R$ filters. This yielded a total of 707 candidate emission line galaxies (after the removal of stellar contaminants) for both fields. Of the 372 and 335 candidates, we observed 301 and 255 through spectroscopic follow-up for the CDFS and S11 fields, respectively. We have identified emission in 189 and 117 candidates and confirmed that around half of these galaxies are \ha{} at $z\sim0.24$. A significant number of galaxies were also found at $z\sim0.21$ by means of their \sii{} emission. Other galaxies found were \oii{} and \hb{}/\oiii{} emitters at $z\sim1.2$ and $z\sim0.6-0.7$, respectively. Through use of the spectroscopy, we refined our colour selection to account for galaxies with a single emission line, leading to a measure of the fraction of \ha{} galaxies as a function of narrowband flux in both of these regions of the sky. We also used the spectroscopy to determine a generic extinction correction using the Balmer decrement. We have determined the \ha{} luminosity function at $z\sim0.24$ separately for both of our fields after correcting for imaging and spectroscopic incompleteness, extinction and contamination from interlopers. We find small differences in their slope and turn-over luminosity while their normalisations were the same. When compared to recent \ha{} surveys, there is remarkable agreement between the luminosity function of our CDFS field with that one the Fabry-Perot imaging survey of \citet{Hippelein03}. Differences between our fields were of the order expected by cosmic variance but less than the scatter between the \ha{} luminosity functions of recent surveys. We surmise that while cosmic variance is a major contributor to this scatter, it is differences in methodology between surveys (mainly differences in selection criteria) that dominate discrepancies between \ha{} luminosity functions and its related observables at $z\sim0.24$. A survey that covers $10-20\times$ the volume of one of our fields is required to get the uncertainty due to cosmic variance to the levels of \citet{Gallego95}. We estimated the star formation density for both our fields to be $\log \dot{\rho} = -1.93 ^{+0.08} _{-0.10}$ and $-2.24 ^{+0.11} _{- 0.14}$ ($\dot{\rho}$ in \Msunyr{}) for the CDFS and S11 fields, respectively, down to our survey limit of $\log F_\mathrm{line} = -16.0$ ($F_\mathrm{line}$ in \lineunits{}) or $\log L_\mathrm{line} = 40.6$ ($L_\mathrm{line}$ in \ergs{}). These values are comparable to other surveys at this redshift when calculated to the same flux limit. Correcting for AGN would decrease these values by 0.02 to 0.04 depending on exactly how much of the \ha{} flux is contributed by the active nucleus rather than by normal star formation. Furthermore, we determined the star formation density in the hypothetical case where \sii{} emitters at $z\sim0.21$ were classified as \ha{} to illustrate the problems associated with solely relying colour selections. The star formation density $\log \dot{\rho}$ of the S11 field does not change by much (+0.02). On the other hand, the star formation density in the CDFS increases by 0.12, due to the large number of foreground \sii{} galaxies at $z\sim0.21$. We explored the amount of star formation with respect to the local environment and found that the star formation rates were typical for the field galaxy densities probed, in agreement with the results of previous work. However, we also found tentative evidence of an increase in star formation rate per galaxy with increasing density of the star forming galaxies. This supports scenarios where merger events are triggers for enhanced star formation, provided it can be demonstrated to be occurring on the smallest scales. We explored this trend by examining the spatial distribution of our fields individually and found that it was largely attributable to one field. A formal study of the clustering statistics of this field is required to confirm this and will be the subject of a future study. | 7 | 10 | 0710.1140 |
0710 | 0710.4881_arXiv.txt | Inclusive neutrino-nucleus cross sections are calculated using a consistent relativistic mean-field theoretical framework. The weak lepton-hadron interaction is expressed in the standard current-current form, the nuclear ground state is described with the relativistic Hartree-Bogoliubov model, and the relevant transitions to excited nuclear states are calculated in the relativistic quasiparticle random phase approximation. Illustrative test calculations are performed for charged-current neutrino reactions on $^{12}$C, $^{16}$O, $^{56}$Fe, and $^{208}$Pb, and results compared with previous studies and available data. Using the experimental neutrino fluxes, the averaged cross sections are evaluated for nuclei of interest for neutrino detectors. We analyze the total neutrino-nucleus cross sections, and the evolution of the contribution of the different multipole excitations as a function of neutrino energy. The cross sections for reactions of supernova neutrinos on $^{16}$O and $^{208}$Pb target nuclei are analyzed as functions of the temperature and chemical potential. | Introduction} Neutrino-nucleus reactions at low energies play an important role in many phenomena in nuclear and particle physics, as well as astrophysics. These reactions present extremely subtle physical processes, not only because they involve the weak interaction, but also because they are very sensitive to the structure of nuclear ground states and excitations, i.e. to the solution of the nuclear many-body problem that includes the strong and electromagnetic interactions. The use of microscopic nuclear structure models in a consistent theoretical framework is therefore essential for a quantitative description of neutrino-nucleus reactions~\cite{Hay.99}. Detailed predictions of neutrino-nucleus cross sections are crucial for the interpretation of neutrino experiments, detection of neutrinos produced in supernova explosions and understanding the underlying nature of these explosions~\cite{sns1}. Neutrino-nucleus reactions which occur in a type II supernova could also contribute to the nucleosynthesis~\cite{Heg.05,Lan.06}, but more data on cross sections are necessary for a more complete understanding of this process, as well as the supernova dynamics. Data on neutrino-nucleus cross sections have been obtained by the LSND~\cite{Alb.95,Ath.97} and KARMEN~\cite{Bod.92,Bod.94,Mas.98} collaborations, and at LAMPF~\cite{Kra.92,Koe.92}, but only for $^{12}$C and $^{56}$Fe target nuclei. New experimental programs are being planned which will provide essential data on the neutrino-nucleus reactions, and also help to improve the reliability of present cross-section calculations. These include the spallation neutron source (SNS) at ORNL, where neutrinos produced by pion decay at rest will enable measurements of cross sections for a wide range of target nuclei~\cite{Avi.03,Efr.05}, and the promising ``beta-beams'' method for the production of pure electron neutrino-beams by using the $\beta$-decay of boosted radioactive ions~\cite{Zuc.02,Vol.04}. Cross sections for neutrino energies in the range of tens of MeV could be measured, and these reactions are particularly interesting for supernova studies~\cite{McL.04}. Weak interaction rates at low energies have been analyzed employing a variety of microscopic approaches, principally in the frameworks of the shell model~\cite{Hax.87,Eng.96,Hay.00}, the random phase approximation (RPA)~\cite{Aue.97,Sin.98,Vol.00,Vol.02}, continuum RPA (CRPA)~\cite{Kol.92,Kol.95,Jac.99,Jac.02,Bot.05}, hybrid models of CRPA and shell model~\cite{Kol.99,Kol.03}, and the Fermi gas model~\cite{Wal.75,Gai.86,Kur.90}. The shell model provides a very accurate description of ground state wave functions. The description of high-lying excitations, however, necessitates the use of large model spaces and this often leads to computational difficulties, making the approach applicable essentially only to light and medium-mass nuclei. For systematic studies of weak interaction rates throughout the nuclide chart, microscopic calculations must therefore be performed using models based on the RPA. Hybrid models combine the shell-model and CRPA in such a way that occupation probabilities of single-particle states in a specific nucleus are determined by shell model calculations, and then inserted into the CRPA~\cite{Kol.99}. In general the CRPA employs different interactions for the calculation of the nuclear ground state (for instance, the Woods-Saxon potential), and in the residual CRPA interaction (G matrix from the Bonn potential, or the Landau-Migdal interaction), and thus additional parameters are required in order to adjust the calculated rates to data. A more consistent approach, based on the quasiparticle RPA with Skyrme effective interactions, has been employed in calculations of weak interaction rates~\cite{Vol.00}. Although in this framework the residual RPA interaction is derived from the same energy functional which determines the nuclear ground state, some terms are usually omitted, and pairing correlations require the adjustment of additional factors. A fully consistent theoretical framework for the description of charge-exchange excitations in open shell nuclei, based on Skyrme effective interactions, has only recently been developed \cite{Fra.05}, but not yet employed in the analysis of neutrino-nucleus reactions. Neutrino-nucleus cross sections have also been calculated using the ab-initio no-core shell model based on realistic NN and three-body interactions~\cite{Hay.03}, and the shell model with an improved Hamiltonian which properly takes into account the spin-isospin interactions~\cite{Suz.06}. The importance of improved calculations of neutrino-nucleus cross sections for neutrino-oscillation studies has been demonstrated in the recent reanalysis of the LSND experiment~\cite{Sam.06}, using the particle-number projected quasiparticle RPA~\cite{Krm.05}, which has shown an enhancement of the neutrino-oscillation probability when compared to previous studies. Although the general expressions for the transition matrix elements relevant for the calculation of neutrino-nucleus cross sections have been known since many years~\cite{Con.72,Don.80}, it is only more recently that systematic calculations have been performed in open-shell nuclei by making use of modern effective interactions in the description of both nuclear ground states and excitations~\cite{Vol.00}. Reliable prediction of weak interaction rates in nuclei necessitates a fully consistent description of the structure of ground states and multipole excitations. Among the relevant charge-exchange excitations, the isobaric analog state (IAS) and Gamow-Teller resonance (GTR) have been the subject of extensive experimental and theoretical studies. Much more limited are the data and theoretical predictions for properties of excitations of higher multipolarities at finite momentum transfer. In this work we analyze charged-current neutrino-nucleus reactions by employing a fully consistent microscopic approach based on relativistic energy density functionals, and also including pairing correlations in the description of open-shell target nuclei. An essential advantage over most current approaches is the use of a single universal effective interaction in calculations of both ground-state properties and multipole excitations of nuclei in various mass regions of the chart of nuclides. Of particular interest for the present study are rates for neutrino-nucleus reactions in the low-energy range below 100 MeV, which play an important role in many astrophysical processes, including stellar nucleosynthesis. A quantitative description of nucleosynthesis of heavy elements during the r-process necessitates accurate predictions of neutrino-nucleus cross sections not only in stable nuclei, but also in nuclei away from the valley of $\beta$-stability. Since nuclei are used as detectors for solar and supernovae neutrinos, as well as in neutrino oscillation experiments, it is important to describe the neutrino detector response in a consistent and fully microscopic theory. Finally, a quantitative estimate of neutrino-nucleus reaction rates will provide information relevant for feasibility studies and simulations of a low-energy beta beam facility, which could be used to produce neutrino beams of interest for particle physics, nuclear physics and astrophysics~\cite{Vol.04}. In Sec.~\ref{sectionII} we outline the basic formalism used in the evaluation of neutrino-nucleus cross sections. In Sec.~\ref{sectionIII} the relativistic Hartree-Bogoliubov model (RHB), and the proton-neutron relativistic quasiparticle random phase approximation (PN-RQRPA) are briefly reviewed. Section~\ref{sectionIV} includes several test cases, and the calculated cross sections are compared with results of previous theoretical studies and the available data. Cross sections for supernova neutrinos are analyzed in Sec.~\ref{sectionV}, and finally Sec.~\ref{sectionVI} summarizes the results of the present investigation and ends with an outlook for future studies. | 7 | 10 | 0710.4881 |
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0710 | 0710.4553_arXiv.txt | We present detailed images of diffuse UV intergalactic light (IGL), situated in a 60\,kpc halo that surrounds the radio galaxy MRC\,1138-262 at z=2. We discuss the nature of the IGL and rule out faint cluster galaxies, nebular continuum emission, synchrotron, inverse Compton, synchrotron self-Compton emission and scattering of galactic stellar light as possible sources of the IGL. Dust scattered quasar light is an unlikely possibility that cannot be ruled out entirely. We conclude that the source of the IGL is most likely to be a young stellar population distributed in a halo encompassing the radio and satellite galaxies, undergoing star formation at a rate greater than 57$\pm8$\Msunpyr. Within 70\,kpc of the radio core, approximately 44\% of the star formation that is traced by UV light occurs in this diffuse mode. The average UV colour of the IGL is bluer than the average galaxy colour, and there is a trend for the IGL to become bluer with increasing radius from the radio galaxy. Both the galaxies and the IGL show a UV colour--surface brightness relation which can be obtained by variations in either stellar population age or extinction. These observations show a different, but potentially important mode of star formation, that is diffuse in nature. Star formation, as traced by UV light, occurs in two modes in the high redshift universe: one in the usual Lyman break galaxy clump-like mode on kpc scales, and the other in a diffuse mode over a large region surrounding massive growing galaxies. Such a mode of star formation can easily be missed by high angular resolution observations that are well suited for detecting high surface brightness compact galaxies. Extrapolating from these results, it is possible that a significant amount of star formation occurs in large extended regions within the halos of the most massive galaxies forming at high redshift. | Distant powerful radio galaxies are important probes of the formation and evolution of massive galaxies, because they are among the most luminous and massive galaxies known in the early Universe, and are likely progenitors of dominant cluster galaxies \cite[e.g][]{Pentericci1997,Villar-Martin2006,Seymour2007}. They are generally embedded in giant (cD-sized) ionized gas halos \cite[e.g][]{Reuland2003} and surrounded by galaxy overdensities \citep{Pentericci2000,Venemans2007}. With radio lifetimes (few $\times10^7$yr) being much smaller than cosmological timescales, the statistics are consistent with every dominant cluster galaxy having gone through a luminous radio phase during its evolution. The deepest image of a distant radio galaxy with the {\it Hubble Space Telescope (HST)} is a recent {\it Advanced Camera for Surveys (ACS)} observation of MRC\,1138-262 at z$ = 2.156$ \citep{Miley2006}. This system has been dubbed the Spiderweb galaxy as its complex morphology resembles a spider's web. The {\it ACS} image shown in Fig.~1 shows tens of UV bright galaxies surrounding a central galaxy (marked with a white cross). As the radio core is cospatial with the central galaxy we refer to this central galaxy as the radio galaxy, and all the surrounding galaxies are referred to as the satellite galaxies. All of these galaxies may merge together to form a single massive galaxy at z=0, therefore the whole complex system is referred to as the Spiderweb system This system is one of the most intensively studied distant radio galaxies \citep{Pentericci1997,Pentericci1998,Pentericci2000} and exhibits the properties expected for a progenitor of a dominant cluster galaxy. The infrared luminosity provides an upper limit to the stellar mass of 10$^{12}$\Msun\ \citep{Seymour2007}, so it is one of the most massive galaxies known at z$>$2. The host galaxy is surrounded by a giant Ly$\alpha$ halo and the high rotation measure of the radio source means that it is embedded in a dense medium with an ordered magnetic field \citep{Pentericci1997}. The radio galaxy is associated with a $>$3\,Mpc-sized structure of galaxies with an estimated mass $>4\times10^{14}$\Msun\ \citep{Venemans2007}, indicating that it is the predecessor of a local rich cluster. Mechanical feedback from the active galactic nucleus (AGN) is sufficient to expel significant fractions of the interstellar medium of the massive gas-rich galaxy \citep{Nesvadba2006}. Therefore the AGN may quench the star formation and allow the radio galaxy to evolve onto the red sequence as suggested by popular models of massive galaxy formation \citep{Croton2006}. The study of the formation of the most massive galaxies through these and other mechanisms has so far mainly been limited to models and simulations \cite[e.g][]{Dubinski1998,Gao2004,DeLucia2007}. Our aim is to use these observations of a high-redshift radio galaxy to study brightest cluster galaxy formation. In this work we present detailed images of widespread UV intergalactic light (IGL) situated in a 60\,kpc halo that lies between the radio galaxy and the UV bright satellite galaxies. We examine the possible origins of this light and present in-situ star formation as the most probable. In section \ref{method} we describe the observations and data reduction, in section \ref{results} we describe the distribution and colour of the IGL, and discuss possible origins in section \ref{nature}. Section \ref{discussion} discusses the implications of widespread star formation in a halo around a massive forming brightest cluster galaxy. Throughout this work we use $H_0=71$, $\Omega_M=0.27$, and $\Omega_\Lambda=0.73$ \citep{Spergel2003}. All magnitudes are AB magnitudes. At a redshift of 2.156, the linear scale is 8.4\,kpc/". | The radio galaxy MRC\,1138-262 is surrounded by a halo of IGL that extends across 60\,kpc. After examining nebular continuum emission, synchrotron, inverse Compton, synchrotron self-Compton emission, scattering, and stripping of stars as possible sources of the IGL, we conclude the most likely origin of the IGL is in-situ star formation. The minimum star formation rate (i.e. uncorrected for dust) of the IGL is 57$\pm8$\Msunpyr, which is comparable to the total star formation rate derived from the UV luminosity integrated over all the galaxies within 70 kpc of the radio galaxy (and including the radio galaxy itself). Applying a minimum dust correction of $E(B-V)\sim0.1$ imply by the red colours of the the IGL increases the star formation rate of the IGL to 142\Msunpyr, and that of the whole Spiderweb system to more than 325\Msunpyr. We estimate that the total observed star formation rate can produce approximately $\sim$7\% of the \lya\ emission in the halo surrounding the galaxy. The radial colour gradient of the IGL indicates that there is a smooth range in extinction or stellar population age from the outer to the inner parts of the halo, where the inner parts are older or dustier than the outer region. While the presence of large quantities of ionized gas found around several remarkable species of high redshift galaxies has been attributed to a number of proposed mechanisms that include AGN feedback, infall, cooling flows, starburst superwinds and mergers, these observations of MRC\,1138-262 show that any successful model should be able to accommodate the mode of extended star formation present in this paper. Our data suggest that the formation of the most massive galaxies is connected with that of their gaseous envelopes and star forming halos, part of which may precede the intracluster light or cD envelopes, or perhaps may contribute to a satellite population. A significant amount of star formation might be occurring in the form of extended low surface brightness features, beyond the typical UV detection limits, as well as in largely obscured extended halos as detected at infrared and sub-mm wavelengths. | 7 | 10 | 0710.4553 |
0710 | 0710.0589.txt | Large-scale asymmetries in the stellar mass distribution in galaxies are believed to trace non-equilibrium situations in the luminous and/or dark matter component. These may arise in the aftermath of events like mergers, accretion, and tidal interactions. These events are key in the evolution of galaxies. In this paper we quantify the large-scale lopsidedness of light distributions in 25155 galaxies at $z < 0.06$ from the Sloan Digital Sky Survey Data Release 4 using the $m = 1$ azimuthal Fourier mode. We show that the lopsided distribution of light is primarily due to a corresponding lopsidedness in the stellar mass distribution. Observational effects, such as seeing, Poisson noise, and inclination, introduce only small errors in lopsidedness for the majority of this sample. We find that lopsidedness correlates strongly with other basic galaxy structural parameters: galaxies with low concentration, stellar mass, and stellar surface mass density tend to be lopsided, while galaxies with high concentration, mass, and density are not. We find that the strongest and most fundamental relationship between lopsidedness and the other structural parameters is with the surface mass density. We also find, in agreement with previous studies, that lopsidedness tends to increase with radius. Both these results may be understood as a consequence of several factors. The outer regions of galaxies and low-density galaxies are more susceptible to tidal perturbations, and they also have longer dynamical times (so lopsidedness will last longer). They are also more likely to be affected by any underlying asymmetries in the dark matter halo. | %\citet{rz95}: 18 galaxies. Galaxy disks exhibit a wide variety of shapes visible in the near-IR. They are due to distortions in surface density distributions, not mass-to-light ratios. Stellar velocities in nonaxisymmetric galaxies differ by 3-6\% from those in symmetric ones. %\citet{zr97}: 60 galaxies, magnitude-limited. 30\% of galaxies have $A_1 > 0.2$. Lopsided mass distributions remain long enough to indicate past interactions, not just ongoing ones. Often the companion is not obvious, either merged or fled. %\citet{rr98}: 54 early-type disk galaxies. Low SFRs to reduce possibility of asymmetric stellar distributions and increase possibilty of light->mass asymmetry tracing. Traces mass dependence because VRI bands show similar lopsidedness and it's OK to go shorter than I and K bands. 20\% of galaxies have lopsidedness above 0.19. %\citet{rrk00}: Correlations between lopsidedness and recent SF, and current SF were found, lopsided = star-forming. Significant fractions of stellar content can be created in a short time from minor mergers and on a similar timescale (1 Gyr). It has long been recognized that galaxies show large-scale asymmetries in their structure \citep{bl+80}. Lopsided galaxies have such asymmetries where one side of their disk is more massive and/or more extended than the opposite side. This lopsidedness can be traced in the spatial structure of the stars \citep{rz95} and/or the HI gas \citep{rs94} and/or in the large-scale kinematics of this material \citep{ss+99}. There are a variety of mechanisms or events that have been proposed to produce the observed lopsidedness. All of them involve a time-dependent non-equilibrium dynamical state, in most cases triggered through an external process. Such external processes are a natural consequence of the standard Lambda Cold Dark Matter cosmological framework. This implies that galaxies assemble hierarchically (a process that is on-going). Examples that can lead to lopsidedness include a minor merger (\citealt{wm+96}; \citealt{zr97}), the tidal interaction resulting from a close encounter between roughly equal-mass galaxies \citep{kl+02}, and the asymmetric accretion of intergalactic gas into the disk (\citealt{bc+05}; \citealt{kk+05}). Other mechanisms involve the dark matter halo: stars and gas orbiting in a lopsided dark matter halo (\citealt{w94}; \citealt{j97}; \citealt{j99}) or a stellar/gas disk that is offset with respect to the center of the dark matter halo (\citealt{ls98}; \citealt{ns+01}). These also involve past tidal interactions and/or mergers that have perturbed the dark matter halo, but such perturbations may be quite long-lived. Finally, dynamical processes internal to the disk that lead to mildly lopsided distributions have also been investigated (\citealt{st+90}; \citealt{st96}; \citealt{mt97}). A variety of programs to study lopsidedness have been undertaken over the past decade. Most of these investigations have studied the lopsided distribution of the stellar component through analysis of optical and near-infrared images. \citet{zr97} studied a magnitude-limited sample of 60 field spiral galaxies. They measured lopsidedness as the radially averaged, azimuthal $m=1$ Fourier amplitude $A_1$ of the light (see Section \ref{sec:error} below) and computed lopsidedness between 1.5 and 2.5 scale lengths in the galactic disks. The value of $A_1$ indicates the typical large-scale variation in mass density from side to opposite side at the same distance from the galactic center. The mass density typically varies from between $1 \pm A_1$ times the average density at the same radius. They found that $\sim 30$\% of field spiral galaxies exhibited significant lopsidedness ($A_1 > 0.2$). \citet{rr98} followed up this work by studying lopsidedness in 54 early-type galaxies and found that $\sim 20$\% had $A_1 > 0.19$. \citet{cbj00} studied a sample of 113 $z < 0.01$ galaxies (elliptical, spiral, and irregular) but used a 180-degree rotational asymmetry measure $A_{180}$. They found that asymmetry is strongly dependent on morphological type, with lower asymmetry in elliptical and lenticular galaxies and higher asymmetry in late-type disk and irregular galaxies. More recently, \citet{bc+05} have measured the Fourier $A_1$ parameter for 149 galaxies in the Ohio State University Bright Galaxy Survey. They confirmed that a large fraction of galaxies have significant lopsidedness in their stellar disks, with late-type galaxies being more lopsided. Lopsidedness in the {\it light} distribution can be produced by either a corresponding asymmetry in the underlying {\it mass} distribution in the stellar population or by large-scale variations in the mass-to-light ratio (e.g., from star formation and dust obscuration). \citet{rz95} investigated this issue with a sample of 18 face-on spiral galaxies imaged in the K$^\prime$ (2.2$\mu$ m) band where the effect of young stars or dust is minimized. They found that about a third of the sample showed significant lopsidedness (similar to results from optical investigations). Similarly, \citet{rr98} found that lopsidedness in early-type disk galaxies is nearly identical when observed in the $V$, $R$, and $I$ bands. They concluded that an asymmetric mass distribution then accounts for the majority of the asymmetry in the light distribution in these galaxies. Lopsidedness has also been studied in the distribution of HI gas. Since the HI can frequently be traced to significantly larger radii than the stars, these investigations are highly complementary to the optical image analysis. Due to the time-consuming nature of HI interferometric mapping, only modest size samples have been analyzed in this way \citep{ss+99}. On the other hand, \citet{rs94} have examined the global HI line profiles for roughly 1700 galaxies, and shown that at least 50\% are significantly asymmetric (confirming that the large-scale HI distribution is frequently lopsided). HI maps also show that apart from a lopsided distribution of the gas the HI rotation curves are often asymmetric \citep{ss+99}. The connection between the phenomena of structural and kinematic lopsidedness in galaxies is not yet clear \citep{ss+99}. Despite these diverse investigations and the abundance of proposed models, the origin of lopsidedness remains unsettled. For models involving tidal interactions or minor mergers, there is an expected link between lopsidedness and the local environment. The evidence in this regard has been mixed (e.g., \citealt{wp04}; \citealt{bc+05}; \citealt{aj+06,aj+07}; \citet{dc+07}). %\citet{dc+07} have undertaken %the most comprehensive investigation so far of this issue. Using the %rotational asymmetry measure $A_{180}$ to study a sample of over 3000 %galaxies, they find that close pairs of galaxies are more asymmetric %than other galaxies and that the asymmetry increases as the pair %separation decreases. They conclude that these global asymmetries %trace recent tidal interactions or mergers. The investigations summarized above have all involved relatively small samples of galaxies, making it difficult to assess the overall distribution of asymmetry or lopsidedness as a function of the basic parameters that characterize the structure of galaxies. This is the first of three papers in which we use the wealth of data available from the Sloan Digital Sky Survey (SDSS) to extend these studies of small samples (of-order one hundred galaxies) to large samples (tens of thousands). In this paper, we describe our sample selection and methodology. We also relate lopsidedness to the basic structural properties of the galaxies. In Paper II we will investigate the connection between lopsidedness and both star formation and black hole growth in galaxies. Finally, in Paper III we will examine the connection between lopsidedness and the local galaxy environment. In \S\ref{sec:data}, we begin by presenting an initial low-redshift sample from the SDSS and describe the observations and properties for its galaxies. Next, we explain our lopsidedness calculation. In \S\ref{sec:syserr}, we address the major data quality issues that limit the reliability of the measurements for portions of the sample. On this basis, we apply cuts on the observational parameters to weed out the problematic cases for our subsequent analysis. Next, \S\ref{sec:lopprops} describes the lopsidedness of galactic light distributions in different optical/near-IR bands, its correspondence with lopsided mass distributions, and its radial dependence. We then examine the relationship between lopsidedness and the basic structural properties of galaxies in \S\ref{sec:lopsfh}. Finally, we summarize our findings in \S\ref{sec:summary}. | } We have measured large-scale galactic asymmetry for a large sample of low-redshift ($z <$ 0.06) galaxies drawn from the Sloan Digital Sky Survey. Our use of lopsidedness, a radially averaged $m=1$ azimuthal Fourier mode, has proven useful for a large fraction of the sample. Images of a minority of galaxies in the sample have poor observational properties that cause significant systematic errors in the lopsidedness calculation, and these galaxies were removed from the sample via cuts on angular size, signal-to-noise, and ellipticity/inclination. Those cuts removed a higher fraction of the high-mass, high-mass-density, and high-concentration galaxies than those with low values of these structural properties. Nonetheless, the resulting sample is well represented by the galaxies of the same range of structural properties as the original sample. We find that there are no systematic differences between the $(g-i)$ colors of the brighter and fainter sides of lopsided galaxies. This implies that there is no systematic difference in the mass/light ratio \citep{k+07}, and hence that the lopsided light distributions are primarily caused by lopsided distributions in the stellar mass. We have verified this through analysis of the relationship between color and mass/light ratio for both model galaxy spectral energy distributions and SDSS galaxy data. However, for our sample the lopsidedness in the $g$-band tends to be slightly greater than in the $r$- and $i$-bands. Thus, some of the lopsidedness in the light does arise from the effects of star-formation and/or dust extinction (which will more strongly affect the $g-$band light). Lopsidedness is a structural property that depends strongly on other structural properties. Galaxies with progressively lower concentration, stellar mass, or stellar mass density tend to have progressively higher lopsidedness. We show that the strongest and most fundamental correlation is between lopsidedness and stellar mass density. We also find that lopsidedness increases systematically with increasing radius, particularly for late-type galaxies. Lopsidedness can be induced through tidal stress associated with interactions with a companion galaxy or through accretion or minor mergers (e.g. \citealt{zr97}; \citealt{bc+05}). Galaxies with low density will be most affected by tidal stress, and the effects of a tidal perturbation will last longer in such systems due to the longer dynamical times. The same arguments pertain to the outer parts of galaxies. Thus, the two above results make good physical sense. Alternatively, if the dark matter halo is lopsided, its effects on the structure of the stellar disk will be more pronounced in the outer region and in galaxies with low mass and low density (where dark matter is more dynamically important). The relatively large values of lopsidedness we measure to be commonplace ($A_1 > 0.1$) appear to be too large to be generated by internally generated dynamical processes (e.g., \citealt{mt97}). Our overall goal in this investigation has been to use lopsidedness as a way of quantifying the signature of moderate or weak global dynamical perturbations. The next step will be to determine the connections between such perturbations and both the on-going/recent star formation and the growth of supermassive black holes in galaxies. These connections can help constrain the processes and conditions that guide the formation and evolution of the galaxies. In future papers we will address these questions using the present sample of galaxies. %are also thought to induce bursts of star formation. Using two well %studied SFH indicators, we have shown that lopsidedness is typically %associated with bursty and fast star formation, and quiescent star %formation is typical in symmetric galaxies. We will extend this %discussion in future work. %Interactions are also thought to induce increased star formation rates %in galaxies. We have found that lopsided galaxies produce stars at %much faster rates per unit stellar mass than symmetric galaxies. %Lopsidedness is then another way to qualitatively describe the stellar %population age of galaxies. Taking lopsidedness as a signature of a %recent interaction, the link between star formation and lopsidedness %may be that they are symptoms of the same interactions. We will %extend this discussion in future work. JB acknowledges the receipt of a FCT post-doctoral grant BPD/14398/2003. We would like to thank Vivienne Wild for reading a draft of the manuscript. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington. | 7 | 10 | 0710.0589 |
0710 | 0710.4547_arXiv.txt | We present an extensive data set of $\sim 150$ localized features from \Cassit{} images of Saturn's Ring~A, a third of which are demonstrated to be persistent by their appearance in multiple images, and half of which are resolved well enough to reveal a characteristic ``propeller'' shape. We interpret these features as the signatures of small moonlets embedded within the ring, with diameters between 40 and 500~meters. The lack of significant brightening at high phase angle indicates that they are likely composed primarily of macroscopic particles, rather than dust. With the exception of two features found exterior to the Encke Gap, these objects are concentrated entirely within three narrow ($\sim 1000$ km) bands in the mid-A Ring that happen to be free from local disturbances from strong density waves. However, other nearby regions are similarly free of major disturbances but contain no propellers. It is unclear whether these bands are due to specific events in which a parent body or bodies broke up into the current moonlets, or whether a larger initial moonlet population has been sculpted into bands by other ring processes. | Saturn's main rings (particularly Ring~A) were determined by the Voyager Radio Science experiment \citep{Zebker85} to be primarily composed of a distribution of icy particles of diameter $D \gtrsim 1$~cm. A steep cutoff in the size-distribution was discerned at $D \sim 20$~m, but particles larger than that value could not be probed due to limitations imposed by the radio experiment's carrier wavelength. At the large end of the particle-size distribution are the two known moonlets embedded in gaps in the outer A~Ring, Pan and Daphnis, of diameters $\sim 28$~km and $\sim 8$~km, respectively \citep{PorcoSci07}. Nothing was known about the distribution of intermediate-size ring particles, between 20~m and 8~km in diameter, until the first evidence of ``missing-link'' particles was found in very-high-resolution images taken by the \Cassit{} spacecraft during its insertion into Saturn orbit \citep{Propellers06}. These intermediate-size moonlets are not directly seen, but rather the propeller-shaped disturbances they create in the ring continuum. This morphology had previously been predicted by numerical simulations \citep{SS00,SSD02,Seiss05}. The sizes of the perturbing moonlets, subject to some ambiguities in interpretation, were given as $D \sim 100$ meters. Their surface densities are quite low relative to smaller particles, corroborating the steep cutoff reported by \citet{Zebker85}. We here present and analyze a data set of 158 localized features in the A~Ring, many of which are well-enough resolved to reveal the characteristic propeller shape. Four of these objects are those reported by \citet{Propellers06}, and another eight were first noted by \citet{Sremcevic07}. Recently, \citet{Espo07} have found evidence for similarly-sized moonlets in the narrow and highly disturbed F~Ring; however, this is not directly applicable to our analysis because there is no particular reason to expect the particle-size distribution of the F~Ring to be simply related to that of the A~Ring. Section~\ref{Propellers} gives further background on the nature and interpretation of propellers. Section~\ref{Observations} summarizes the imaging sequences we used in compiling our data set, and Section~\ref{Analysis} describes the process by which features were identified and characterized with one of two models (``resolved'' or ``unresolved''). Our results are summarized in Section~\ref{Results}. Section~\ref{Interpretation} contains further discussion of the interpretation of propeller features. A full and unabridged presentation of our data set is given as an Appendix. \begin{figure}[!t] \begin{center} \includegraphics[width=10cm,keepaspectratio=true]{f1.pdf} \caption{Particle trajectories under Hill's equations \citep[see, e.g.,][ch.~3.13]{MD99}, showing the propeller-shaped chaotic zones as well as associated moonlet wakes. Radial ($r$)) and azimuthal ($\ell$) coordinates are each shown in units of Hill radii; the perturbing mass is located at the origin, and the $L_1$ and $L_2$ Lagrange points are at $\ell = 0$, $r = \pm 1$ The direction towards Saturn is down, and the orbital direction is to the right. Scattered trajectories deflected by $>5R_H$ are not shown. This simple model, which neglects inter-particle collisions as well as self-gravity, is presented for conceptual purposes only. \label{Illustration}} \end{center} \end{figure} | 7 | 10 | 0710.4547 |
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0710 | 0710.1648_arXiv.txt | We have carried out a search for radio emission from six X-ray dim isolated neutron stars (XDINSs) observed with the Robert C. Byrd Green Bank Radio Telescope (GBT) at 820~MHz. No bursty or pulsed radio emission was found down to a $4\sigma$ significance level. The corresponding flux limit is 0.01--0.04~mJy depending on the integration time for the particular source and pulse duty cycle of 2\%. These are the most sensitive limits yet on radio emission from these objects. | The \emph{ROSAT} mission discovered a group of seven nearby low-luminosity isolated neutron stars, termed the X-ray dim isolated neutron stars (XDINSs). They share very similar properties and are characterized by soft blackbody-like spectra in the range $\sim 40$--100 eV, very faint optical counterparts ($V>25$), long spin periods of 3--12~s (see Table~\ref{table}). For a recent review see \citet{haberl2007}. So far, no confident detections of pulsed radio emission were found from XDINSs. The aim of this project is to search for pulsed and bursty radio emission from XDINSs to link them in their evolutionary scenarios with other classes of neutron stars, such as magnetars and rotating radio transients (RRATs). All three populations of neutron stars have similar properties, such as period, period derivative, age, and magnetic field so that connections between them are plausible. | The sporadicity of the RRATs' radio emission led to immediate suggestions that they are related to other classes of traditionally ``radio-quiet'' neutron stars such as XDINSs and magnetars. \citet{popov2006} have shown that the implied birthrate of RRATs is more consistent with that of XDINSs than that of magnetars. As shown in the P-\.P diagram on the Figure~\ref{ppdot}, RRATs and XDINSs also have similar periods and period derivatives, implied ages and magnetic fields. However, the RRATs spin-down properties are also consistent with those of the normal pulsar population and X-ray observations of one RRAT~\citep{mmclaugh2007} reveal properties similar to those of both normal radio pulsars and XDINSs. \begin{figure} \includegraphics[scale=0.44]{p-pdot_bw.ps} \caption{P-\.P diagram. Lines on the top mark the period values for the objects with unknown yet period derivatives: seven RRATs (solid), one magnetar SGR~1627$-$41 and two candidates AX~J1845$-$03 and CXO~J164710.2$-$455216 (dashed), and one XDINS RX~J1605.3+3249 (longer dashed). } \label{ppdot} \end{figure} The RRATs are powerful sources of isolated radio bursts, and we have not detected such bursts, or any periodic emission, from the six XDINSs we have observed. Because the distances to the XDINSs are believed to be much smaller than those to the RRATs, we should have had high sensitivity to RRAT-like radio emission. However, our non-detection of such emission does not necessarily mean that there is no relationship between these two source classes. XDINSs may simply be ``radio-quiet'', but it is also likely that perhaps the narrow radio beams from these XDINSs are simply misaligned with our line-of-sight. It is possible that searches at lower frequencies, where radio emission beams are believed to be wider, may be more sensitive to radio emission from XDINSs. Indeed, Malofeev and co-authors~\citep{malofeev2005, malofeev2007} reported detection of radio emission from RX~J1308.6+2127 and RX~J2143.7+0654 at the low frequency of 111~MHz. On the other hand, XDINSs could have very steep spectral indices. If the detection of Malofeev and co-authors is real, our non-detection of radio emission from these two XDINSs at 820~MHz sets a lower limit on the spectral index of 3.6. Finally, it is possible that our non-detection of radio emission from these XDINSs is due to the large amount of contamination from RFI. Our search highlights the importance of improved excision algorithms for impulsive, broadband terrestrial interference. \begin{theacknowledgments} SZ thanks STFC (ex-PPARC) for support through an AF. The Robert C. Byrd Green Bank Telescope (GBT) is operated by the National Radio Astronomy Observatory which is a facility of the U.S. National Science Foundation operated under cooperative agreement by Associated Universities, Inc. \end{theacknowledgments} | 7 | 10 | 0710.1648 |
0710 | 0710.1683.txt | The lensing cross section of triaxial halos depends on the relative orientation between a halo's principal axes and its line of sight. Consequently, a lensing subsample of randomly oriented halos is not, in general, randomly oriented. Using an isothermal mass model for the lensing galaxies and their host halos, we show that the lensing subsample of halos that produces doubles is preferentially aligned along the lines of sight, whereas halos that produce quads tend to be projected along their middle axes. These preferred orientations result in different projected ellipticity distributions for quad, doubles, and random galaxies. We show that $\approx 300$ lens systems must be discovered to detect this effect at the $95\%$ confidence level. We also investigate the importance of halo shape for predicting the quad-to-double ratio and find that the latter depends quite sensitively on the distribution of the short-to-long axis ratio, but is otherwise nearly independent of halo shape. Finally, we estimate the impact of the preferred orientation of lensing galaxies on their projected substructure mass fraction, and find that the observed alignment between the substructure distribution and the mass distribution of halos result in a negligible bias. | Statistics of lensing galaxies have been used as cosmological and galaxy formation probes since early in the modern history of gravitational lensing \citep[][]{turneretal84}. Lensing rates can be used to constrain dark energy \citep[][]{fukugitaetal92, chae03,mitchelletal05, chae07,ogurietal07}, to probe the structure of lensing galaxies \citep[][]{keeton01d,kochanekwhite01, chae05}, and to probe galaxy evolution \citep[][]{chaemao03, ofeketal03,rusinkochanek05}. While the use of lensing statistics as a cosmological probe has had mixed success, particularly early on, it remains a unique probe with entirely different systematics from more traditional approaches. Consequently, lensing statistics are likely to remain a fundamental cross-check of our understanding of cosmology and galaxy evolution. One of the difficulties that confronts the study of lensing statistics is that, in general, the halo population that produces gravitational lenses can in fact be a highly biased subsample of the general halo population. For instance, it has long been known that while early type galaxies compose only $\approx 30\%$ of all luminous galaxies, the majority of lensing galaxies are in fact early type since these tend to be more massive and reside in more massive halos than their late counterparts. By the same token, lensing early type galaxies tend to have higher luminosity and velocity dispersions than non-lensing early type galaxies \citep[][]{moelleretal06, boltonetal06}. Overall, then, when interpreting lensing statistics, one ought to always remember that by selecting lensing galaxies one is automatically introducing an important selection effect that can significantly bias the distribution of any galaxy observable that has an impact on the lensing probabilities. Here, we consider one such source of bias, the triaxiality of galaxy halos.\footnote{Throughout this work, we will be using the term galaxy and halo more or less interchangeably. The reason for this is that we are primarily focused on the impact of halo triaxiality on the lensing cross section, and the latter depends only on the {\it total} matter density. Consequently, differentiating between halo and galaxy would only obfuscate presentation and introduce unnecessary difficulties. For instance, while modeling the total matter distribution as isothermal is a reasonable approximation, neither the baryons nor the dark matter by itself is isothermally distributed. Thus, it is much simpler to adopt an isothermal model, and refer to the baryons plus dark matter as a single entity, than to try to differentiate between the two. Likewise, when discussing triaxiality, what is important in this work is the triaxiality of the total matter distribution.} That halo triaxiality can have important consequences for lensing statistics has been known for several years. For instance, \citet[][]{ogurikeeton04} have shown that triaxiality can significantly enhance the optical depth of large image separation lenses. Similar conclusions have been reached concerning the formation of giant arcs by lensing clusters \citep[see e.g.][and references therein]{ogurietal03, rozoetal06c, hennawietal07}. Curiously, however, little effort has gone into investigating how observational properties of lensing galaxies can be different from those of the galaxy population as a whole due to the triaxial structure of galactic halos. This work addresses this omission. The first observable we consider is the projected axis ratio of lensing galaxies. Roughly speaking, given that non-zero ellipticities are needed in order to produce quad systems, one would generically expect lenses that lead to this image configuration to be more elliptical than the overall galaxy population. Likewise, lensing galaxies that produce doubles should, on average, be slightly more circular than a random galaxy. There can, however, be complications for these simple predictions due to halo triaxiality. For instance, given a prolate halo, projections along the long axis of the lens will result in highly concentrated, very circular profiles. Will the increase in Einstein radius of such projections compensate for the lower ellipticity of the system, implying most quads will be projected along their long axis, or will it be the other way around? Clearly, the relation between ellipticity and lensing cross sections is not straightforward once triaxiality of the lensing galaxies is taken into account, but it seems clear that there should be some observable difference between the ellipticity distribution of lensing galaxies and that of all early types. Interestingly, no such difference has been observed \citep[][]{keetonetal97,rusintegmark01}, which seems to fly in the face of our expectations \citep[though see also the discussion in][]{keetonetal98}. Is this actually a problem, or will a quantitative analysis show that the consistency of the two distributions is to be expected? Here, we explicitly resolve this question, and demonstrate that current lens samples are much too small to detect the expected differences. Having considered the ellipticity distribution of random and lensing galaxies, it is then a natural step to investigate the impact of halo triaxiality on predictions of the quad-to-double ratio. Specifically, it is well known that the quad-to-double ratio is sensitive to the ellipticity distribution of lensing galaxies \citep[][]{keetonetal97}, so if lensing can bias the distribution of ellipticities in lensing galaxies, then it should also affect the predicted quad-to-double ratios. This is an important point because it has been argued that current predictions for the quad-to-double ratio are at odds with observations. More specifically, the predicted quad-to-double ratio for the CLASS \citep[Cosmic Lens All-Sky Survey,][]{myersetal03,browneetal03} sample of gravitational lenses is too low relative to observations \citep[][]{rusintegmark01,hutereretal05}. Curiously, however, recent work on the quad-to-double ratio observed in the SQLS \citep[Sloan Digital Sky Survey Quasar Lens Search,][]{ogurietal06,inadaetal07}. suggests that the exact opposite is true for the latter sample, namely, theoretical expectations are too high relative to observations \citep[][]{oguri07}. In either case, it is of interest to determine how exactly does triaxiality affects theoretical predictions, especially since the aforementioned difficulties with the CLASS sample has led various authors to offer possibilities as to how one might boost the expected quad-to-double ratios. Specifically, one can boost the quad-to-double ration in the class sample either from the effect of massive satellite galaxies near the lensing galaxies \citep[][]{cohnkochanek04}, or through the large-scale environment of the lensing galaxy \citep[][]{keetonzabludoff04}. Clearly, we should determine whether halo triaxiality can be added to this list. This brings us then to the final problem we consider here, namely whether the substructure population of lensing galaxies is different from that of non-lensing galaxies. Specifically, we have argued that lensing galaxies will not be isotropically distributed in space. Since the substructure distribution of a dark matter halo is typically aligned with its parent halo's long axis \citep[][]{zentneretal05,libeskindetal05,agustssonbrainerd06,azzaroetal06}, it follows that the projected distribution of substructures for lensing galaxies may in fact be different for lensing halos than for non-lensing halos. Such an effect could be quite important given the claimed tension between the Cold Dark Matter (CDM) predictions for the substructure mass fraction of halos \citep[see][]{maoetal04} and their observed values \citep[][]{dalalkochanek02a,kochanekdalal04}. Likewise, such a bias would impact the predictions for the level of astrometric and flux perturbations produced by dark matter substructures in gravitational lenses \citep[][]{rozoetal06,chenetal07}. Here, we wish to estimate the level at which the projected substructure mass fraction of lensing halos could be affected due to lensing biasing. The paper is organized as follows: in section \ref{sec:biases} we derive the basic equations needed to compute how observable quantities will be biased in lensing galaxy samples due to halo triaxiality. Section \ref{sec:model} presents the model used in this work to quantitatively estimate the level of these biases, and discusses how lensing halos are oriented relative to the line of sight as a function of the halos' axes ratios. Section \ref{sec:axis} investigates the projected axis ratio distributions of lensing versus non-lensing galaxies, and demonstrates that present day lensing samples are too small to detect the triaxiality induced biases we have predicted. Section \ref{sec:ratio} discusses the problem of the quad to double ratio, and section \ref{sec:subs} demonstrates that halo triaxiality biases the projected substructure mass fraction in lensing halos by a negligible amount. Section \ref{sec:caveats} discusses a few of the effects we have ignored in our work and how these may alter our results, and finally section \ref{sec:summary} summarizes our work and presents our conclusions. %---------------------------------------------- | \label{sec:summary} The triaxial distribution of mass in galactic halos implies that the probability that a galaxy becomes a lens is dependent on the relative orientation of the galaxy's major axis to the line of sight. Consequently, a subsample of randomly oriented galaxies that act as strong lenses will {\it not} be randomly oriented in space. The relative orientation and the strength of the alignment depends on the shape of the matter distribution, and on the type of lens under consideration: prolate doubles have a high probability of being project along their long axis, whereas the distribution of oblate doubles is nearly isotropic. Prolate quads are most often projected along their middle axis, though the degree to which alignment occurs is not as strong as for prolate doubles. Interestingly, highly prolate quads are also more likely to be projected along their short axis than along the long axis, though this very quickly changes as halos become more triaxial and less prolate. Oblate quads strongly avoid projections along the short axis of the lens, but projections along the other two axis are almost equally likely. An important consequence of the differences in the distribution of halo orientations for quad lenses, double lenses, and the galaxy population as a whole is that the ellipticity distribution of these various samples must be different, even if the distribution of halo shapes is the same. Specifically, we predict that quad lenses are typically more elliptical than random galaxies, and that the ellipticity distribution of doubles is very slightly more circular than that of random galaxies. While current data do not show any indication of these trends, we have shown that $\approx 300\ (1,400)$ lenses are necessary to obtain a $2\sigma\ (5\sigma)$ detection of the effect. The fact that halo triaxiality affects the ellipticity distribution of lensing galaxies also means that halo triaxiality needs to be properly taken into account in lensing statistics. Consequently, we estimate how the biased lensing cross sections of galaxies depend on halo shape, and find that they are nearly independent of the halo shape parameter $T$. Instead, the mean biased cross section of a lens depends almost exclusive on the distribution on the short-to-long axis ratio $q_2$ (often denoted by $s$). Finally, given that the distribution of substructures in numerical simulations is observed to be preferentially aligned with the long axis of the host halos, we estimate how the preferred orientation of lensing galaxies affects their predicted substructure mass fraction. We find that biases due to non-isotropic distribution of halos relative to the line of sight have an insignificant impact on the mean substructure mass fraction of lensing galaxies. {\bf Acknowledgements: } ER would like to thank Christopher Kochanek for numerous discussions and valuable comments on the manuscript which have greatly improved both the form and content of this work. The authors would also like to thank to Emilio Falco for kindly providing the isophotal axis ratio data that was needed for producing Figure \ref{fig:cumq}, and to Charles Keeton for a careful reading of the manuscript. ER was funded by the Center for Cosmology and Astro-Particle Physics (CCAPP) at The Ohio State University. ARZ has been funded by the University of Pittsburgh, the National Science Foundation (NSF) Astronomy and Astrophysics Postdoctoral Fellowship program through grant AST 0602122, and by the Kavli Institute for Cosmological Physics at The University of Chicago. This work made use of the National Aeronautics and Space Administration Astrophysics Data System. % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- %--------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- % -------------------------------------------------------------------------------- | 7 | 10 | 0710.1683 |
0710 | 0710.5407_arXiv.txt | {Based on the analogy with non-minimal $SU(2)$ symmetric Wu-Yang mono\-pole with regular metric, the solution describing a non-minimal $U(1)$ symmetric Dirac monopole is obtained. In order to take into account the curvature coupling of gravitational and electromagnetic fields, we reconstruct the effective metrics of two types: the so-called associated and optical metrics. The optical metrics display explicitly that the effect of birefringence induced by curvature takes place in the vicinity of the non-minimal Dirac monopole; these optical metrics are studied analytically and numerically. } \email 1 {[email protected]}% \email 2 {[email protected]}% | The Dirac monopole as a specific static spherically symmetric solution to the minimal Einstein-Maxwell equations has became a subject of discussion in tens papers, reviews and books (see, e.g., \cite{Dirac,000,00,0,1,3,2}). The motion of massive and massless particles, which possess electric charge or are uncharged, is studied in detail (see, e.g., \cite{4,5} and references therein). In the paper \cite{BaZa07} we introduced and discussed the $SU(2)$ symmetric Wu-Yang monopole of the new type, namely, the non-minimal monopole with regular metric. Since the non-minimal Wu-Yang monopole is effectively Abe\-lian, it is naturally to consider the corresponding analog of that solution in the framework of non-minimal electro\-dynamics. Mention that non-minimal models with mag\-netic charge have been discussed earlier (see, e.g., \cite{Horn,MHS}), but the exact analytical regular solution obtained here as direct reduction to the $U(1)$ symmetry is the new one. The second novelty of the presented paper is the investigation of photon dynamics in the vicinity of the Dirac monopole accounting the curvature coupling of the gravitational and electromagnetic fields. In the presence of non-minimal interaction (induced by curva\-ture) the master equations for electromagnetic and gravitational fields in vacuum can be rewritten as the master equations in some effective anisotropic (quasi)me\-dium \cite{B1,BL05}. This means that two effective (optical) metrics can be introduced \cite{HehlObukhov,Perlick,BZ05}, so that the photon propagation in vacuum interacting with curvature is equivalent to the photon motion in the effective space-time with the first or second optical metric, depending on the photon polarization. Even if the real space-time has a regular metric, the optical metrics can be singular, admitting the interpretation in terms of the so-called ``trapped surfaces'' and ``inaccessible zones'' \cite{AG1,AG2}. We discuss this problem in Sect. 3. Numerical modeling of the photon orbits, presented in Sect. 4, supplement our conclusions. | Qualitative and numerical analysis of the photon orbits in the vicinity of non-minimal Dirac monopole with regular metric has demonstrated the following interesting features. 1. Propagation of the electromagnetic waves in the vicinity of non-minimal Dirac monopole is characterized by birefringence, induced by curvature, i.e., the phase velocities of waves depend on their polarization. Two different optical metrics should be introduced to describe two principal states of polarization. 2. The metric of the non-minimal Dirac monopole, obtained and discussed in this paper, is regular, thus, all the singularities of the optical metrics have a dynamic origin and are supported by the non-minimal (curvature induced) interaction of the gravitational and electromagnetic fields. 3. The points of self-intersection, the points of the closest approach, the reverse points, etc. in the photon trajectories can be recognized and catalogued for different combinations of the values of the impact parameter and the parameter of non-minimal coupling. We discussed here only the principal pictures. \Acknow This work was partially supported by the DFG through project No. 436RUS113/487/0-5. The authors are grateful to Claus L\"ammerzahl for valuable remarks. \small | 7 | 10 | 0710.5407 |
0710 | 0710.5294_arXiv.txt | {Thanks to remarkable progress, radial velocity surveys are now able to detect terrestrial planets at habitable distance from low-mass stars. Recently, two planets with minimum masses below 10~$M_{\oplus}$have been reported in a triple system around the M-type star Gliese 581. These planets are found at orbital distances comparable to the location of the boundaries of the habitable zone of their star. } {In this study, we assess the habitability of planets Gl~581c and Gl 581d (assuming that their actual masses are close to their minimum masses) by estimating the locations of the habitable-zone boundaries of the star and discussing the uncertainties affecting their determination. An additional purpose of this paper is to provide simplified formulae for estimating the edges of the habitable zone. These may be used to evaluate the astrobiological potential of terrestrial exoplanets that will hopefully be discovered in the near future.} {Using results from radiative-convective atmospheric models and constraints from the evolution of Venus and Mars, we derive theoretical and empirical habitable distances for stars of F, G, K, and M spectral types. } {Planets Gl~581c and Gl~581d are near to, but outside, what can be considered as the conservative habitable zone. Planet `c' receives 30\% more energy from its star than Venus from the Sun, with an increased radiative forcing caused by the spectral energy distribution of Gl~581. This planet is thus unlikely to host liquid water, although its habitability cannot be positively ruled out by theoretical models due to uncertainties affecting cloud properties and cloud cover. Highly reflective clouds covering at least 75\% of the day side of the planet could indeed prevent the water reservoir from being entirely vaporized. Irradiation conditions of planet `d' are comparable to those of early Mars, which is known to have hosted surface liquid water. Thanks to the greenhouse effect of CO$_2$-ice clouds, also invoked to explain the early Martian climate, planet `d' might be a better candidate for the first exoplanet known to be potentially habitable. A mixture of several greenhouse gases could also maintain habitable conditions on this planet, although the geochemical processes that could stabilize such a \textit{super-greenhouse} atmosphere are still unknown. } {} | The M-type star Gl~581 hosts at least 3 planets, which were detected using radial velocity measurements by Bonfils et al. \citeyearpar{2005A&A...443L..15B} (planet 'b') and Udry et al. \citeyearpar{2007A&A...469L..43U} (planets `c' and `d'). The properties of this star and its planets are given in Table~\ref{tab:gl581}. Before this discovery, only two exoplanets were known to have a minimum mass below 10~$M_{\oplus}$, which is usually considered as a boundary between terrestrial and giant planets, the latter having a significant fraction of their mass in an H$_2$-He envelope. The first one was GJ~876d, a very hot planet ($P\leq 2$~days) with a minimum mass of 5.9~M$_{\oplus}$ \citep{2005ApJ...634..625R}. The other one is OGLE-05-390L b, found to be a $\sim$5.5~M$_{\oplus}$ cold planet at 2.1~AU from its low-mass parent star thanks to a microlensing event \citep{2006Natur.439..437B,2006ApJ...651..535E}. Neither of these two planets is considered as habitable, even with very loose habitability criteria. In the case of Gl~581, and as already mentioned by Udry et al. (2007), the locations of planet `c' and `d' must be fairly close to the inner and outer edges, respectively, of the habitable zone (HZ). In this paper, we investigate the atmospheric properties that would be required to make the habitability of these planets possible. Because of its equilibrium temperature of $\sim$300~K when calculated with an albedo of 0.5, it has been claimed that the second planet of this system, Gl~581c, is potentially habitable (Udry et al. 2007), with climatic conditions possibly similar to those prevailing on Earth. After a brief discussion about the relationship between the equilibrium temperature and habitability, we summarize in this paper what are usually considered as the boundaries of the circumstellar HZ and the uncertainties on their precise location. In Sect.~\ref{sec:HZstars} we provide parameterizations to determine such limits as a function of the stellar luminosity and effective temperature. These can be used to evaluate the potential habitability of the terrestrial exoplanets that should soon be discovered. We then discuss the specific case of the system around Gl~581. \begin{table}[!t] \caption{Properties of the star Gl~581 and its 3 detected planets, from Udry et al. (2007).} \begin{center} \begin{tabular}{lllll} \hline Star & $T_{\rm eff}$ (K) &$M$/M$_{\odot}$ & $R$/R$_{\odot}$ & $L$/L$_{\odot}$ \\ \hline Gl~581 & 3200 &0.31 & 0.38 & 0.0135 \\ \hline &&& \\ &&& \\ \hline Planets & $a$~(AU) & $M_{\rm min}$/M$_{\oplus}$ & $R_{\rm min}$/R$_{\oplus}$ & stellar flux\\ & &$^{*}$ & $^{**}$ & $S/S_{0}$$^{***}$\\ \hline b & 0.041 & 15.6 & 2.2-2.6 & 8.1 \\ \rowcolor{lightgray} c & 0.073 & 5.06 & 1.6-2.0 & 2.55 \\ \rowcolor{lightgray} d & 0.253 & 8.3 & 1.8-2.2 & 0.21 \\ \hline \end{tabular} \end{center} The potential habitability of planets `c' and `d', highlighted in grey, is discussed in this paper. \\ $^{*}$ $M_{\rm min}=M \sin i$, where $i$ is the orbital inclination.\\ $^{**}$ Radius for a rocky and ocean planet, respectively \citep{2007Sotin,2007ApJ...665.1413V}.\\ $^{***}$ $S_0$ is the solar flux at 1 AU: 1360 W m$^{-2}$.\\ \label{tab:gl581} \end{table} | According to our present knowledge, based on available models of planetary atmospheres, and assuming that the actual masses of the planets are the minimum masses inferred from radial velocity measurements, Gl~581c is very unlikely to be habitable, while Gl~581d could potentially host surface liquid water, just as early Mars did. Because of the uncertainties in the precise location of the HZ boundaries, planets at the edge of what is thought to be the HZ are crucial targets for future observatories able to characterize their atmosphere. At the moment, our theory of habitability is only confirmed by the divergent fates of Venus and the Earth. We will have to confront our models with actual observations to better understand what makes a planet habitable. The current diversity of exoplanets (planets around pulsars, hot Jupiters, hot Neptunes, super-Earths, etc) has already taught us that Nature has a lot more imagination when building a variety of worlds than we expected from our former models inspired by the Solar System. It is obvious that the idealized model of a habitable planet atmosphere, where the two important constituents are CO$_2$ and H$_2$O, CO$_2$ being controlled by the carbonate-silicate cycle, is likely to represent only a fraction of the diversity of terrestrial planets that exist at habitable distances from their parent star. As an example, planets fully covered by an ocean may be common, either because they are richer in water than Earth or because the distribution between surface and mantle water is different, or perhaps simply because, for a given composition, the mass-to-surface ratio and thus the water-to-surface ratio increases with the planetary mass, as noted by Lissauer \citeyearpar{1999Lissauer}. Without emerged continents, it is not at all clear that the carbonate-silicate cycle could operate. The planets around Gl~581 can fall into this category since they are significantly more massive than the Earth (especially the $>$8~$M_{\oplus}$ planet Gl~581d) and also because they may have started their formation in the outer and more water-rich region of the protoplanetary disk. Darwin/TPF-I and TPF-C could eventually reveal what the actual properties of the atmosphere of Gl~581c and Gl~581d are. From their thermal light curves we could infer if a thick atmosphere is making the climate more or less uniform on both the day and night hemispheres of these planets, despite a (nearly?) synchronized rotation \citep{2004ASPC..321..170S}. Visible and mid-IR water vapor bands could be searched in the atmosphere of Gl~581d to confirm its habitability. Mid-IR spectra of this planet could also reveal other greenhouse gases at work. Spectral observations of Gl~581c could potentially distinguish between a Venus-like atmosphere dominated by CO$_2$ or an H$_2$O-rich atmosphere. The detection of O$_2$ on this planet would generate a fascinating debate about its possible origin: as either a leftover of H$_2$O photolysis and H escape or a biological release. There is certainly no doubt that Gl~581c and Gl~581d are prime targets for exoplanet characterization missions. | 7 | 10 | 0710.5294 |
0710 | 0710.0364_arXiv.txt | We explore the cosmological consequences of Modified Gravity (MOG), and find that it provides, using a minimal number of parameters, good fits to data, including CMB temperature anisotropy, galaxy power spectrum, and supernova luminosity-distance observations without exotic dark matter. MOG predicts a bouncing cosmology with a vacuum energy term that yields accelerating expansion and an age of $\sim$13 billion years. | The preferred model of cosmology today, the $\Lambda$CDM model, provides an excellent fit to cosmological observations, but at a substantial cost: according to this model, {\em about 95\% of the universe is either invisible or undetectable, or possibly both} \cite{Komatsu2008}. This fact provides a strong incentive to seek alternative explanations that can account for cosmological observations without resorting to dark matter or Einstein's cosmological constant. For gravitational theories designed to challenge the $\Lambda$CDM model, the bar is set increasingly higher by recent discoveries. Not only do such theories have to explain successfully the velocity dispersions, rotational curves, and gravitational lensing of galaxies and galaxy clusters, the theories must also be in accord with cosmological observations, notably the acoustic power spectrum of the cosmic microwave background (CMB), the matter power spectrum of galaxies, and the recent observation of the luminosity-distance relationship of high-$z$ supernovae, which is seen as evidence for ``dark energy''. Modified Gravity (MOG) \citep{Moffat2006a} has been used successfully to account for galaxy cluster masses \citep{Brownstein2006b}, the rotation curves of galaxies \citep{Brownstein2006a}, velocity dispersions of satellite galaxies \citep{Moffat2007}, and globular clusters \citep{Moffat2007a}. It was also used to offer an explanation for the Bullet Cluster \citep{Brownstein2007} without resorting to cold dark matter. Remarkably, MOG also meets the challenge posed by cosmological observations. In this paper, it is demonstrated that MOG produces an acoustic power spectrum, a matter power spectrum, and a luminosity-distance relationship that are in good agreement with observations, and require no dark matter nor Einstein's cosmological constant. In the arguments presented here, we rely on simplified analytical calculations. We are not advocating these as substitutes for an accurate numerical analysis. However, a thorough numerical analysis requires significant time and resources; before these are committed, it is useful to be able to demonstrate if a theory is viable, and if the additional effort is warranted. In the next section, we review the key features of MOG. This is followed by sections presenting detailed calculations for the luminosity-distance relationship of high-$z$ supernovae, the acoustic power spectrum of the CMB, and the galaxy power spectrum. A concluding section summarizes our results and maps out future steps. | In this paper, we demonstrated how Modified Gravity can account for key cosmological observations using a minimum number of free parameters. Although MOG permits the running of its coupling and scaling constants with time and space, we made very little use of this fact. Throughout these calculations, we used consistently the value of $\alpha\simeq 19$ for the MOG coupling constant, consistent with a flat universe with $\Omega_b\simeq 0.05$ visible matter content, no dark matter, nor Einstein's cosmological constant. In nearly all cosmological calculations, we set the MOG scaling constant $\mu$ to the inverse of the radius of the visible universe, which is a natural choice. The only exception is the mass power spectrum calculation, where the scaling constant enters in conjunction with the wave number $k$, and describes gravitational interactions between nearby concentrations of matter, not on the cosmological scale. The theory requires {\em no other parameters} to obtain the remarkable fits to data that have been demonstrated here. At all times, $\lim\limits_{r\rightarrow 0}G=G_N$, i.e., the effective gravitational constant at short distances remains Newton's constant of gravitation. For this reason, the predictions of MOG are {\em not contradicting our knowledge of the processes of the initial nucleosynthesis}, taking place at redshifts of $z\simeq 10^{10}$, since the interactions take place over distance scales that are much shorter than the horizon scale. Our calculations relied on analytical approximations. This is dictated by necessity, not preference. We recognize that numerical methods, including high-accuracy solutions of coupled systems of differential equations, as in {\tt CMBFAST} \citep{Seljak1996}, or $N$-body simulations, can provide superior results, and may indeed help either to confirm or to falsify the results presented here. Nevertheless, our present work demonstrates that at the very least, MOG provides a worthy alternative to $\Lambda$CDM cosmology. | 7 | 10 | 0710.0364 |
0710 | 0710.2157_arXiv.txt | We investigate the incidence of major mergers creating ${\rm M}_{\rm star}>10^{11} {\rm M}_{\sun}$ galaxies in the dense environments of present-day groups and clusters more massive than ${\rm M}_{\rm halo}=2.5\times10^{13}{\rm M}_{\sun}$. We identify 38 pairs of massive galaxies with mutual tidal interaction signatures selected from $>5000$ galaxies with ${\rm M}_{\rm star}\geq5\times10^{10} {\rm M}_{\sun}$ that reside in a halo mass-limited sample of 845 groups. We fit the images of each galaxy pair as the line-of-sight projection of symmetric models and identify mergers by the presence of residual asymmetric structure associated with both progenitors, such as nonconcentric isophotes, broad and diffuse tidal tails, and dynamical friction wakes. At the resolution and sensitivity of the SDSS, such mergers are found in 16\% of the high-mass, galaxy-galaxy pairs with $\leq1.5$ $r$-band magnitude differences and $\leq30$ kpc projected separations. Relying on automated searches of major pairs from the SDSS spectroscopic galaxy sample will result in missing 70\% of these mergers owing to spectroscopic incompleteness in high-density regions. We find that 90\% of these mergers are between two nearly equal-mass progenitors with red-sequence colors and centrally-concentrated morphologies, in agreement with numerical simulations that predict that an important mechanism for the formation of massive elliptical galaxies is the dissipationless (gas-poor or so-called dry) major merging of spheroid-dominated galaxies. We identify seven additional ${\rm M}_{\rm star}>10^{11} {\rm M}_{\sun}$ mergers with disturbed morphologies and semi-resolved double nuclei. Mergers at the centers of massive groups are more common than between two satellites, but both types are morphologically indistinguishable and we tentatively conclude that the latter are likely located at the dynamical centers of large subhalos that have recently been accreted by their host halo, rather than the centers of distinct halos seen in projection. We find that the frequency of central and satellite merging diminishes with group mass in a manner that is consistent with dynamical friction. Based on reasonable assumptions, the centers of these massive halos are gaining stellar mass at a rate of 1--9\% per Gyr on average. Compared to the merger rate for the overall population of luminous red galaxies, we find that the rate is 2--9 times greater when restricted to these dense environments. Our results imply that the massive end of the galaxy population continues to evolve hierarchically at a measurable level, and that the centers of massive groups are the preferred environment for the merger-driven assembly of massive ellipticals. | Understanding the formation of the most-massive galaxies (${\rm M}_{\rm star}>10^{11}{\rm M}_{\sun}$) remains an important challenge in astrophysics. The tip of the stellar mass function is dominated by elliptical galaxies with intrinsically spheroidal mass distributions that are supported by anisotropic stellar motions \citep{kormendy96,burstein97}. Numerical simulations have long demonstrated that ``major'' mergers between smaller galaxies of comparable mass could produce the observed shapes and dynamics of ellipticals \citep{toomre77,barnes96,naab03,cox06}. Moreover, massive ellipticals are found in greater abundance in high-density structures like large groups and clusters of galaxies \citep[e.g.,][]{dressler80a,postman84,hashimoto99,smith05}, which naturally grow through the hierarchical merging of dark-matter halos over cosmic time as expected in the $\Lambda$CDM cosmological model \citep{blumenthal84,davis85,cole00}. There is, therefore, a clear expectation for galaxy-galaxy and halo-halo merging to be physically linked \citep{maller06,hopkins06b,delucia07}. Indeed, modern galaxy formation models predict that massive ellipticals form by major dissipationless (so-called ``dry'') merging of likewise spheroidal and gas-poor progenitors \citep{boylan06,naab06a}, that a large fraction of today's massive ellipticals had their last major merger since redshift $z=0.5$ \citep[e.g.,][]{delucia06}, and that the most-massive systems form at the centers of large dark-matter halos \citep{dubinski98,aragon98}. Yet, direct evidence for the major-merger assembly of massive galaxies at present times has been lacking, and finding such systems is needed to place constraints on their rates, progenitor properties, and environmental dependencies. To this end we look for close pairs of massive interacting galaxies within a complete and well-defined sample of over 5000 galaxies with $z\leq0.12$ and ${\rm M}_{\rm star}\geq5\times10^{10}{\rm M}_{\sun}$, selected from galaxy groups in the Sloan Digital Sky Survey (SDSS) with dark-matter halo masses above ${\rm M}_{\rm halo}=2.5\times10^{13}{\rm M}_{\sun}$. Ellipticals galaxies make up the bulk of the massive end of the red-sequence population with optical colors indicative of their non-star-forming and old stellar nature. Despite a quiet star-formation history over the last 6--8 billion years \citep{bell05a}, the total stellar mass density on the red sequence has roughly doubled over this interval \citep{bell04b,blanton06,borch06,faber07,brown07} and now accounts for more than half of the present-day budget \citep{hogg02,bell03b}, providing strong observational evidence for the ongoing hierarchical growth of the massive galaxy population. These results were derived from red galaxy number densities over a wide range of stellar masses above and below $10^{11}{\rm M}_{\sun}$. Owing to the scarcity of the highest-mass galaxies, cosmic variance, and systematic uncertainties in stellar mass estimates, any increase in the number density of ${\rm M}_{\rm star}>10^{11}{\rm M}_{\sun}$ galaxies is poorly constrained, resulting in controversy over whether this population has continued to grow slowly \citep[e.g.,][]{brown07} or has been effectively static \citep[e.g.,][]{scarlata07}, since $z\sim1$. Besides number density evolution, mergers of sufficiently massive galaxies could provide a more clear indication for some continued stellar mass growth in the high-mass galaxy population. The existence of a handful of massive red mergers over the redshift interval $0.1<z<0.9$ \citep{vandokkum99,tran05,bell06a,lotz06,rines07} proves that the growth is non-zero at high stellar masses and implies that this mechanism does contribute to the assembly of galaxies at the top of the food chain. Yet, the importance of this process and the related rate of mass growth are highly uncertain given the tiny samples over this large cosmic time interval. Indirect measures such as the presence of faint tidal debris or shells around many local massive ellipticals \citep{vandokkum05,mihos05}, the isophotal properties of giant ellipticals \citep{kang07}, the lack of evolution of the stellar mass-size relation of red spheroids since $z=1$ \citep{mcintosh05b}, and the lack of morphological evolution on the red sequence since $z=0.7$ \citep{bell04a} provide a variety of limits to the importance of dissipationless mergers. Perhaps the most powerful method for obtaining estimates for the stellar mass growth rate via major merging is based on small-scale clustering statistics that provide an accurate measurement of close pair frequencies in real space \citep{masjedi06,bell06b,masjedi07}. However, this method likely yields an overestimate of the merger frequency because it assumes that all close pairs will merge. All estimates of merger-driven growth rates are limited by uncertainties in the time interval over which a pair will merge, and over what duration an object could be identified as interacting. \citet{masjedi07} find a very small growth rate (1--2\% per Gyr) at $z\sim0.25$ for major mergers involving at least one progenitor drawn from the SDSS Luminous Red Galaxy sample \citep[LRG,][]{eisenstein01}; LRGs have typical masses of several times $10^{11}{\rm M}_{\sun}$. To date there remains no direct evidence of ongoing merger-driven assembly of massive galaxies at $z<0.1$, and the LRG result implies that this formation process is no longer important. These facts motivate a thorough search for the existence/nonexistence of ongoing examples in the present-day universe. While the aforementioned statistical method for finding close physical pairs is powerful, it does not isolate actual merging systems and thus provides no information on the progenitor properties of massive merger remnants. Recent numerical simulations and models make a range of predictions regarding the progenitor morphologies at the time of the last major merger \citep{khochfar03,naab06a,kang07}, yet robust observational constraints are missing for ${\rm M}_{\rm star}>10^{11} {\rm M}_{\sun}$ systems. Many studies have identified major-merger candidates by either close pairs \citep{carlberg94,patton00,carlberg00,patton02,bundy04,lin04} or disturbed morphologies \citep{lefevre00,conselice03,lotz06}, but these samples mostly contain major mergers between lower-luminosity galaxies that tend to be gas-rich spiral disks. Numerical simulations show that such dissipative merging of disk galaxies will not produce massive pressure-supported ellipticals \citep[e.g.,][]{naab06d}. As mentioned above, only circumstantial evidence and a small number of red galaxy pairs with $z<0.9$ support the existence of mergers likely to produce massive ellipticals. Our understanding of the progenitors is therefore very limited. Here we present a thorough census of 38 massive merger pairs from SDSS, providing an order-of-magnitude increase in the number of such detections at $z<0.5$ and allowing an improved understanding of their progenitor properties. While many estimates of major merger rates are found in the literature, to date no measure of the environmental dependence of merger-driven mass growth has been attempted. In the standard cosmological model, there is a trade off between the expansion of the universe and the gravitational collapse of dark and luminous matter. Therefore, the rate at which stellar mass is assembled at the centers of the largest dark matter halos over recent cosmic history is a fundamental aspect of the ongoing formation of large-scale structure, and the rate that high-mass galaxies form by mergers as a function of halo mass constrains galaxy formation theories. Some theories predict that the mergers producing massive ellipticals occur preferentially in groups rather than in high-density cluster or low-density field environments because the smaller velocity dispersions allow more galaxy interactions \citep{cavaliere92}; also dynamical friction is more efficient in lower-mass halos \citep[e.g.][]{cooray05d}. Others predict that the brightest cluster galaxies (BCGs) grow by hierarchical merging (``galactic cannibalism'') at the centers of the dark-matter potential wells of large clusters \citep{ostriker75,merritt85,dubinski98,cooray05d}. A handful of low-redshift BCGs show multiple nuclei suggesting cannibalism in the form of multiple minor mergers \citep{lauer88}, but there are no observations of major mergers at the centers of clusters. In this paper we make use of the statistically large SDSS group catalog \citep{yang05a,weinmann06a} to show that major mergers occur in present-day dense environments, and to explore the halo-mass dependence and central/satellite identity of merger-driven massive galaxy assembly. Throughout this paper we calculate comoving distances in the $\Lambda$CDM concordance cosmology with $\Omega_{\rm m} = 0.3$, $\Omega_{\Lambda} = 0.7$, and assume a Hubble constant of $H_0 = 70 $\,km\,s$^{-1}$\,Mpc$^{-1}$. SDSS magnitudes are in the AB system. | \label{sec:Disc} We find the first direct observational evidence for an important population of galaxy-galaxy mergers with total stellar masses above $10^{11} {\rm M}_{\sun}$ in the local universe. These objects provide an unprecedented census of the progenitor properties for the merger-driven assembly of high-mass galaxies, which we compare to recent predictions from numerical models of galaxy formation and evolution. Moreover, the existence of these mergers prove that a measurable amount of stellar mass growth continues in the massive galaxy population at present times, and we compare estimates based on this sample with other estimates in the literature. Finally, we have identified mergers restricted to reside in large SDSS groups and clusters with $z\leq0.12$, thus allowing the first constraints on the halo-mass dependencies of recent massive merger activity. While it is well-established that massive galaxies are more common in such high-density environments, we are missing much more than 50\% of the population with ${\rm M}_{\rm star}<4\times10^{11} {\rm M}_{\sun}$ in the local volume, as Table \ref{mstar_counts} shows. Therefore, we must keep this caveat in mind when interpreting the conditions for which our results hold. In an upcoming study, we are examining the role of major mergers as a function of stellar mass over the full range of environments hosting galaxies more massive than ${\rm M}_{\rm star}=5\times10^{10}{\rm M}_{\sun}$. \subsection{Massive Merger Progenitors: Observations Meet Theories} \label{sec:Disc1} Establishing the luminosity dependence of elliptical (E) galaxy properties \citep{davies83,bender88,bender92} set the stage for theories regarding the types of merger progenitors that would produce the characteristics of low and high-mass early-type galaxies (ETGs)\footnote{The distinction between elliptical and early-type galaxies is often blurred in the literature. We consider Es to be a morphological subset of ETGs, which are concentrated and spheroid-dominated systems including Es, lenticulars (S0s), and Sa spirals. When referencing other authors we remain faithful to their choice of nomenclature.} galaxies \citep{bender92,kormendy96,faber97}. We concentrate on modern numerical simulations and semi-analytic models that attempt to reproduce the kinematic, photometric, and structural properties observed in massive Es through major merging \citep{naab99,naab03,khochfar03,khochfar05,naab06a,boylan06,kang07}. For this discussion we make the straight-forward assumption that the major mergers that we have identified will produce remnants that are {\it not unlike} the ${\rm M}_{\rm star}>10^{11} {\rm M}_{\sun}$ galaxy population already in place. We can only guess at remnant properties (see Fig. \ref{fig:cm.remn}), but in general, massive galaxies on the red-sequence are typically early-type. As we show in Figure \ref{fig:prog.mass}, the progenitor masses are comparable for the most part, and quantitatively consistent with the LRG-LRG merger mass spectrum from \citet{masjedi07} under the assumption that companions merge on dynamical friction time scales. $N$-body simulations \citep[e.g.][]{naab99} have long shown that ${\rm M}_1/{\rm M}_2\approx1$ are necessary to produce the lack of significant rotation observed in massive Es. Yet, a near unity mass ratio alone is not sufficient to produce the predominance of boxy and anisotropic Es found at high luminosity \citep{naab03,naab06a}. To match the decreasing fraction of rotational support and increasing fraction of boxiness in more luminous Es, the role of gas dissipation must be significantly reduced at high masses \citep{bender92,khochfar05,naab06d,kang07}, and recent ETG-ETG merger simulations have demonstrated this numerically \citep{naab06a}. Figures \ref{fig:prog.type} and \ref{fig:cm.prog} show that 90\% of the progenitors in this study have concentrated light profiles and red-sequence colors, both common attributes of ETGs, with little or no cold gas content. In addition, the tidal signatures of the bulk of these massive mergers (see Figs. \ref{fig:dpairs1} \& \ref{fig:dpairs2}) match those of observed \citep{bell06a} and simulated \citep{naab06a} major dissipationless (or gas-poor) merging of ETGs. Thus, our sample represents a more than order-of-magnitude increase in the number of such known systems with $z<0.2$, and demonstrates that dissipationless merging is indeed an important channel for the formation of massive galaxies. Finally, we compare the observed high fraction of ETG-ETG mergers ($f_{\rm ETG-ETG}=0.9$) with several semi-analytic predictions. Recall that we have looked for signs of interaction in $>200$ major pairs from a total sample of $>5000$ massive galaxies (i.e., sampM), yet only 10\% of the 38 mergers we identify could possibly form a ${\rm M}_{\rm star}>10^{11}{\rm M}_{\sun}$ remnant by other than an ETG-ETG merger. The progenitor morphologies of this study best match the predictions of \citet{khochfar03}, who find $f_{\rm E-E}=0.75$ for the last major merger of $4L^{\ast}$ remnants, independent of environment. We find much larger ETG-ETG fractions than \citet{naab06a} who predict only 20-35\% (also independent of environment) over the estimated mass range of our merger remnants ($11.1<\log_{10}{({\rm M}_{\rm star}/{\rm M}_{\sun})}<11.7$), and \citet{kang07} who predict $f_{\rm ETG-ETG}<0.1$ for $\log_{10}{({\rm M}_{\rm star}/{\rm M}_{\sun})}>11$. We note that these predicted progenitor morphologies for present-day Es are based on the final major mergers that could occur over a large redshift range out to $z\sim1$, which could be different in nature to those that occur in the short time interval that we observe. Moreover, we focus on high-density environments known to have very few massive late-type (blue) galaxies \citep{butcher78a}, which might explain the low number of ``mixed'' (early-late or elliptical-spiral) mergers that we find. Hence, for these models to be consistent with our data, either (1) the $f_{\rm ETG-ETG}$ of present-day major mergers depends on halo mass (i.e., environment), or (2) the relative importance of major mixed mergers has decreased significantly since $z=1$. \subsection{Estimating Stellar Mass Accretion Rates} The existence of massive dissipationless mergers at low redshift is direct observational evidence that the growth of ${\rm M}_{\rm star}>10^{11}{\rm M}_{\sun}$ galaxies continues at present times in agreement with many cosmologically-motivated simulations \citep{khochfar05,delucia06,kaviraj07,kang07}. Moreover, even under conservative assumptions that limit the amount of companion mass that is added to massive CEN galaxies, all of our sample will still result in remnants with ${\rm M}_{\rm star}>10^{11}{\rm M}_{\sun}$. Previously, the observational evidence for recent merger-based assembly of $z\sim0$ massive Es was limited to luminous/massive galaxy clustering statistics \citep{masjedi06,bell06b,masjedi07} or post-merger signatures that cannot distinguish between minor and major merging; e.g., tidal shells \citep{malin83}, fine structure \citep{schweizer92}, faint tidal features \citep{vandokkum05,mihos05}, or kinematic/photometric properties \citep[e.g.,][]{kang07}. With the merger sample presented here we can quantify directly the amount of growth, occurring in dense environments, at the high-mass end of the stellar mass function. Going from the observed merger counts to an inferred merger rate is limited mostly by the uncertainty in the merger timescale ($t_{\rm merg}$) that one assumes. Numerical models show that the time interval for two galaxies to interact and finally merge into a single remnant depends critically on the orbital parameters, progenitor mass ratios and densities, and the degree to which the merger is dissipationless. For major mergers of massive galaxies a number of different $t_{\rm merg}$ have been put forth in the literature based on simple orbital timescale arguments. For example, \citet{masjedi06} derived a reasonable lower limit of $t_{\rm merg}=0.2$ Gyr for a close ($d_{1,2}=10$ kpc) pair of LRG galaxies with a velocity dispersion of $\sigma=200$ \kms. Naturally, bound pairs with $d_{1,2}>10$ kpc separation will take longer to merge. \citet{bell06b} made a similar calculation for somewhat less-massive galaxies typically separated by $d_{1,2}=15$ kpc and estimated $t_{\rm merg}=0.4$ Gyr and argued for at least a factor of two uncertainty in this time. The mergers in this study have an average projected separation of 15.5 kpc (see Fig. \ref{fig:pair.mergVnon}), so in what follows, we adopt $t_{\rm merg}=0.4^{+0.4}_{-0.2}$ Gyr with conservative error bars that encompass the range of uncertainties discussed in the literature. Here, we compute the rate of stellar mass accretion by major merging onto massive galaxies in large groups. First, we find that the total mass accreted onto the centers of the $N_{\rm CEN}=845$ halos that we study is $\sum{ f{\rm M}_{{\rm s},i}}=3.9(3.5)\times10^{12} {\rm M}_{\sun}$, if we include (exclude) the four SAT-SAT mergers at their host's dynamical center (see \S \ref{sec:cenmerging}). ${\rm M}_{{\rm s},i}$ is the stellar mass of the secondary (SAT) galaxy in the $i^{\rm th}$ CEN-SAT merger, and $f$ is the fraction of ${\rm M}_{{\rm s},i}$ that winds up as part of the CEN galaxy. The rate of stellar mass buildup per massive CEN galaxy is therefore \begin{equation} \dot{{\rm M}}_{\rm CEN} = \frac{\sum{ f{\rm M}_{{\rm s},i} }}{N_{\rm CEN}} \times \frac{1}{t_{\rm merg}} , \label{eq:4} \end{equation} or between $1.0^{+1.0}_{-0.5}\times10^{10}{\rm M}_{\sun}{\rm Gyr}^{-1}$ and $1.2^{+1.1}_{-0.6}\times10^{10}{\rm M}_{\sun}{\rm Gyr}^{-1}$, depending on which sample of CEN-SAT mergers that we consider. The lopsided error bars result from the range of accretion rates for $t_{\rm merg}=0.4^{+0.4}_{-0.2}$ Gyr, as described above. If we divide all of these accretion rates by $2.69\times10^{11} {\rm M}_{\sun}$, the average stellar mass of the 845 CEN galaxies in this study, we find that each CEN is growing by 1--9\% per Gyr. Finally, these values can be decreased by assuming $f<1$ in (\ref{eq:4}), but as we discuss in \S \ref{sec:cmplane}, $f=0.5$ represents a likely lower limit. Likewise, the total stellar mass accreted onto all galaxies in sampM is $\sum{ f{\rm M}_{{\rm s},i}}+\sum{ {\rm M}_{{\rm s},j}}=5.1\times10^{12} {\rm M}_{\sun}$, where ${\rm M}_{{\rm s},j}$ is the mass of the secondary (SAT) galaxy in the $j^{\rm th}$ SAT-SAT merger. Therefore, the growth per ${\rm M}_{\rm star}\geq5\times10^{10} {\rm M}_{\sun}$ galaxy in high-mass groups is \begin{equation} \dot{{\rm M}}_{(\geq5\times10^{10} {\rm M}_{\sun})} = \frac{\sum{ f{\rm M}_{{\rm s},j} } + \sum{ {\rm M}_{{\rm s},j}} }{ (N_{\rm CEN}+N_{\rm SAT}-N_{\rm s,sampM})} \times \frac{1}{t_{\rm merg}} , \end{equation} where $N_{\rm s,sampM}=12$ is the number of secondary SAT galaxies in sampM that are involved in major mergers and must be subtracted to avoid double counting. We find $\dot{{\rm M}}_{(\geq5\times10^{10} {\rm M}_{\sun})}=2.4^{+2.4}_{-1.2}\times10^{9}{\rm M}_{\sun}{\rm Gyr}^{-1}$; if we assume $f=0.5$ for CEN-SAT mergers only we find $\dot{{\rm M}}_{(\geq5\times10^{10} {\rm M}_{\sun})}=1.6^{+1.5}_{-0.7}\times10^{9}{\rm M}_{\sun}{\rm Gyr}^{-1}$. Given that the average stellar mass of sampM galaxies is $1.04\times10^{11} {\rm M}_{\sun}$, we find that every massive galaxy is growing by 1--5\% per Gyr. Even though SAT-SAT mergers may occur as frequently as CEN-SAT mergers in these massive groups, the centers are where much of the mass growth takes place. It is clear from Figure \ref{fig:prog.mass} that mostly only ${\rm M}_{\rm star}>10^{11} {\rm M}_{\sun}$ galaxies build up in mass by major mergers in groups with ${\rm M}_{\rm halo}>2.5\times10^{13}{\rm M}_{\sun}$. In contrast, we find few mergers among the $5\times 10^{10}<{\rm M}_{\rm star}<10^{11} {\rm M}_{\sun}$ galaxies in these high-mass groups, which make up the bulk (60\%) of sampM. This suggests that if major merging is playing an important role in the strong mass growth observed on the red sequence below M* \citep{bell04b,blanton06,borch06,faber07,brown07}, it is occurring in lower-mass groups than we study here. Rather than mass growth rates we can use the same line of reasoning to estimate massive galaxy-galaxy merging rates of $(21+4)/845/t_{\rm merg}=0.074^{+0.074}_{-0.037}{\rm Gyr}^{-1}$ for CEN-SAT and $(38+7)/(845+4531-12)/t_{\rm merg}=0.021^{+0.021}_{-0.011}{\rm Gyr}^{-1}$ for all galaxies in sampM. For these estimates we included the seven additional major mergers (4 CEN, 3 SAT) we identified by their highly-disturbed appearance. \citet{masjedi06} found a strict upper limit to the LRG-LRG rate of only $0.006{\rm Gyr}^{-1}$. We estimate that LRGs have a stellar mass range of $11.4<\log_{10}{({\rm M}_{\rm star}/{\rm M}_{\sun})}<12.0$, based on typical red-sequence colors and luminosities between $4L^{\ast}$ and $25L^{\ast}$. Within these mass limits, we find a merger rate of $5/462/t_{\rm merg}=0.027^{+0.027}_{-0.014}{\rm Gyr}^{-1}$ on the red sequence, or 2--9 times the LRG-LRG rate. In Table \ref{mstar_counts}, we show that the high-mass groups that we study contain $>70\%$ of the very-massive, red galaxy population in the $z\leq0.12$ volume of DR2, with the vast majority being CENs. Yet, the same group selection contains only 30\% of the population of $11.4<\log_{10}{({\rm M}_{\rm star}/{\rm M}_{\sun})}<11.6$ systems. These numbers show that a significant portion of the local counterparts to LRGs are found in groups with ${\rm M}_{\rm halo}<2.5\times10^{13}{\rm M}_{\sun}$. Therefore, we conclude that LRG-LRG merging occurs more frequently in the more massive groups. | 7 | 10 | 0710.2157 |
0710 | 0710.1542_arXiv.txt | {The initial-final mass relationship of white dwarfs, which is poorly constrained, is of paramount importance for different aspects in modern astrophysics. From an observational perspective, most of the studies up to now have been done using white dwarfs in open clusters.} {In order to improve the initial-final mass relationship we explore the possibility of deriving a semi-empirical relation studying white dwarfs in common proper motion pairs. If these systems are comprised of a white dwarf and a FGK star, the total age and the metallicity of the progenitor of the white dwarf can be inferred from the detailed analysis of the companion.} {We have performed an exhaustive search of common proper motion pairs containing a DA white dwarf and a FGK star using the available literature and crossing the SIMBAD database with the Villanova White Dwarf Catalog. We have acquired long-slit spectra of the white dwarf members of the selected common proper motion pairs, as well as high resolution spectra of their companions. From these observations, a full analysis of the two members of each common proper motion pair leads to the initial and final masses of the white dwarfs.} {These observations have allowed us to provide updated information for the white dwarfs, since some of them were misclassified. In the case of the DA white dwarfs, their atmospheric parameters, masses, and cooling times, have been derived using appropriate white dwarf models and cooling sequences. From a detailed analysis of the FGK stars spectra we have inferred the metallicity. Then, using either isochrones or X-ray luminosities we have obtained the main-sequence lifetime of the progenitors, and subsequently their initial masses.} {This work is the first one in using common proper motion pairs to improve the initial-final mass relationship, and has also allowed to cover the poorly explored low-mass domain. As in the case of studies based on white dwarfs in open clusters, the distribution of the semi-empirical data presents a large scatter, which is higher than the expected uncertainties in the derived values. This suggests that the initial-final mass relationship may not be a single-valued function.} | White dwarfs are the final remnants of low- and intermediate-mass stars. About 95\% of main-sequence stars will end their evolutionary pathways as white dwarfs and, hence, the study of the white dwarf population provides details about the late stages of the life of the vast majority of stars. Since white dwarfs are long-lived objects, they also constitute useful objects to study the structure and evolution of our Galaxy (Liebert et al.~2005a; Isern et al. 2001). For instance, the initial-final mass relationship (IFMR), which connects the properties of a white dwarf with those of its main-sequence progenitor, is of paramount importance for different aspects in modern astrophysics. It is required as an input for determining the ages of globular clusters and their distances, for studying the chemical evolution of galaxies, and also to understand the properties of the Galactic population of white dwarfs. Despite its relevance, this relationship is still poorly constrained, both from the theoretical and the observational points of view. The first attempt to empirically determine the initial-final mass relationship was undertaken by \cite{wei77}, who also provides a recent review on this subject (Weidemann 2000). It is still not clear how this function depends on the mass and metallicity of the progenitor, its angular momentum, or the presence of a strong magnetic field. The total age of a white dwarf can be expressed as the sum of its cooling time and the main-sequence lifetime of its progenitor. The latter depends on the metallicity of the progenitor of the white dwarf, but it cannot be determined from observations of single white dwarfs. This is because white dwarfs have such strong surface gravities that gravitational settling operates very efficiently in their atmospheres, and any information about their progenitors (e.g. metallicity) is lost in the very early evolutionary stages of the cooling track. Moreover, the evolution during the AGB phase of the progenitors is essential in determining the size and composition of the atmospheres of the resulting white dwarfs, since the burning processes that take place in H and He shells determine their respective thicknesses and their detailed chemical compositions, which are crucial ingredients for determining the evolutionary cooling times. A promising approach to circumvent the problem, and also to directly test the initial-final mass relationship, is to study white dwarfs for which external constraints are available. This is the case of white dwarfs in open and globular clusters (Ferrario et al.~2005, Dobbie et al.~2006) or in non-interacting binaries, for instance, common proper motion pairs (Wegner 1973, Oswalt et al.~1988). Focusing on the latter, it is sound to assume that the members of a common proper motion pair were born simultaneously and with the same chemical composition. Since the components are well separated (100 to 1000 AU), mass exchange between them is unlikely and it can be considered that they have evolved as isolated stars. Thus, important information of the white dwarf, such as its total age or the metallicity of the progenitor, can be inferred from the study of the companion. In particular, if the companion is an F, G or K type star the metallicity can be derived with high accuracy from detailed spectral analysis. On the other hand, the age can be obtained using different methods. In particular, we will use stellar isochrones when the star is moderately evolved, or the X-ray luminosity if the star is very close to the ZAMS. The purpose of this work is to present our spectroscopic analysis of both members of some common proper motion pairs containing a white dwarf, and the semi-empirical intial-final mass relationship that we have derived from this study. The paper is organized as follows. In \S 2 we present the observations done so far and describe the data reduction. Section 3 is devoted to discuss the classification and the analysis of the observed white dwarfs, whereas in \S 4 we present the analysis of the companions. This is followed by \S 5 where we present our main results and finally in \S 6 we elaborate our conclusions. | We have studied a sample of common proper motion pairs comprised of a white dwarf and a FGK star. We have performed high signal-to-noise low resolution spectroscopy of the white dwarf members, which led us to carry out a full analysis of their spectra and to make a re-classification when necessary. From the fit of their spectra to white dwarf models we have derived their atmospheric parameters. Then, using different cooling sequences --- namely those of \cite{sal00} and \cite{fon01} --- their masses and cooling times were obtained. Simultaneously, we have performed independent high resolution spectroscopic observations of their companions. Using the available photometry we have obtained their effective temperatures. Then, from a detailed analysis of their spectra and using either isochrones or X-ray luminosities, we have derived their metallicities and ages (i.e., the metallicities of the progenitors of the white dwarfs and their total ages). These observations allowed us to obtain the initial and final masses of six white dwarfs in common proper motion pairs, four of them corresponding to initial masses below $2\,\rm M_{\sun}$, a range which has not been previously covered by the open cluster data. Our semi-empirical relation shows significant scatter, compatible with the results obtained by \cite{fer05} and \cite{dob06}, which are mainly based on open cluster data. However, the dispersion of the results is higher than the error bars, which leaves some open questions that should be studied in detail (e.g., rotation or magnetic fields). We have shown that common proper motion pairs containing white dwarfs can be useful to improve the initial-final mass relationship, since they cover a wide range of ages, masses and metallicities, and they are also representative of the disk white dwarf population. We have seen that the accuracy in the total ages depends almost exclusively on the evolutionary state of the low-mass companions. Such relative accuracy becomes poor when the star is close to the ZAMS. However, this limitation may not be critical to many common proper motion pairs. Planned deep surveys like GAIA, LSST or the Alhambra Survey will discover thousands of new white dwarfs, some of them belonging to wide binaries. In the meantime, our most immediate priority is to further extend the sample of wide binaries valid for this study. We are working in the search for more wide binaries of our interest in the NLTT catalog (Gould \& Chanam\'e 2004) and also in the LSPM-north catalog (L\'epine \& Bongiorno 2007). Detailed study of the current and future common proper motion pairs of this type should help to explain the scatter in the semi-empirical initial-final mass relationship and to discern whether this is a single-valued function. If consistency between observations and theoretical calculations is found, this would have a strong impact on stellar astrophysics, since this relationship is used in many different areas, such as chemical evolution of galaxies, the determination of supernova rates or star formation and feedback processes in galaxies. | 7 | 10 | 0710.1542 |
0710 | 0710.4151_arXiv.txt | The discovery of dark energy (DE) as the physical cause for the accelerated expansion of the Universe is the most remarkable experimental finding of modern cosmology. However, it leads to insurmountable theoretical difficulties from the point of view of fundamental physics. Inflation, on the other hand, constitutes another crucial ingredient, which seems necessary to solve other cosmological conundrums and provides the primeval quantum seeds for structure formation. One may wonder if there is any deep relationship between these two paradigms. In this work, we suggest that the existence of the DE in the present Universe could be linked to the quantum field theoretical mechanism that may have triggered primordial inflation in the early Universe. This mechanism, based on quantum conformal symmetry, induces a logarithmic, asymptotically-free, running of the gravitational coupling. If this evolution persists in the present Universe, and if matter is conserved, the general covariance of Einstein's equations demands the existence of dynamical DE in the form of a running cosmological term, $\CC$, whose variation follows a power law of the redshift. | Modern Cosmology incorporates the notion of dark energy (DE) as an experimental fact that accounts for the physical explanation of the observed accelerated expansion\,\cite{SNe,WMAP3Y}. Although the nature of the DE is not known, one persistent possibility is the 90-years-old cosmological constant (CC) term, $\CC$, in Einstein's equations. In recent times, one is tempted to supersede this hypothesis with another, radically different, one: viz. a slowly evolving scalar field $\phi$ (``quintessence'') whose potential, $V(\phi)\gtrsim 0$, could explain the present value of the DE and whose equation of state (EOS) parameter $\omega_{\phi}= p_{\phi}/\rho_{\phi}\simeq -1+\dot\phi^2/V(\phi)$ is only slightly larger than $-1$ (hence insuring a negative pressure mimicking the $\CC$ case)\,\cite{weinberg}. The advantage to think this way is that the DE can then be a dynamical quantity taking different values throughout the history of the Universe. However, this possibility can not explain why the DE is entirely due to such an \textit{ad hoc} scalar field and why the contributions to the vacuum energy from the other fields (e.g. the electroweak Standard Model ones) must not be considered. In short, it does not seem to be such a wonderful idea to invent the field $\phi$ and simply replace $\rL=\CC/8\pi\,G$ (the energy density associated to $\CC$, where $G$ is Newton's constant) with $\rho_{\phi}\simeq V(\phi)$. One has to explain, too, why the various contributions (including the additional one $V(\phi)$!) must conspire to generate the tiny value of the DE density at present -- the ``old CC problem''\,\cite{weinberg}. While we cannot solve this problem at this stage, the dynamical nature of the DE makes allowance for this possibility. Furthermore, since there is no obvious gain in the quintessence idea, we stick to the CC approach, although we extend it to include the possibility of a dynamical (``running'') $\CC$ term\,\cite{JHEPCC1,SSRev}. The obvious question now is: where this dynamics could come from? One possibility is that it could originate from the fundamental mechanism of inflation\,\cite{inflation}, which presumably took place in the very early Universe and could have left some loose end or remnant -- kind of ``fossil'' -- in our late Universe, which we don't know where to fit in now. However, what mechanism of inflation could possibly do that? There is in principle a class of distinct possibilities, in particular see \cite{Sahni98,Mongan01}, but our very source of inspiration here is the quantum theory of the conformal factor, which was extensively developed in\,\cite{Conformal1}. For a recent discussion, see e.g. \cite{Conformal2,Conformal3} and references therein. More specifically, we start from the idea of ``tempered anomaly-induced inflation'', which was first proposed in \cite{shocom,DESY} (see also \cite{PST}). It leads essentially to a modified form of the original Starobinsky model\,\cite{Starobinsky}. In the present paper, we push forward the possibility that the mechanism that successively caused, stabilized, slowed down (``tempered'') and extinguished the fast period of inflation in our remote past could have left an indelible imprint in the current Universe, namely a very mild (logarithmically) running Newton's coupling $G$. We show that, if matter is covariantly conserved, this necessarily implies an effective renormalization group (RG) running of the ``cosmological constant'' energy density, $\rL=\rL(a)$, which takes the form of a cubic law of $a^{-1}=1+z$ during the matter dominated epoch ($a$ being the scale factor and $z$ the cosmological reshift). | In this work we have suggested that the presence of dynamical dark energy (DE) in the current Universe is actually a consistency demand of Einstein equations under the two assumptions of: i) matter conservation, and ii) the existence of a period of primordial inflation in the early Universe, especially when realized as ``tempered anomaly-induced inflation''. Based essentially on the previous works\,\cite{shocom,DESY,PST} and on the general setting of the quantum theory of the conformal factor\,\cite{Conformal1,Conformal2}, we have found that if the inflationary mechanism is caused by quantum effects on the effective action of conformal quantum field theory in curved space-time, then the gravitational coupling $G$ becomes a running quantity of the scale factor, $G(a)=G_0/(1-\f\ln a)$, $\f$ being the coefficient of the $\beta$-function for the conformal Newton's coupling. The effect of this coupling on the inflationary dynamics is to efficiently ``temper'' the regime of stable inflation presumably into the FLRW regime. The rigorous high energy calculation of $\f$ in QFT in curved space-time shows that both fermions and bosons produce non-negative contributions ($\f\geq 0$). As a consequence, $G$ becomes an asymptotically-free coupling of the scale factor. Intriguingly enough, we have suggested the possibility that this running might persist in the present Universe and, if so, it could provide a \textit{raison d'\^etre} for the existence of the (dynamical) DE, which would appear in the form of running cosmological vacuum energy $\rL$. In fact, the logarithmic evolution of $G$ induces a power-law running of $\rL$, which is essentially driven by the soft-decoupling terms $\sim H^2\,M_i^2$ (hence by the heaviest particle masses). The result is a Universe effectively filled with a mildly-dynamical DE, which can be perfectly consistent with the present observations. To summarize, from the point of view of the ``RG-cosmology'' under consideration, the current Universe appears as FLRW-like while still carrying some slight imprints of important physical processes that determined the early stages of the cosmic evolution. Most conspicuously, the smooth dynamics of $G$ and $\rL$ can be thought of as ``living fossils'' left out of the quantum field theoretical mechanism that triggered primordial inflation. Remarkably, this framework fits with previous attempts to describe the renormalization group running of the cosmological term\,\cite{JHEPCC1,RGTypeIa,IRGA03,SSS1,croat,SS12,LXCDM12,Bilic07} and could provide an attractive link between all stages of the cosmic evolution. It is reassuring to find that there is a large class of RG models behaving effectively the same way. Differences between them could probably be resolved at the level of finer tests, such as those based on cosmological perturbations and structure formation. For example, in references\,\cite{FSS1,GOPS} it is shown that the study of cosmological perturbations within models of running cosmological constant puts a limit on the amount of running, which is more or less stringent depending on the peculiarities of the model. Similarly, a particular study of perturbations would be required in the present framework (which includes the variation of both $\Lambda$ and $G$) to assess the implications on the parameter $\f$. This study is beyond the scope of the present work. \vspace{1cm} \textit{Acknowledgements}. I am very grateful to Ilya Shapiro for discussions on different aspects of this work and for the fruitful collaboration maintained on the cosmological constant ptoblem over the years. I thank also Ana Pelinson for interesting discussions in the early stages of this work. The author has been supported in part by MECYT and FEDER under project 2004-04582-C02-01, and also by DURSI Generalitat de Catalunya under 2005SGR00564 and the Brazilian agency FAPEMIG. I am thankful for the warm hospitality at the Dept. of Physics of the Univ. Federal de Juiz de Fora, where part of this work was carried out. | 7 | 10 | 0710.4151 |
0710 | 0710.1318_arXiv.txt | We present preliminary results on the calculation of synthetic spectra obtained with the stellar model atmospheres developed by Cardona, Crivellari, and Simonneau. These new models have been used as input within the {\sc Synthe} series of codes developed by Kurucz. As a first step we have tested if {\sc Synthe} is able to handle these models which go down to $\log{\tau_{\rm Ross}}= -13$. We have successfully calculated a synthetic solar spectrum in the wavelength region 2000--4500~\AA\ at high resolution ($R=522\,000$). Within this initial test we have found that layers at optical depths with $\log{\tau_{\rm Ross}} < -7$ significantly affect the mid-UV properties of a synthetic spectrum computed from a solar model. We anticipate that these new extended models will be a valuable tool for the analysis of UV stellar light arising from the outermost layers of the atmospheres. \end {abstract} | \label{sec:1} A set of spectral energy distributions (SEDs) is a very useful tool to analyze stellar spectra and the integrated spectral properties of stellar systems (via some evolutionary population synthesis code, see e.g. Buzzoni 1995). Observational atlases of stellar SEDs generally lack of homogeneous and complete coverage of the main stellar parameters ($T_{\rm eff}, \log{g}$ and [M/H]), therefore many of the recent population analyses rely on results of stellar atmosphere modelling, a practise that has been eased with the development of faster computers and sophisticated computational codes. In order to calculate a theoretical stellar SED it is necessary to have a model atmosphere which describes the physical quantities at different depths and, ideally, a complete set of opacities that account for the absorption of the radiation passing throughout the atmosphere. During the last decades several groups have developed computational codes capable to calculate model atmospheres and spectra at high resolution, some of whom have allowed the public use of their codes, among others, the codes {\sc Atlas9}, {\sc Atlas12} and {\sc Synthe} built by Kurucz\footnote{http://kurucz.harvard.edu/} (1993a,b) and {\sc Tlusty} and {\sc Synspec} constructed by Hubeny \& Lanz\footnote{http://nova.astro.umd.edu/} (1992). In particular, {\sc Atlas9} and {\sc Synthe} codes have been used by our group to calculate the UVBLUE grid of theoretical SEDs and to investigate its potential for stellar and populations studies in the ultraviolet (UV) wavelength interval (see Rodriguez-Merino et al. 2005 for more details). The UV wavelength range has historically been challenging in many branches of modern astrophysics since the observed stellar spectra are not well reproduced by predictions of theoretical models. One possible reason is the missing opacity problem (see Holweger 1970; Gustafsson et al. 1975), but another probable reason is actually that most of the model atmospheres do not provide the atmosphere structure near the stellar surface, where most of the UV radiation emerges. Therefore, it is crucial to calculate new models which describe the outermost layers of the stellar atmospheres. In this work we briefly describe the structure of a new model atmosphere for the Sun which incorporates layers down to $\log{\tau_{\rm Ross}}=-13$. That is, optical depths more than five orders of magnitude thinner compared to classical models currently in use. This model has been couplet to {\sc Synthe} codes for testing their compatibility and exploring the effects of such layers on the UV flux. | \label{sec:4} The main result of this work is that {\sc Synthe} series of codes is capable of treating the CCS model atmospheres, which reach very low values of $\tau_{\rm Ross}$. The analysis of the effects of extending the atmosphere indicates that at mid-UV wavelengths the effects are significant while negligible in the blue. The following steps are to extend the analysis to models with atmospheric parameters different of the Sun, to complement the opacity (both continuous and of lines) as well as to introduce more chemical species. We are in the process of including convection for intermediate and cool star models. A detailed comparison at high resolution with an observed solar atlas (Kurucz et al. 1984) is also underway. | 7 | 10 | 0710.1318 |
0710 | 0710.0128_arXiv.txt | % Preliminary VLBA polarisation results on 6 ``blazars'' from 6.5-cm to 7-mm are presented here. Observing at several different wavelengths, separated by short and long intervals, enabled reliable information about the magnetic (B) field structure to be obtained and for the effect of Faraday Rotation to be determined and corrected. For all sources the magnitude of the core Rotation Measure (RM) derived from the shorter wavelength data was greater than that derived from the longer wavelength data, consistent with a higher electron density and/or B-field strength closer to the central engine. A transverse RM gradient was detected in the jet of 0954+658, providing evidence for the presence of a helical B-field surrounding the jet. The RM in the core region of 2200+420 (BL Lac) displays sign changes in different wavelength intervals (on different spatial scales); we suggest an explanation for this in terms of modest bends in a helical B-field surrounding the jet. | The Faraday effect causes a rotation of the plane of linear polarisation, described by: $\Delta\chi = RM \lambda^2$, with the rotation measure (RM) determined by the integral of the electron density and the dot product of the magnetic (B) field and the path length along the line of sight (LoS). A positive/negative RM indicates that the LoS B-field is pointing towards/away from the observer. $\Delta\chi$ is the change in the electric vector position angle (EVPA). Previous results indicated the presence of different RM signs in the core regions of 6 blazars in different wavelength intervals (O'Sullivan \& Gabuzda 2006). This has two main possible origins: (1) the LoS B-field changes with distance from the centre of activity, or (2) since the previous observations were not simultaneous, it could be due to an intrinsic change in the overall jet B-field structure between observing epochs. Our new 8 wavelength observations are designed to test these possibilities. | Our results for 2200+420 confirm the presence of an RM sign reversal in the core region. Since the dominant jet B-field is transverse to the jet and remains transverse while the jet bends, we will suppose a helical B-field surrounds the jet. The observed RM sign reversal can be explained by a slight bend of the jet, due, for example, to a collision with material in the parent galaxy or some instabilities inherent in the jet itself. (A longitudinal jet B-field with a change in the angle to the line of sight could also cause a RM sign reversal, but this does not correspond to the observed B-field.) A side-on view of a helical B-field (Fig.1 Top right) will have a RM that will be equally strong on both sides of the jet, hence, a zero net RM will be observed for an unresolved jet. This would occur when the source is viewed at $1/\Gamma$ in the observer's frame. For a tail-on view of a helical B-field (ie. $\theta > 1/\Gamma$) (Fig.1 Middle), the dominant RM will be from the bottom half of the jet and a negative RM will be observed because the dominant LoS B-field will be pointing away from us. (Assuming the jet is not fully resolved in the transverse direction.) Conversely, for a head-on view of a helical B-field (ie. $\theta < 1/\Gamma$) (Fig.1 Bottom), a positive RM will be observed. Therefore, regions with different RM signs in the jets of AGN can be explained within a helical B-field model as places where the jet is observed at angles greater than or less than $~1/\Gamma$, due to bends in the jet. Since VLBI resolution is usually not sufficient to completely resolve the true optically thick core, 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 could be derived from observations at different wavelengths (ie. probing different scales of the inner-jet). In our future work, we will attempt to reconstruct the 3D path of the jet through space using the combined information from the observed distributions of the total intensity, linear polarisation, spectral index and rotation measure. | 7 | 10 | 0710.0128 |
0710 | 0710.5682_arXiv.txt | { {\it Context: } One of the most debated issues about sub-mJy radio sources, which are responsible for the steepening of the 1.4 GHz source counts, is the origin of their radio emission. Particularly interesting, from this point of view, is the possibility of combining radio spectral index information with other observational properties to assess whether the sources are triggered by star formation or nuclear activity. {\it Aims:} The aim of this work is to study the optical and near infrared properties of a complete sample of 131 radio sources with $S>0.4$ mJy, observed at both 1.4 and 5 GHz as part of the ATESP radio survey. The availability of multi--wavelength radio and optical information is exploited to infer the physical properties of the faint radio population. {\it Methods:} We use deep multi--colour (UBVRIJK) images, mostly taken in the framework of the ESO \emph{Deep Public Survey}, to optically identify and derive photometric redshifts for the ATESP radio sources. Deep optical coverage and extensive colour information are available for 3/4 of the region covered by the radio sample. Typical depths of the images are $U\sim 25$, $B\sim 26$, $V\sim 25.4$, $R\sim 25.5$, $I\sim 24.3$, $19.5\leq K_{s}\leq 20.2$, $J\leq 22.2$. We also add shallower optical imaging and spectroscopy obtained previously in order to perform a global analysis of the radio sample. {\it Results:} Optical/near infrared counterparts are found for $\sim 78\%$ (66/85) of the radio sources in the region covered by the deep multi--colour imaging, and for 56 of these reliable estimates of the redshift and type are derived. We find that many of the sources with flat radio spectra are characterised by high radio--to--optical ratios ($R>1000$), typical of classical powerful radio galaxies and quasars. Flat--spectrum sources with low $R$ values are preferentially identified with early type galaxies, where the radio emission is most probably triggered by low--luminosity active galactic nuclei. Considering both early type galaxies and quasars as sources with an active nucleus, such sources largely dominate our sample (78\%). Flat--spectrum sources associated with early type galaxies are quite compact ($d<10-30$ kpc), suggesting core-dominated radio emission. | \label{sec:introduction} The faint (sub-mJy) radio population consists of a mixture of different classes of objects. Since the early seventies it has been known that the strongest sources are almost exclusively associated with either active galactic nuclei (AGNs) or giant ellipticals, the latter of which are also known as radio galaxies (99\% above 60 mJy, \citealt{Windhorst90}). More recent work on mJy and sub-mJy sources has revealed that faint sources are also found to be associated with normal elliptical, spiral and star-forming galaxies, with the early type galaxies being the dominant component \citep{Gruppioni1999, Georgakakis1999, Magliocchetti2000, Prandoni2001b, Afonso2006}, while at $\mu$Jy levels star-forming galaxies prevail (see e.g. \citealt{Richards1999}). In spite of the progress made in our understanding of the faint radio population, many questions remain open. For example, the relative fractions of the different types of objects are still quite uncertain, and our knowledge of their dependence on limiting flux density is still incomplete. The reason is, of course, that very little is known about the faint ends of the various luminosity functions, and even less is known about the cosmological evolution of different kinds of objects. This uncertainty is due to the incompleteness of optical identification and spectroscopy, since faint radio sources usually have very faint optical counterparts. Clearly {\it very}\, deep ($I\apprge 25$) optical imaging and spectroscopy, for reasonably large deep radio samples, are critical if one wants to investigate these radio source populations. Since the radio emission comes from different types of objects an important question is what are the physical processes that trigger this emission. It is natural to assume that in the case of star-forming galaxies the emission traces the history of galaxy formation and subsequent evolution by merging and interaction, while the emission in AGNs will reflect black hole accretion history. To make matters more complicated, both processes may be present at the same time. Although research in this field proceeds slowly due to very time--consuming spectroscopy much progress has been made in recent years thanks to strong improvement in the photometric redshift technique. Several multi--colour/multi--object spectroscopy surveys overlapping deep radio fields have recently been undertaken, including the Phoenix Deep Survey \citep{Hopkins1998,Georgakakis1999,Afonso2006} and the Australia Telescope ESO Slice Project (ATESP) survey \citep{Prandoni2000a,Prandoni2000b,Prandoni2001b}. In other cases, deep multi--colour/multi--wavelength surveys have been complemented by deep radio observations (see e.g. the VLA--VIRMOS, \citealt{Bondi2003}; and the COSMOS, \citealt{Schinnerer2006}). Multi--frequency radio observations are also important in measuring the radio spectral index, which may help to constrain the origin of the radio emission in the faint radio sources. This approach is especially meaningful when high resolution radio images are available and radio source structures can be inferred. However, multi--frequency radio information is available for very few, and small, sub-mJy radio samples. The largest sample with multi--frequency radio coverage available so far is a complete sample of 131 radio sources with $S>0.4$ mJy, extracted from a square degree region observed at both 1.4 and 5 GHz as part of the ATESP radio survey \citep{Prandoni2000a,Prandoni2000b,Prandoni2006}. The $1.4-5$~GHz radio spectral index analysis of the ATESP radio sources was presented in the first paper of this series (\citealt{Prandoni2006}, hereafter Paper I). We found a flattening of the radio spectra with decreasing radio flux density. At mJy levels most sources have steep spectra ($\alpha \sim -0.7$, assuming $S\sim \nu^{\alpha}$), typical of synchrotron radiation, while at sub-mJy flux densities a composite population is present, with up to $\sim 60\%$ of the sources showing flat ($\alpha > -0.5$) spectra and a significant fraction ($\sim 30\%$) of inverted-spectrum ($\alpha>0$) sources. This flattening at sub-mJy fluxes confirms previous results based on smaller samples (\citealt{Donnelly1987,Gruppioni1997,Ciliegi2003}). Flat spectra in radio sources usually indicate the presence of a self-absorbed nuclear core, but they can also be produced on larger scales by thermal emission from stars. It is possible to combine the spectral index information with other observational properties and infer the nature of the faint radio population. This is especially important with respect to the class of flat/inverted--spectrum sources as it permits us to study the physical processes that trigger the radio emission in those sources. This kind of analysis needs information about the redshifts and types of the galaxies hosting the radio sources. A detailed radio/optical study of the sample above is possible, thanks to the extensive optical/infrared coverage mostly obtained in the ESO \emph{Deep Public Survey} (DPS, \citealt{Mignano2007,Olsen2006}). We give a brief discussion of all the data collected so far in Sect.~\ref{sec:datacoverage}, followed by a more detailed analysis of the DPS optical data in Sect.~\ref{sec:dpsanalysis}, where we derive the UBVRI colour catalogue and photometric redshifts for the DPS galaxies in the region covered by the ATESP survey, assessing the reliability of the photometric redshifts themselves. In Sects.~\ref{sec:optid} and \ref{sec:radiozphot}, respectively, we use the DPS UBVRIJK optical data to identify the ATESP radio sources and to derive photometric redshifts. A radio/optical analysis of the optically identified radio sources is presented in Sect.~\ref{sec:comp}, while in Sect.~\ref{sec:nature} we discuss the nature of the mJy and sub--mJy population on the basis of all the radio and optical data available to the ATESP sample. The main results are briefly summarised in Sect.~\ref{sec:summary}.\\ Throughout this paper we use the $\Lambda$CDM model, with $H_0=70$, $\Omega_m=0.3$ and $\Omega_{\Lambda}=0.7$. | \label{sec:summary} In this paper we have discussed the nature of the faint, sub-mJy, radio population, using a sample of 131 radio sources that were observed at 1.4 and 5 GHz with the ATCA (the ATESP--DEEP1 sample). A smaller sample of 85 radio sources is covered by deep multi--colour images. These were optically identified down to very faint magnitudes, which was possible thanks to the availability of very deep multi--colour optical material (in U, B, V, R, I, and sometimes J and K bands). The high percentage of identifications ($\sim 78\%$) makes this a sample that is well suited for follow up studies concerning the composition of the sub-mJy population and, in general, the cosmological evolution of the various classes of objects associated with faint radio sources. We summarise our main results here. \begin{itemize} \item For 85\% of the identification sample we succeeded in deriving reliable photometric redshifts, based on the available accurate colours (UBVRIJK). \item Based on spectral types determined either directly from spectroscopy or from the photometry (or both), we find that at the sub-mJy level the large majority of sources are associated with objects that have early type (64\%) and AGN (14\%) spectra; these are of course what we would normally call radio galaxies and quasars. \item Although earlier work (based on shallower optical imaging and spectroscopy) revealed the presence of a conspicuous component of late type and star-burst objects, such objects appear to be important only at brighter magnitudes ($I<19$), and are rare at fainter magnitudes ($19<I<23.5$). \item From an overall comparison of the radio spectral index with other radio and optical properties of the entire ATESP--DEEP1 sample, we find that most sources with flat radio spectra have high radio-to-optical ratios, as expected for classical radio galaxies and quasars. Flat-spectrum sources with low radio-to-optical ratios are preferentially associated with ETS, in which the radio emission is most plausibly triggered by nuclear activity as well, while star-forming galaxies are associated to steep-spectrum radio sources. \item ETS with flat or inverted spectra are mostly compact, with linear size $<10-30$ kpc, suggesting core-dominated radio emission. Their low radio luminosities (in the range $10^{22}-10^{24}$ W/Hz at 1.4 GHz) and the absence of emission lines in their spectra (when available) suggest that they are FRI sources, although these would normally have steeper spectra and be more extended. They may therefore represent specific phases in the life of a radio source, or may be similar to the low power compact radio sources discussed by \cite{Giroletti2005}. \end{itemize} | 7 | 10 | 0710.5682 |
0710 | 0710.2541_arXiv.txt | Preliminary results are presented from a simple, single-antenna experiment designed to measure the all-sky radio spectrum between 100 and 200~MHz. The system used an internal comparison-switching scheme to reduce non-smooth instrumental contaminants in the measured spectrum to 75~mK. From the observations, we place an initial upper limit of $450$~mK on the relative brightness temperature of the redshifted 21~cm contribution to the spectrum due to neutral hydrogen in the intergalactic medium (IGM) during the epoch of reionization, assuming a rapid transition to a fully ionized IGM at a redshift of 8. With refinement, this technique should be able to distinguish between slow and fast reionization scenarios. To constrain the duration of reionization to $\Delta z>2$, the systematic residuals in the measured spectrum must be reduced to 3~mK. | The transition period at the end of the cosmic ``Dark Ages'' is known as the epoch of reionization (EOR). During this epoch, radiation from the very first luminous sources---early stars, galaxies, and quasars---succeeded in ionizing the neutral hydrogen gas that had filled the intergalactic medium (IGM) since the recombination event following the Big Bang. Reionization marks a significant shift in the evolution of the Universe. For the first time, gravitationally-collapsed objects exerted substantial feedback on their environments through electromagnetic radiation, initiating processes that have dominated the evolution of the visible baryonic Universe ever since. The epoch of reionization, therefore, can be considered a dividing line when the relatively simple evolution of the early Universe gave way to more complicated and more interconnected processes. Although the Dark Ages are known to end when the first luminous sources ionized the neutral hydrogen in the IGM, precisely when this transition occurred remains uncertain. The best existing constraints on the timing of the reionization epoch come from two sources: the cosmic microwave background (CMB) anisotropy and absorption features in the spectra of high-redshift quasars. The amplitude of the observed temperature anisotropy in the CMB is affected by Thomson scattering due to electrons along the line of sight between the surface of last scattering and the detector, and thus, it is sensitive to the ionization history of the IGM through the electron column density. In addition, if there is sufficient optical depth to CMB photons due to free electrons in the IGM after reionization, some of the angular anisotropy in the unpolarized intensity can be converted to polarized anisotropy. This produces a peak in the polarization power spectrum at the angular scale size equivalent to the horizon at reionization with an amplitude proportional to the optical depth \citep{1997ApJ...488....1Z}. Measurements by the WMAP satellite of these effects indicate that the redshift of reionization is $z_r\approx11\pm4$ \citep{2007ApJS..170..377S}, assuming an instantaneous transition. Lyman-$\alpha$ absorption by neutral hydrogen is visible in the spectra of many high-redshift quasars and, thus, offers the second currently feasible probe of the ionization history of the IGM. Continuum emission from quasars is redshifted as it travels through the expanding Universe to the observer. Neutral hydrogen along the line of sight creates absorption features in the continuum at wavelengths corresponding to the local rest-frame wavelength of the Lyman-$\alpha$ line. Whereas CMB measurements place an integrated constraint on reionization, quasar absorption line studies are capable of probing the ionization history in detail along the sight-lines. There is a significant limitation to this approach, however. The Lyman-$\alpha$ absorption saturates at very low fractions of neutral hydrogen (of order $x_{HI} \approx 10^{-4}$). Nevertheless, results from these studies have been quite successful and show that, while the IGM is highly ionized below $z\lesssim6$ (with typical $x_{HI}\lesssim10^{-5}$), a significant amount of neutral hydrogen is present above, although precisely how much remains unclear \citep{2001ApJ...560L...5D, 2001AJ....122.2850B, 2002AJ....123.1247F, 2003AJ....125.1649F, 2004Natur.427..815W, 2006AJ....132..117F}. The existing CMB and quasar absorption measurements are somewhat contradictory. Prior to these studies, the reionization epoch was assumed generally to be quite brief, with the transition from an IGM filled with fully neutral hydrogen to an IGM filled with highly ionized hydrogen occurring very rapidly. These results, however, open the possibility that the ionization history of the IGM may be more complicated than previously believed \citep{2003ApJ...595....1H, 2003ApJ...591...12C, 2003MNRAS.344..607S, 2004ApJ...604..484M}. Direct observations of the 21~cm (1420~MHz) hyperfine transition line of neutral hydrogen in the IGM during the reionization epoch would resolve the existing uncertainties and reveal the evolving properties of the IGM. The redshifted 21~cm signal should appear as a faint, diffuse background in radio frequencies below $\nu<200$~MHz for redshifts above $z>6$ (according to $\nu=1420/[1+z]$~MHz). For diffuse gas in the high-redshift ($z\approx10$) IGM, the expected unpolarized differential brightness temperature of the redshifted 21~cm line relative to the pervasive CMB is readily calculable from basic principles and is given by \citep[their \S~2]{2004ApJ...608..622Z} \begin{equation} \begin{array}{rl} \label{eqn_intro_temp} \delta T_{21}(\vec{\theta}, z) \approx~& 23~(1+\delta)~x_{HI} \left ( 1 - \frac{T_\gamma}{T_S} \right ) \\ & \times \left ( \frac{\Omega_b~h^2}{0.02} \right ) \left [ \left ( \frac{0.15}{\Omega_m~h^2} \right ) \left ( \frac{1+z}{10} \right ) \right ]^{1/2} \mbox{mK}, \end{array} \end{equation} where $\delta(\vec{\theta},z)$ is the local matter over-density, $x_{HI}(\vec{\theta},z)$ is the neutral fraction of hydrogen in the IGM, $T_\gamma(z) = 2.73~(1+z)$~K is the temperature of CMB at the redshift of interest, $T_S(\vec{\theta},z)$ is the spin temperature that describes the relative population of the ground and excited states of the hyperfine transition, and $\Omega_b$ is the baryon density relative to the critical density, $\Omega_m$ is the total matter density, and $h$ specifies the Hubble constant according to $H_0=100~ h$~km~s$^{-1}$~Mpc$^{-1}$. From Equation~\ref{eqn_intro_temp}, we see that perturbations in the local density, spin temperature, and neutral fraction of hydrogen in the IGM would all be revealed as fluctuations in the brightness temperature of the observed redshifted 21~cm line. The differential brightness temperature of the redshifted 21~cm line is very sensitive to the \hi spin temperature. When the spin temperature is greater than the CMB temperature, the line is visible in emission. For $T_S \gg T_\gamma$, the magnitude of the emission saturates to a maximum (redshift-dependent) brightness temperature that is about 25 to 35~mK for a mean-density, fully neutral IGM between redshifts 6 and 15, assuming a $\Lambda$CDM cosmology with $\Omega_m=0.3$, $\Omega_\Lambda=0.7$, $\Omega_b=0.04$, and $h=0.7$. At the other extreme, when the spin temperature is very small and $T_S \ll T_\gamma$, the line is visible in absorption against the CMB with a potentially very large (and negative) relative brightness temperature. A number of factors are involved in predicting the typical differential brightness temperature of the redshifted 21~cm line as a function of redshift. In particular, the spin temperature must be treated in detail, including collisional coupling between the spin and kinetic temperatures of the gas, absorption of CMB photons, and heating by ultra-violet radiation from the first luminous sources. We direct the reader to \citet{2006PhR...433..181F} for a good introduction to the topic. The results of several efforts to predict the evolution of the differential brightness temperature of the redshifted 21~cm line have yielded predictions that are generally consistent in overall behavior, but vary highly in specific details \citep{1997ApJ...475..429M, 1999A&A...345..380S, 2004ApJ...608..611G, 2006MNRAS.371..867F}. These models tend to agree that, for a finite period at sufficiently high redshifts ($z\gtrsim20$), the \hi hyperfine line should be seen in absorption against the CMB, with relative brightness temperatures of up to $|\delta T_b|\lesssim100$~mK. This is because the IGM initially cools more rapidly than the CMB following recombination \citep{1994ApJ...427...25S, 1997ApJ...475..429M}. During this period, fluctuations in the differential brightness temperature of the redshifted 21~cm background should track the underlying baryonic matter density perturbations \citep{1972A&A....20..189S, 1979MNRAS.188..791H, 1990MNRAS.247..510S, 2002ApJ...572L.123I, 2003MNRAS.341...81I, 2004PhRvL..92u1301L, 2005ApJ...626....1B}. Eventually, however, the models indicate that the radiation from the first generations of luminous sources will elevate the spin temperature of neutral hydrogen in the IGM above the CMB temperature and the redshifted 21~cm line should be detected in emission with relative brightness temperatures up to the expected maximum values (of order $25$~mK). Finally, during the reionization epoch, the neutral hydrogen becomes ionized, leaving little or no gas to produce the \hi emission, and the apparent differential brightness temperature of the redshifted 21~cm line falls to zero as reionization progresses. As the gas is ionized, a unique pattern should be imprinted in the redshifted 21~cm signal that reflects the processes responsible for the ionizing photons and that evolves with redshift as reionization progresses \citep{1997ApJ...475..429M, 2000ApJ...528..597T, 2003ApJ...596....1C, 2004ApJ...608..622Z, 2004ApJ...613...16F}. The details of the specific timing, duration, and magnitude of these features remains highly variable between theoretical models due largely to uncertainties about the properties of the first luminous sources. Measuring the brightness temperature of the redshifted 21~cm background could yield information about both the global and the local properties of the IGM. Determining the average brightness temperature over a large solid angle as a function of redshift would eliminate any dependence on local density and temperature perturbations and constrain the evolution of the product $\overline{x_{HI}(1-T_\gamma/T_S)}$, where we use the bar to denote a spatial average. During the reionization epoch, it is, in general, believed to be a good approximation to assume that $T_S\gg T_\gamma$ and, therefore, that the brightness temperature is proportional directly to $\bar{x}_{HI}$. Global constraints on the brightness temperature of the redshifted 21~cm line during the EOR, therefore, would directly constrain the neutral fraction of hydrogen in the IGM. Such constraints would provide a basic foundation for understanding the astrophysics of reionization by setting bounds on the duration of the epoch, as well as identifying unique features in the ionization history (for example if reionization occurred in two phases or all at once). They would also yield improvements in estimates of the optical depth to CMB photons and, thus, would help to break existing degeneracies in CMB measurements between the optical depth and properties of the primordial matter density power spectrum \citep{2006PhRvD..74l3507T}. \begin{figure} \centering \includegraphics[width=20pc]{f1_color.eps} \caption[Photograph of EDGES deployed at Mileura Station]{ \label{f_edges_photos} EDGES deployed at Mileura Station in Western Australia. The left panel shows the full antenna and ground screen in the foreground and the analog-to-digital conversion and data acquisition module in the background. The right panel is a close-up view of the amplifier and switching module connected directly to the antenna (through the balun).} \end{figure} For these reasons, several efforts are underway to make precise measurements of the radio spectrum below $\nu<200$~MHz ($z>6$). In this paper, we report on the initial results of the Experiment to Detect the Global EOR Signature (EDGES). In \S~\ref{s_edges_method}, we describe the specific approach used for EDGES to address the issue of separating the redshifted 21 signal from the foreground emission. We then give an overview of the EDGES system in \S~\ref{s_edges_system}, followed by the results of the first observing campaign with the system in \S~\ref{s_edges_results}, along with a discussion of the implications for future single-antenna measurements. | In principle, useful measurements of the redshifted 21~cm background can be carried out with a small radio telescope. These measurements would be fundamental to understanding the evolution of the IGM and the EOR. In particular, the global evolution of the mean spin temperature and mean ionization fraction of neutral hydrogen in the high redshift IGM could be constrained by very compact instruments employing individual radio antennas. We have reported preliminary results to probe the reionization epoch based on this approach from the first observing campaign with the EDGES system. These observations were limited by systematic effects that were an order of magnitude larger than the anticipated signal and, thus, ruled out only an already unlikely range of parameter space for the differential amplitude of the redshifted 21~cm brightness temperature and for the duration of reionization. Nevertheless, the results of this experiment indicate the viability of the simple global spectrum approach. Building on the experiences of these initial efforts, modifications to the EDGES system are underway to reduce the residual systematic contribution in the measured spectrum and to expand the frequency coverage of the system down to 50~MHz or lower in order to place constraints on the anticipated transition of the hyperfine line from absorption to emission as the IGM warms before the EOR. Constraining the redshift and intensity of this feature would be very valuable for understanding the heating history of the IGM and, since the transition has the potential to produce a step-like feature in the redshifted 21~cm spectrum with a magnitude over 100~mK (up to a factor of 4 larger than the amplitude of the step during the reionization epoch), it may be easier to identify than the transition from reionization---although the sky noise temperature due to the Galactic synchrotron foreground increases significantly at the lower frequencies, as well. Through these and other global spectrum efforts, the first contribution to cosmic reionization science from measurements of the redshifted 21~cm background will hopefully be achieved in the near future. | 7 | 10 | 0710.2541 |
0710 | 0710.1175_arXiv.txt | {The physical mechanism responsible for the short outbursts in a recently recognized class of High Mass X--ray Binaries, the Supergiant Fast X--ray Transients (SFXTs), is still unknown. Two main hypotheses have been proposed to date: the sudden accretion by the compact object of small ejections originating in a clumpy wind from the supergiant donor, or outbursts produced at (or near) the periastron passage in wide and eccentric orbits, in order to explain the low ($\sim$10$^{32}$~erg~s$^{-1}$) quiescent emission. Neither proposed mechanisms seem to explain the whole phenomenology of these sources. } {Here we propose a new explanation for the outburst mechanism, based on the X--ray observations of the unique SFXT known to display periodic outbursts, IGR~J11215-5952. } {We performed three Target of Opportunity observations with \sw, \xmm\ and \inte\ at the time of the fifth outburst, expected on 2007 February 9. \sw\ observations of the February 2007 outburst have been reported elsewhere. Another ToO with \sw\ was performed in July 2007, in order to monitor the supposed ``apastron'' passage. } {\xmm\ observed the source on 2007 February 9, for 23~ks, at the peak of the outburst, while \inte\ started the observation two days later, failing to detect the source, which had already undergone the decaying phase of the fast outburst. The Swift campaign performed in July 2007 reveals a second outburst occurring on 2007 July 24, as bright as that observed about 165~days before. } {The new X--ray observations allow us to propose an alternative hypothesis for the outburst mechanism in SFXTs, linked to the possible presence of a second wind component, in the form of an equatorial disk from the supergiant donor. We discuss the applicability of the model to the short outburst durations of all other Supergiant Fast X--ray Transients, where a clear periodicity in the outbursts has not been found yet. The new outburst from \src\ observed in July suggests that the true orbital period is $\sim$165~days, instead of 329~days, as previously thought. } | The Galactic plane monitoring performed by the \inte\ satellite in the last 5 years has allowed the discovery of a number of new High Mass X--ray Binaries (HMXBs). Several of these new sources are intrinsically highly absorbed and were difficult to discover with previous missions (e.g. IGR~J16318--4848, \citealt{Walter2003}). Others are transient HMXBs (associated with OB supergiant) displaying short outbursts (few hours, typically less than a day; \citealt{Sguera2005}), and form the recently recognized new class of Supergiant Fast X--ray Transients (SFXTs). \object{IGR~J11215--5952} is a hard X--ray transient discovered by \inte\ during a fast outburst in April 2005 \citep{Lubinski2005}. The short duration of this outburst led \citet{Negueruela2005a} to propose that \src\ could be a new member of the class of Supergiant Fast X-ray Transients. The optical counterpart is indeed a B-type supergiant, \object{HD~306414} located at a distance of 6.2~kpc (\citealt{Negueruela2005b}, \citealt{Masetti2006}, \citealt{Steeghs2006}). From the analysis of archival \inte\ observations and the discovery of two previously unnoticed outbursts, a recurrence period in the X--ray activity of $\sim$330~days has been found \citep*[hereafter Paper I]{SidoliPM2006}, likely linked to the orbital period of the binary system. This periodicity was later confirmed by the fourth outburst from \src\ observed with $RXTE/PCA$ on 2006 March 16--17, 329~days after the previous one \citep{Smith2006a}. The \inte\ spectrum was well fitted by a hard power-law with a high energy cut-off around 15~keV (Paper~I). Assuming a distance of 6.2 kpc, the peak fluxes of the outbursts correspond to a luminosity of $\sim 3 \times$10$^{36}$~erg~s$^{-1}$ (5--100~keV; Paper~I). The $RXTE/PCA$ observations showed a possible pulse period of $\sim$195\,s \citep{Smith2006b}, later confirmed during the February 2007 outburst, yielding P=$186.78\pm0.3$\,s \citep{Swank2007}. All these findings confirmed \src\ as a member of the class of SFXTs, and the first one displaying periodic outbursts. Based on the known periodicity, an outburst was expected for 2007 February 9. This allowed us to obtain several Target of Opportunity (ToO) observations with $Swift/XRT$, $XMM-Newton$ and $INTEGRAL$. The $Swift/XRT$ results of the February 2007 outburst have been reported in \citealt{Romano2007} (hereafter Paper~II). Here we report the results of the \xmm\ and \inte\ observations of the February 2007 outburst, and of a $Swift/XRT$ campaing performed in July 2007, in order to monitor the supposed apastron passage. | \label{conclusions} We have proposed here a new explanation for the short outbursts in SFXTs, based on our results of a monitoring campain of \src. An equatorial wind from the supergiant companion is suggested, based on the narrow and steep shape of the X--ray lightcurve observed during the latest outburst. The short outburst is suggestive of a deviation from coplanarity ($\theta$$>$0) of the equatorial plane of the companion with the orbital plane (as, e.g., in PSR~B1259-63, \citealt{Wex1998} or in PSR~J0045--7319, \citealt{Kaspi1994}), and of a some degree of eccentricity (e$>$0). Both orbital eccentricity and no-coplanarity can be explained by a substantial supernova ``kick'' (e.g. \citealt{Eggleton2001}) at birth. This could suggest that SFXTs are likely young systems, probably younger than persistent HMXBs. We are aware that the derived quantities for the proposed supergiant equatorial disk are highly uncertain and still speculative. This uncertainty clearly emerges from the fact that the observed X--ray lightcurve can be reproduced with both orbital periods (329~days and 164.5~days), different eccentricities and different wind properties. A determination of the orbital eccentricity of the system, for example, would be essential in better constraining the expected properties of the wind which match with the level of X--ray emission during the outbursts. These considerations highlight the need for an as complete as possible deep monitoring of the X--ray emission along the orbit, together with optical observations in order to constrain the supergiant wind properties. \begin{figure*} \vbox{ \includegraphics[height=6.5cm,angle=0]{8137fig10a.ps} \includegraphics[height=6.5cm,angle=0]{8137fig10b.ps} \includegraphics[height=6.5cm,angle=0]{8137fig10c.ps} \includegraphics[height=6.5cm,angle=0]{8137fig10d.ps} \includegraphics[height=6.5cm,angle=0]{8137fig10e.ps} \includegraphics[height=6.5cm,angle=0]{8137fig10f.ps}} \caption[]{Results of the proposed model compared with the $Swift/XRT$ lightcurve ({\em upper panels}) the expected luminosity variations ({\em middle panels}) and wind density along the whole neutron star orbit ({\em lower panels}) in two cases: {\em on the left} geometry case ``a'' (with orbital period of 164.5~days and an assumed eccentricity of 0.4) is shown (see Fig.~\ref{fig:geom}), while {\em on the right} the results for case ``b'' are reported (orbital period of 329~days, circular orbit and two similar outbursts per orbit). Both models assume a blue supergiant with a mass of 39~M$_\odot$ and radius of 42~R$_\odot$, a ``polar wind'' component \textbf{(``PW'')} with a terminal velocity of 1800~km~s$^{-1}$. The X--ray lightcurve observed with $Swift/XRT$ is better reproduced assuming a ``polar wind'' mass loss rate of 5$\times$10$^{-6}$~M$_\odot$~yrs$^{-1}$ for case ``a'' and 9$\times$10$^{-7}$ ~M$_\odot$~yrs$^{-1}$ for case ``b''. The second wind component, in form of an equatorial disk (``ED''), has a variable velocity ranging from 750~km~s$^{-1}$ to 1400~km~s$^{-1}$ (for case ``a''), and from 850~km~s$^{-1}$ to 1600~km~s$^{-1}$ for case ``b''. The wind density profiles assumed here are reported in the two lower panels. For a magnetic field of 10$^{12}$~G the centrifugal barrier is open along all the orbit in both cases. Dashed line in the first upper right figure shows the model prediction based on a ten times smaller wind density with respect to that assumed to reproduce the peak luminosity (maintaining fixed all the other parameters). } \label{fig:model} \end{figure*} \onltab{2}{ % \begin{table*} \begin{center} \caption{Observation log.} \label{igr112:tab:alldata2} \begin{tabular}{lllll} \hline \hline \noalign{\smallskip} Sequence & Start time (MJD) & Start time (UT) & End time (UT) & Net Exposure$^{\mathrm{a}}$ \\ & & (yyyy-mm-dd hh:mm:ss) & (yyyy-mm-dd hh:mm:ss) &(s) \\ \noalign{\smallskip} \hline \noalign{\smallskip} 00030881024 & 54256.7449 & 2007-06-05 17:52:35 & 2007-06-05 21:15:58 & 1519 \\ 00030881025 & 54263.4918 & 2007-06-12 11:48:11 & 2007-06-12 16:57:58 & 1522 \\ 00030881026 & 54270.3808 & 2007-06-19 09:08:17 & 2007-06-19 11:05:57 & 1843 \\ 00030881027 & 54277.1485 & 2007-06-26 03:33:52 & 2007-06-26 05:27:57 & 2109 \\ 00030881028 & 54284.8571 & 2007-07-03 20:34:14 & 2007-07-03 23:59:57 & 2347 \\ 00030969001 & 54294.5627 & 2007-07-13 13:30:15 & 2007-07-14 07:10:56 & 2126 \\ 00030881030 & 54298.0425 & 2007-07-17 01:01:12 & 2007-07-17 02:58:57 & 2069 \\ 00030881031 & 54305.0012 & 2007-07-24 00:01:42 & 2007-07-24 13:16:55 & 1551 \\ 00030881032 & 54312.2343 & 2007-07-31 05:37:23 & 2007-07-31 10:26:56 & 1695 \\ \noalign{\smallskip} \hline \end{tabular} \end{center} \begin{list}{}{} \item[$^{\mathrm{a}}$] The exposure time is spread over several snapshots (single continuous pointings at the target) during each observation. \end{list} \end{table*} } % | 7 | 10 | 0710.1175 |
0710 | 0710.2189.txt | {Minimal walking technicolor models can provide a nontrivial solution for cosmological dark matter, if the lightest technibaryon is doubly charged. Technibaryon asymmetry generated in the early Universe is related to baryon asymmetry and it is possible to create excess of techniparticles with charge ($-2$). These excessive techniparticles are all captured by $^4He$, creating \emph{techni-O-helium} $tOHe$ ``atoms'', as soon as $^4He$ is formed in Big Bang Nucleosynthesis. The interaction of techni-O-helium with nuclei opens new paths to the creation of heavy nuclei in Big Bang Nucleosynthesis. Due to the large mass of technibaryons, the $tOHe$ ``atomic'' gas decouples from the baryonic matter and plays the role of dark matter in large scale structure formation, while structures in small scales are suppressed. %Due to Nuclear interactions with matter slow down cosmic techni-O-helium in Earth below the threshold of underground dark matter detectors, thus escaping severe CDMS constraints. On the other hand, these nuclear interactions %of techni-O-helium are not sufficiently strong to exclude this form of Strongly Interactive Massive Particles by constraints from the XQC experiment. Experimental tests of this hypothesis are possible in search for $tOHe$ in balloon-borne experiments (or on the ground) and for its charged techniparticle constituents in cosmic rays and accelerators. The $tOHe$ ``atoms'' can cause cold nuclear transformations in matter and might form anomalous isotopes, offering possible ways to exclude (or prove?) their existence.} | The question of the existence of new quarks and leptons is among the most important in the modern high energy physics. This question has an interesting cosmological aspect. If these quarks and/or charged leptons are stable, they should be present around us and the reason for their evanescent nature should be found. Recently, at least three elementary particle frames for heavy stable charged quarks and leptons were considered: (a) A heavy quark and heavy neutral lepton (neutrino with mass above half the $Z$-boson mass) of a 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}; (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}; (c) AC-leptons, based on the approach of almost-commutative geometry \cite{5,book}, that can form evanescent AC-atoms, playing the role of dark matter \cite{5,FKS,Khlopov:2006uv}. In all these recent models, the predicted stable charged particles escape experimental discovery, because they are hidden in elusive atoms, composing the dark matter of the modern Universe. It offers a new solution for the physical nature of the cosmological dark matter. Here we show that such a 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). This approach differs from the idea of dark matter composed of primordial bound systems of superheavy 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, superheavy 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 show how the problems, related to various dark matter scenarios with composite atom-like systems, can find an elegant solution on the base of the minimal walking technicolor model. The approaches (b) and (c) try to escape the problems of free charged dark matter particles \cite{Dimopoulos:1989hk} by hiding opposite-charged particles in atom-like bound systems, which interact weakly with baryonic matter. However, in the case of charge symmetry, when primordial abundances of particles and antiparticles are equal, annihilation in the early Universe suppresses their concentration. If this primordial abundance still permits these particles and antiparticles to be the dominant dark matter, the explosive nature of such dark matter is ruled out by constraints on the products of annihilation in the modern Universe \cite{Q,FKS}. Even in the case of charge asymmetry with primordial particle excess, when there is no annihilation in the modern Universe, binding of positive and negative charge particles is never complete and positively charged heavy species should retain. Recombining with ordinary electrons, these heavy positive species give rise to cosmological abundance of anomalous isotopes, exceeding experimental upper limits. To satisfy these upper limits, the anomalous isotope abundance on Earth should be reduced, and the mechanisms for such a reduction are accompanied by effects of energy release which are strongly constrained, in particular, by the data from large volume detectors. These problems of composite dark matter models \cite{Glashow,5} revealed in \cite{Q,Fargion:2005xz,FKS,I}, can be avoided, if the excess of only $-2$ charge $A^{--}$ %($(\bar U \bar U \bar u)$, $(\bar U \bar U \bar U)$) % or neutral %$(\bar U u)$ particles is generated in the early Universe. Here we show that in walking technicolor models, technilepton and technibaryon excess is related to baryon excess and the excess of $-2$ charged particles can appear naturally for a reasonable choice of model parameters. It distinguishes this case from other composite dark matter models, since in all the previous realizations, starting from \cite{Glashow}, such an excess was put by hand to saturate the observed cold dark matter (CDM) density by composite dark matter. After it is formed in Big Bang Nucleosynthesis, $^4He$ screens the $A^{--}$ charged particles in composite $(^4He^{++}A^{--})$ ``atoms''. These neutral primordial nuclear interacting objects saturate the modern dark matter density and play the role of a nontrivial form of strongly interacting dark matter \cite{Starkman,McGuire:2001qj}. The active influence of this type of dark matter on nuclear transformations seems to be incompatible with the expected dark matter properties. However, it turns out that the considered scenario is not easily ruled out \cite{FKS,I} and challenges the experimental search for techni-O-helium and its charged techniparticle constituents. The structure of the present paper is as follows. Starting with a review of possible dark matter candidates offered by the minimal walking technicolor model, we reveal the possibility for the lightest techniparticle(s) to have electric charge $\pm 2$ (Section II). % In the framework of the considered approach, the %excess of techniparticles relative to their antiparticles can be %generated in early Universe and related to baryon asymmetry. In Section III we show how the minimal technicolor model can provide substantial excess of techniparticles with electric charge $-2$. %can be generated in the form of antitechnibaryons %$(\bar{U}\bar{U})^{-2}$, of technileptons $(\zeta^{-2}$ or of %their mixture. In Section IV we show how all these $-2$ charge particles can be captured by $^4He$, after its formation in the Standard Big Bang Nucleosynthesis (SBBN), % and neutral %$^4He^{+2}(\bar{U}\bar{U})^{-2}$, $^4He^{+2}\zeta^{-2}$ making neutral techni-O-helium ``atoms" that can account for the modern dark matter density. Techni-O-helium catalyzes a path for heavy element formation in SBBN, but we stipulate in Section IV a set of arguments, by which the considered scenario can avoid immediate contradiction with observations. Gas of heavy techni-O-helium ``atoms" decouples from the plasma and radiation only at a temperature about few hundreds eV, so that small scale density fluctuations are suppressed and gravitational instability in this gas develops more close to warm dark matter, rather than to cold dark matter scenario (subsection A of Section V). We further discuss in Section V the possibility to detect charged techniparticle components of cosmic rays (subsection B), effects of techni-O-helium catalyzed processes in Earth (subsection C), and possibilities of direct searches for techni-O-helium (subsection D). The problems, signatures, and possible experimental tests of the techni-O-helium Universe are considered in Section VI. Details of our calculations are presented in the Appendices 1 and 2. % In spite of technilepton $(\zeta^{-2}$ excess, in the %difference with the case of technibaryons, the frozen out %abundance of antiparticles $(\bar \zeta^{+2}$ is not negligible. %Since $(\bar \zeta^{+2}$ represents a form of anomalous helium, it %might cause the problem of anomalous helium overproduction. %Cosmological evolution of technileptons and mechanisms of %suppression for primordial $(\bar \zeta^{+2}$ abundance are %considered in Appendix 2. | Discussion} In this paper we explored the cosmological implications of a walking technicolor model with doubly charged technibaryons and technileptons. The considered model escapes most of the drastic problems of the Sinister Universe \cite{Glashow}, related to the primordial $^4He$ cage for $-1$ charge particles and a consequent overproduction of anomalous hydrogen \cite{Fargion:2005xz}. These charged $^4He$ cages pose a serious problem for composite dark matter models with single charged particles, since their Coulomb barrier prevents successful recombination of positively and negatively charged particles. The doubly charged $A^{--}$ techniparticles considered in this paper, bind with $^4He$ in the techni-O-helium neutral states. %catalyzers of $AC$ binding and AC-leptons may thus escape this trap. To avoid overproduction of anomalous isotopes, an excess of $-2$ charged techniparticles over their antiparticles should be generated in the Universe. In all the previous realizations of composite dark matter scenarios, this excess was put by hand to saturate the observed dark matter density. In our paradigm, this abundance of techibaryons and/or technileptons is connected naturally to the baryon relic density. A challenging problem that we leave for future work 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. Another aspect of the considered approach is more clear. For reasonable values of the techiparticle mass, the amount of primordial $^4He$, bound in this atom like state is significant and should be taken into account in comparison to observations. 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. To conclude, walking technicolor cosmology can naturally resolve most of problems of composite dark matter. Therefore, the model considered in this paper with stable $-2$ charged particles might provide a realistic physical basis for a composite dark matter scenario. | 7 | 10 | 0710.2189 |
0710 | 0710.5560_arXiv.txt | We describe the time- and position-dependent point spread function (PSF) variation of the Wide Field Channel (WFC) of the Advanced Camera for Surveys (ACS) with the principal component analysis (PCA) technique. The time-dependent change is caused by the temporal variation of the $HST$ focus whereas the position-dependent PSF variation in ACS/WFC at a given focus is mainly the result of changes in aberrations and charge diffusion across the detector, which appear as position-dependent changes in elongation of the astigmatic core and blurring of the PSF, respectively. Using $>400$ archival images of star cluster fields, we construct a ACS PSF library covering diverse environments of the $HST$ observations (e.g., focus values). We find that interpolation of a small number ($\sim20$) of principal components or ``eigen-PSFs'' per exposure can robustly reproduce the observed variation of the ellipticity and size of the PSF. Our primary interest in this investigation is the application of this PSF library to precision weak-lensing analyses, where accurate knowledge of the instrument's PSF is crucial. However, the high-fidelity of the model judged from the nice agreement with observed PSFs suggests that the model is potentially also useful in other applications such as crowded field stellar photometry, galaxy profile fitting, AGN studies, etc., which similarly demand a fair knowledge of the PSFs at objects' locations. Our PSF models, applicable to any WFC image rectified with the Lanczos3 kernel, are publicly available. | } Even in the absence of atmospheric turbulence, the finite aperture of Hubble Space Telescope ($HST$) causes light from a point source to spread at the focal plane with the diffraction pattern mainly reflecting the telescope's aperture and optical path difference function. Although the point-spread-function (PSF) of $HST$ is already far smaller than what one can achieve with any of the current ground-based facilities, astronomers' endless efforts to push to the limits of their scientific observations with $HST$ ever increase the demand for the better knowledge of the instrument's PSF. Especially, since the installation of the Advanced Camera for Surveys (ACS) on $HST$, there have been concentrated efforts to carefully monitor and understand the instrument's PSFs, and to utilize the unparalleled resolution and sensitivity of ACS in gravitational weak-lensing (e.g., Jee et al. 2005a; Heymans et al. 2005; Schrabback et al. 2007; Rhodes et al. 2007). Modeling the PSFs of ACS has proven to be non-trivial because of its complicated time- and position-dependent variation. The time-dependent change occurs due to the variation in the $HST$ focus, which relates to the constant shrinking of the secondary mirror truss structure and the thermal breathing of $HST$. The former is the main cause of the long-term focus change, and the secondary mirror position has been occasionally adjusted to compensate for this shrinkage (Hershey 1997). The latter is responsible for the short-term variation of the $HST$ focus and is affected by the instrument's earth heating, sun angle, prior pointing history, roll angle, etc. Even at a fixed focus value of $HST$, the PSFs of ACS also significantly change across the detector from the variation of the CCD thickness and the focal plane errors, which appear as position-dependent changes in charge diffusion and elongation of the astigmatic cores, respectively. The strategies to model these PSF variations can be categorized into two types: an empirical approach based on real stellar field observations and a theoretical prediction based on the understanding of the instrument's optics. The first method treats the optical system of the instrument nearly as a blackbox and mainly draws information from observed stellar images. Although the PSF variation pattern can be most straightforwardly described by the variation of the pixel intensity as a function of position (e.g., Anderson \& King 2006), frequently orthogonal expansion of the observed PSFs (e.g., Lauer 2002; Bernstein \& Jarvis 2002; Refregier 2003) have been utilized to make the description compact and tractable. On the other hand, the second approach mainly relies on the careful analysis of the optical configurations of the instrument and receives feedbacks from observations to fine-tune the existing optics model. The TinyTim software (Krist 2001) is the unique package of this type applicable to most instruments of $HST$. In this paper, we extend our previous efforts of the first kind (Jee et al. 2005a; 2005b; 2006; 2007) to describe the time- and position-dependent PSF variations of ACS/WFC now with the principal component analysis (PCA). In our previous work, we used ``shapelets'' (Bernstein \& Jarvis 2002; Refregier 2003) to perform orthogonal expansion of the PSFs. Shapelets are the polar eigenfunctions of two-dimensional quantum harmonic oscillators, which form a highly localized orthogonal set. Although the decomposition of the stars with shapelets is relatively efficient and has proven to meet the desired accuracy for cluster weak-lensing analyses, the scheme is less than ideal in some cases. One important shortcoming is that it is too localized to capture the extended features of PSFs (Jee et al. 2007; also see \textsection\ref{section_basis_function}). In principle, the orthonormal nature of shapelets should allow us to represent virtually all the features of the target image when the number of basis functions are sufficiently large. However, this is not a viable solution not only because the convergence is slow, but also because the orthonormality breaks down in pixelated images for high orders as the function becomes highly oscillatory within a pixel. The PCA technique provides us with a powerful scheme to obtain the optimal set of basis functions from the data themselves. Unlike ``shapelets'', the basis functions derived from the PCA are by nature non-parametric, discrete, and highly customized for the given dataset. Therefore, it is possible to summarize the multi-variate statistics, with a significantly small number of basis functions (i.e., much smaller than the dimension of the problem). For example, PCA has been applied to the classification of object spectra in large area surveys (Connolly et al. 1995; Bromley et al. 1998; Madgwick et al. 2003). It has been shown that only a small number ($10\sim 20$) of the basis functions or $eigenspectra$ are needed to reconstruct the sample. The application of PCA to the PSF decomposition is used by the Sloan Digital Sky Survey (SDSS) to model the PSF variations (e.g., Lupton et al. 2001; Lauer 2002). Jarvis \& Jain (2004) used the PCA technique to describe the variation in the PSF pattern in the CTIO 75 square-degree survey for cosmic shear analyses. They fit the ``rounding'' kernel component with PCA, not the PSF shape directly. This scheme is motivated by their shear measurement technique (i.e., reconvolution to remove systematic PSF anisotropy). However, in the current study we choose to fit the PSF shapes directly because this is more general in the sense that the rounding kernel components are not uniquely determined for a given PSF. In addition, our PSF library generating the PSF shapes directly has more uses in other studies. We aim to construct a high-quality PSF library for the broadband ACS filters (F435W, F475W, F555W, F606W, F625W, F775W, F814W, and F850LP) from $>400$ archival stellar images, which sample a wide range of the $HST$ environments (e.g., the focus values). Our PSF models describe ACS PSFs in rectified images, specifically, drizzled using the Lanczos3 kernel with an output pixel scale of 0.05\arcsec (see \textsection\ref{section_drizzling_kernel} for the justification of this choice). The results from this work are made publicly available on-line via the ACS team web site\footnote{The full PSF library of ACS will become available at http://acs.pha.jhu.edu/$\sim$mkjee/acs\_psf/.}. We will present our works as follows. The justification and the basic mathematical formalism of PCA are briefed in \textsection\ref{section_PCA}. In \textsection\ref{section_application_acs}, we demonstrate how the technique can be applied to ACS data with some test results. Focus dependency of the ACS PSFs, comparison with TinyTim, and strategies to find matching templates are discussed in \textsection\ref{section_discussion} before we conclude in \textsection\ref{section_conclusion}. | } We showed that the time- and position-dependent ACS/WFC PSF can be robustly described through PCA. The PCA technique allows us to perform orthogonal expansion of the observed PSFs with as few as 20 eigen-PSFs derived from the data themselves. This method is superior to our previous shapelet-based decomposition of the PSFs, capturing more details of the diffraction pattern of the instrument PSF. By interpolating the position-dependent variation of the eigen-PSFs with 5th order polynomials, we are able to recover the observed pattern of the PSF ellipticity and width variation. Although the TinyTim software provides a good approximation of the observed PSFs, we demonstrate that there are some important mismatches between the TinyTim prediction and the real PSFs, which cannot be attributed to CTE degradation of WFC over time. The CTE charge trailing effect should be negligible for these bright high S/N stars, and we do not observe any long-term variation of the pattern (i.e., increasing elongation in parallel read-out direction with time) due to the CTE degradation. Because typical science observations require integration of one or more orbits in broadband filters, the background levels are high ($\sim200$ $e^{-}$ for integration of one orbit). These high background photons are supposed to fill the charge traps and thus mitigate the CTE effects. Therefore, we argue that the CTE-induced elongation is not likely to limit the application of our PSF models extracted from short-exposure observations to long-exposure science images. We have compiled WFC PSFs from $>400$ stellar field observations, which span a wide range of $HST$ focus values. Although the current paper mainly deals with the ACS/WFC PSF issue in the context of weak-lensing analysis, we believe that our PSF model can be used in a wide range of the astronomical data analyses where the knowledge of the position-dependent WFC PSF is needed (e.g., crowded field stellar photometry, robust profile fitting of small objects, weak-lensing analyses, etc.). ACS was developed under NASA contract NAS5-32865, and this research was supported by NASA grant NAG5-7697. | 7 | 10 | 0710.5560 |
0710 | 0710.0345_arXiv.txt | We present results of a population synthesis study aimed at examining the role of spin-kick alignment in producing a correlation between the spin period of the first-born neutron star and the orbital eccentricity of observed double neutron star binaries in the Galactic disk. We find spin-kick alignment to be compatible with the observed correlation, but not to alleviate the requirements for low kick velocities suggested in previous population synthesis studies. Our results furthermore suggest low- and high-eccentricity systems may form through two distinct formation channels distinguished by the presence or absence of a stable mass transfer phase before the formation of the second neutron star. The presence of highly eccentric systems in the observed sample of double neutron stars may furthermore support the notion that neutron stars accrete matter when moving through the envelope of a giant companion. | Recent observations of single and binary pulsars have sparked new questions and challenged accepted ideas on the formation of neutron stars (NSs) and the nature of supernova (SN) kicks. In particular, measurements of pulsar radio emission polarization and pulsar wind nebulae symmetry axis directions have provided increasingly compelling evidence for the alignment of pulsar proper motions and pulsar rotation axes \citep[e.g.][]{2005MNRAS.364.1397J, 2006ApJ...639.1007W, 2006ApJ...644..445N, 2007ApJ...660.1357N, 2007ApJ...664..443R, 2007arXiv0708.4251J}. Assuming the proper motion and pulsar spin axis directions are representative of the natal kick and NS progenitor rotation axis, the alignment suggests that natal kicks are preferentially aligned with the progenitor's rotation axis. Moreover, the increasing sample of double neutron star (DNS) binaries has revealed a possible correlation between the spin period $P_{\rm spin}$ of the first born NS and the binary orbital eccentricity $e$ (see Fig.~\ref{f1}) \citep{2005ASPC..328...43M, 2005ApJ...618L.119F}. Such a relation arises naturally for symmetric SN explosions, but is highly constraining for asymmetric explosions. Our aim in this paper is to examine the role of spin-kick alignment in establishing the observed $P_{\rm spin}$--$e$ correlation. \begin{figure} \includegraphics[height=.24\textheight]{observedpspinns1e.ps} \caption{Orbital eccentricities and spin periods of observed DNSs in the Galactic disk. The dashed line represents the best-fitting log-linear curve $e=-1.7 + 1.2 \log P_{\rm spin}$.} \label{f1} \end{figure} | We investigated the role of spin-kick alignment in establishing a correlation between the spin period of the first-born NS and the orbital eccentricity of DNSs using the StarTrack binary population synthesis code and a simple prescription for the spin-up of a NS due to mass accretion. For DNSs forming through the standard evolutionary channel, spin-kick alignment is compatible with the observed $P_{\rm spin}$--$e$ relation, but does not alleviate the requirement for low kick velocities proposed in previous population synthesis studies based on isotropic kick distributions. This is puzzling considering the stringent lower limits on the kick velocities imparted to the second-born NS in PSR\,B1534+12 and PSR\,B1913+16. Moreover, if low kicks are imparted to the second-born NS in DNSs, the standard formation channel cannot explain the formation of the high-eccentricity systems PSR\,B1913+16 and PSR\,J1811-1736. A possible resolution would be a dichotomous formation channel where low-eccentricty DNSs are formed through the standard formation channel with low kicks imparted to the second-born NS, and high-eccentricity systems are formed through a formation channel where no mass transfer occurs after the common envelope phase of the second NS's progenitor and "normal" kicks typical of isolated radio pulsars are imparted to the second-born NS. The low-eccentricity systems on the $P_{\rm spin}$--$e$ relation can be formed with either polar or isotropic kicks. However, to obtain a $P_{\rm spin}$--$e$ relation at high eccentricities, some spin-kick alignment is required to counteract the increased spread in the post-SN orbital eccentricities introduced by the larger kicks. We will explore this in the continuation of this investigation. \begin{theacknowledgments} This work is partially supported by a Packard Foundation Fellowship, a NASA BEFS grant (NNG06GH87G), and a NSF CAREER grant (AST-0449558) to VK. \end{theacknowledgments} | 7 | 10 | 0710.0345 |
0710 | 0710.3085_arXiv.txt | We present the results of recent observations of phase-dependent variations in brightness designed to characterize the atmospheres of hot Jupiters. In particular, we focus on recent observations of the transiting planet HD~189733b at 8~\micron~using the \emph{Spitzer Space Telescope}, which allow us to determine the efficiency of the day-night circulation on this planet and estimate the longitudinal positions of hot and cold regions in the atmosphere. We discuss the implications of these observations in the context of two other successful detections of more sparsely-sampled phase variations for the non-transiting systems $\upsilon$~And~b and HD~179949b, which imply a potential diversity in the properties of the atmospheres of hot Jupiters. Lastly, we highlight several upcoming \emph{Spitzer} observations that will extend this sample to additional wavelengths and more transiting systems in the near future. | There are currently more than 20 known transiting planetary systems, of which the majority are gas-giant planets orbiting extremely close ($<$0.05~A.U.) to their parent stars\footnote{See http://vo.obspm.fr/exoplanetes/encyclo/encycl.html for the latest count}. These planets, known as ``hot Jupiters'', receive $>$10,000 times more radiation from their stars than Jupiter does from the sun, heating them to temperatures as high as 2000~K \citep{har07}. The picture is further complicated by the fact that most of these planets are expected to be tidally locked, with permanent day and night sides. As a result, the equilibrium compositions and circulation patterns in these atmospheres are expected to differ significantly from those of the gas-giant planets in the solar system. It is not surprising, then, that simple models for the atmospheres of these planets \citep[see, for example, ][]{seag05,bar05,fort06b,bur07b} can sometimes differ significantly in both their assumptions and the conclusions that they reach; this is particularly true of the circulation models discussed in \S\ref{models} Fortunately, it is possible to test the predictions of these models using observations of transiting systems. From the wavelength-dependent depth of the transit, when the planet passes in front of the star, we can search for features in the transmission spectrum of the planet near the day-night terminator \citep{char02,vid03,vid04,knut07a,bar07,tin07,ehren07}. Similarly, by measuring the depth of the secondary eclipse, when the planet moves behind the star, we can characterize the properties of the emission from the day side of the planet \citep{char05,dem05,dem06,dem07,demory07,grill07,rich07,har07,knut07b,knut07c}. Lastly, by measuring the changes in brightness over a significant fraction of a planet's orbit we can directly constrain the longitudinal temperature distribution across the planet's atmosphere \citep{har06,cow07,knut07b}. This last type of observation is particularly informative for hot Jupiters, which may have extreme changes in temperature between the day and night sides of the planet. It is also a difficult observation to make in practice as it requires a very high precision over time scales of several days; this is why all of the successful detections to date have been made from space. Although the shape of the ingress and egress during secondary eclipse also contain information about the brightness distribution on the day side of the planet \citep{will06,rau07}, these variations are smaller and take place on shorter time scales, and there have not been any clear detections of this effect. In this paper we focus on observations of phase variations, discussing the predictions of the current generation of atmospheric circulation models, the results of the three successful measurements of phase variations to date, and upcoming plans for more such observations in the near future. | The above discussion highlights the significant advances that have been made in the year since \citet{har06} reported the first detection of the phase variation of an extrasolar planet. However, the current sample is limited, and there is clearly much to be gained from observations of additional transiting systems, for which the planetary radius, orbital inclination, and dayside flux are all well-constrained. There are a number of exciting programs scheduled for the upcoming \emph{Spitzer} observing cycle which will address this issue. First, we have a program which will observe HD 189733b continuously at 24~\micron~over the same half orbit as \citet{knut07b}, which should probe a different region of the atmosphere than the previous 8~\micron~observations. As a part of this same program we will also observe HD~209458b over half an orbit at both 8 and 24~\micron, which will allow us to directly compare the wavelength-dependent properties of the atmospheric circulation for these two planets. To date all of the observations we have described have examined planets with effectively circular orbits; however, there is another promising set of upcoming observations by G. Laughlin and G. Bakos that will observe two highly eccentric planets (HD 80606b and HAT-P-2b) continuously over $\sim30$ hours centered on periastron passage. Although only one of these planets, HAT-P-2b, is transiting, both observations should detect the rapid increase in brightness as the planetary atmosphere is flash-heated during this passage, providing a direct measurement of the radiative time constant. Of the two methods described above, it is the more time-intensive, continuous observations of transiting systems that have provided the most detailed information about circulation in the atmospheres of hot Jupiters to date. However, the detections of phase variations for $\upsilon$~And~b and HD~179949b highlight the advantages of more sparsely-sampled observations in facilitating a much wider-ranging survey of a large number of planetary systems. These two approaches are inherently complementary, and we expect that both will make valuable contributions to our understanding of these unusual planets in the near future. | 7 | 10 | 0710.3085 |
0710 | 0710.4170_arXiv.txt | {We present the analysis and results of recent high-energy gamma-ray observations of the high energy-peaked BL Lac (HBL) object 1ES 1218+304 with the Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE). 1ES 1218+304 is an X-ray bright HBL at a redshift z=0.182. It has been predicted to be a gamma-ray emitter above 100 GeV, detectable by ground-based Cherenkov telescopes. Recently, this source has been detected by MAGIC and VERITAS, confirming these predictions. STACEE's sensitivity to astrophysical sources at energies above 100 GeV allows it to explore high energy sources such as X-ray bright active galaxies and gamma-ray bursts. We present results from STACEE observations of 1ES 1218+304 in the 2006 and 2007 observing seasons. } \begin{document} | Active galaxies of the ``blazar'' class include BL Lac objects and flat-spectrum radio quasars (FSRQs), and are characterized by non-thermal continuum emission that extends from radio to high energy gamma rays. The spectral energy distributions (SEDs) of these sources typically have two broad peaks, one at low energies (radio to X-ray) and the other at higher energies (keV to TeV). In the framework of relativistic jet models, these objects are highly beamed sources, emitting plasma in relativistic motion (e.g. \cite{Urry&Padovani1995}). In blazar SEDs, the low energy peak is explained as synchrotron emission from high energy electrons in the jet, while the high energy peak is probably due to inverse Compton emission. Several competing ``leptonic'' and ``hadronic'' jet model explanations exist for the high energy emission (e.g. see \cite{Boettcher2002} \& \cite{Muckeetal2003} for reviews), and further broadband observation of blazars are needed to distinguish between these models. Since the discovery of blazars as high energy gamma-ray sources by the Energetic Gamma-Ray Experiment Telescope (EGRET) on board the {\sl Compton Gamma Ray Observatory} (CGRO) \cite{Hartman1999} and the first detection of a TeV ($10^{12}$ eV) blazar by the Whipple Observatory (Mrk 421 \cite{Punch1992}), the search has been on for more TeV blazars. The number count of TeV blazars is growing, with the advent of new generation atmospheric Cherenkov telescopes (ACTs) \cite{Hessblazar19xx}. To date, almost all confirmed blazars detected at TeV energies are high-frequency-peaked BL Lac objects (HBLs), as opposed to quasars that constitute the majority of the EGRET detections. HBLs are a sub-class of BL Lac objects, characterized by lower luminosity than FSRQs and a synchrotron peak in the X-ray band \cite{Urry&Padovani1995}. Blazars are categorized into different sub-classes based on the peak frequencies and the relative power in the low and high energy peaks of their SEDs \cite{Fossati1998}. Given the high synchrotron peak frequencies of HBLs, indicating the presence of high energy electrons, these sources have been predicted to be good candidates for TeV emission, based on synchrotron self-Compton (SSC) emission models \cite{Costamante&Ghisellini2002} as well as hadronic models \cite{Mannheim1993}. Several of the ``extreme'' synchrotron BL Lacs \cite{Costamante2001} have been detected at TeV energies, confirming these predictions. 1ES 1218+304 is an X-ray bright (flux at 1 keV $> 2$ $\mu$Jy) HBL, categorized as an ``extreme'' BL Lac, and predicted to be a TeV source. The source was recently detected by both MAGIC \cite{Albert2006} and VERITAS \cite{Fortin2007}, at energies $>100$ GeV. At a redshift of $z=0.182$, 1ES 1218+304 is one of the most distant blazars detected to date. The source was never detected by EGRET, indicating that ACTs are sensitive to a different population of gamma-ray blazars than EGRET. The Solar Tower Atmospheric Cherenkov Effect Experiment (STACEE) is a ground-based experiment that is sensitive to gamma rays above 100 GeV. STACEE observations of AGN are motivated by the need to understand particle acceleration and emission mechanisms in blazars, as well as their interaction with the extragalactic background light (EBL). Despite what is already known, a great deal remains to be discovered regarding the physics of blazars. STACEE's extragalactic observing program has included both HBLs, as well as LBLs \cite{Mukherjee2006, Mukherjee2005}. Recent observations of 1ES 1218+304 with STACEE were motivated by the detection of TeV emission from the source by MAGIC, providing further evidence that X-ray bright HBLs tend to be strong VHE sources. STACEE observations of 1ES 1218+304 were carried out in the 2006 and 2007 observing seasons. In this paper we present a summary of STACEE observations of the HBL 1ES 1218+304. | STACEE observed the high frequency-peaked BL Lac object 1ES 1218+304 in 2006 and 2007. After all cuts and padding 28.3 hr of data yielded an ON-source excess with a significance of $2.3\sigma$ consistent with no detected flux. Simulated effective areas (Figure 3) were used to derive flux upper limits. For the combined 2006 and 2007 data sets the differential flux upper limit at the 99\% confidence level was derived to be $< 5.2 \times 10^{-6}$ m$^{-2}$ s$^{-1}$ TeV$^{-1}$ at 150 GeV, the energy threshold of STACEE. The upper limit was calculated assuming that the differential flux of photons follows a power law with an index of $-3.0$, as measured by MAGIC \cite{Albert2006}. The STACEE upper-limit is shown overlaid on the MAGIC spectrum in Figure 4. These numbers are consistent with the spectrum measured by MAGIC \cite{Albert2006}, and with the recent VERITAS results indicating a weak gamma-ray source at $\sim5$\% of the Crab flux \cite{Fortin2007}. \begin{figure} \begin{center} \includegraphics*[width=0.48\textwidth,height=2.5in]{rmukherjee_0403_fig4.ps} \end{center} \caption{Gamma-ray spectrum of 1ES 1218+304, as measured by MAGIC (figure from \cite{Albert2006}). The STACEE 99\% flux upper limit is overlaid on the plot at 150 GeV, the energy threshold of STACEE, as obtained from detector simulations. The upper limit was calculated assuming the power-law spectral index of $-3.0$ measured by MAGIC. } \label{fig4} \end{figure} \bigskip \vskip 2in {\bf Acknowledgements: } Many thanks go to the staff of the National Solar Tower Test Facility, who have made this work possible. This work was funded in part by the US National Science Foundation, the Natural Sciences and Engineering Research Council of Canada, Fonds Quebecois de la Recherche sur la Nature et les Technologies, the Research Corporation, and the University of California at Los Angeles. R. M. acknowledges support from NSF grant 0601112. | 7 | 10 | 0710.4170 |
0710 | 0710.1080_arXiv.txt | HCN and CO line diagnostics provide new insight into the OH megamaser (OHM) phenomenon, suggesting a dense gas trigger for OHMs. We identify three physical properties that differentiate OHM hosts from other starburst galaxies: (1) OHMs have the highest mean molecular gas densities among starburst galaxies; nearly all OHM hosts have $\bar{n}({\rm H}_2)=10^3$--$10^4$~cm$^{-3}$ (OH line-emitting clouds likely have $n({\rm H}_2)>10^4$~cm$^{-3}$). (2) OHM hosts are a distinct population in the nonlinear part of the IR-CO relation. (3) OHM hosts have exceptionally high dense molecular gas fractions, $L_{\rm HCN}/L_{\rm CO}>0.07$, and comprise roughly half of this unusual population. OH absorbers and kilomasers generally follow the linear IR-CO relation and are uniformly distributed in dense gas fraction and $L_{\rm HCN}$, demonstrating that OHMs are independent of OH abundance. The fraction of non-OHMs with high mean densities and high dense gas fractions constrains beaming to be a minor effect: OHM emission solid angle must exceed $2\pi$ steradians. Contrary to conventional wisdom, IR luminosity does not dictate OHM formation; both star formation and OHM activity are consequences of tidal density enhancements accompanying galaxy interactions. The OHM fraction in starbursts is likely due to the fraction of mergers experiencing a temporal spike in tidally driven density enhancement. OHMs are thus signposts marking the most intense, compact, and unusual modes of star formation in the local universe. Future high redshift OHM surveys can now be interpreted in a star formation and galaxy evolution context, indicating both the merging rate of galaxies and the burst contribution to star formation. | OH megamasers (OHMs) are rare luminous 18~cm masers associated with major galaxy merger-induced starbursts. The hosts of OHMs are (ultra)luminous IR galaxies ([U]LIRGs), and the OHM fraction in (U)LIRGs peaks at about 1/3 % in the highest luminosity mergers \citep{darling02a}. It is not known whether all major mergers experience an OHM stage or what detailed physical conditions produce OHMs, but it is clear that OHMs are a radically different phenomenon from the aggregate OH maser emission associated with ``normal'' (Galactic) modes of star formation in galaxies. \citet{lo05} posed a key question: why do $80\%$ of LIRGs show no OHM activity? To reframe the question: given two merging systems with similar global IR and radio continuum properties in the same morphological stage of merging, why does one show OHM emission while the other does not? What is the difference between the two systems? Perhaps there is no difference and the fraction of OHMs among mergers simply reflects beaming or OH abundance. Or perhaps OHM activity depends on small-scale conditions that are decoupled from global properties of mergers. The provenance of OHM emission vis-\`{a}-vis the host galaxy has been extensively investigated in the radio through X-ray bands by comparing samples of OHM galaxies to similarly selected non-masing control samples. For example, \citet{darling02a} and \citet{baan06} studied radio and IR properties vis-\`{a}-vis the AGN versus starburst contributions to OHM activity, \citet{baan1992} and \citet{darling02a} investigated the OHM fraction in (U)LIRGs versus star formation rate and IR color, \citet{baan1998} and \citet{darling06} used optical spectral classification to distinguish populations and to quantify AGN fraction in OHM hosts, and \citet{vignali05} conducted an X-ray study of the contribution of AGNs to OHM hosts. While some of these studies pointed to minor differences in statistical samples of OHM hosts versus nonmasing systems, they could not identify on a case-by-case basis which systems would harbor OHMs and which would not based on any observable quantity except the OH line itself. Theoretical modeling of OHM formation has seen a recent renaissance: \citet{parra2005} model the $\sim50$~pc molecular torus in III~Zw~35 and show how OHM emission is a stochastic amplification of unsaturated emission by multiple overlapping clouds, and \citet{lockett07} show how the general excitation of OHMs is fundamentally different from Galactic OH maser emission and predict that a single excitation temperature governs all 18~cm OH lines. While the physics of OHMs is crystallizing, and models predict that beaming is not likely to be the dominant factor in the OHM fraction among (U)LIRGs, it remains unclear on a case-by-case basis what conditions found in starbursts drive or prohibit OHM formation. Here we describe a dense gas trigger for OHM formation, at last identifying physical observable properties that differentiate OHMs from nonmasing mergers. We identify OHMs, OH absorbers, OH kilomasers, and OH non-detections in the Gao \& Solomon (2004a; hereafter GS04a) sample (\S \ref{sec:sample}) and employ CO($1-0$) and HCN($1-0$) molecular gas tracers to show that while OH absorbers appear nearly uniformly distributed in $L_{\rm IR}$ and $L_{\rm HCN}$, OHMs represent the {\it majority} of the nonlinear population in the IR-CO relation (\S \ref{sec:results}). In combination with a Kennicutt-Schmidt-based star formation model of CO line emission by \citet{krumholz07}, we identify a high mean molecular density driving OHM emission, and from the HCN/CO ratio we find that OHM galaxies are exclusively high dense gas fraction starbursts (\S \ref{sec:results}). Now that we can at last observe quantities that are highly predictive of OHM activity, we can employ OHMs at high redshifts as probes of major galaxy mergers and extreme star formation (\S \ref{sec:discussion}). | We have identified three closely related physical properties that differentiate OHMs from other starburst galaxies: OHM hosts have the highest mean molecular gas densities, they are a distinct population in the nonlinear part of the IR-CO relation, and they reside in galaxies with exceptionally high dense molecular gas fractions. We conclude that molecular gas must be concentrated and massive in order to reach the mean density required to form an OHM in a galactic nucleus. IR luminosity is not a condition for OHM formation; both star formation and OHM activity are consequences of the tidal density enhancements accompanying galaxy interactions. The fraction of OHMs in dense starbursts constrains OHM beaming to be a minor effect: OHM solid angle emission must be greater than $2\pi$ steradians. These conclusions are in good agreement with the stochastic cloud-cloud overlap amplification model by \citet{parra2005}. The rather uniform distribution of OH absorbers in IR, HCN, and CO luminosity suggests that OH abundance is not a significant factor in OHM formation. The main caveat to these conclusions is that the sample of OHMs with HCN observations remains small, and should be expanded, particularly to higher redshifts to include ``typical'' OHMs. OHMs are signposts of the most intense, compact, and unusual modes of star formation in the local universe, and surveys for OHMs will now provide detailed information about the detected host galaxies and their mode of star formation. The missing datum required for a complete interpretation of OHM surveys, however, is the OHM lifetime. | 7 | 10 | 0710.1080 |
0710 | 0710.4941_arXiv.txt | The bulk composition of an exoplanet is commonly inferred from its average density. For small planets, however, the average density is not unique within the range of compositions. Variations of a number of important planetary parameters---which are difficult or impossible to constrain from measurements alone---produce planets with the same average densities but widely varying bulk compositions. We find that adding a gas envelope equivalent to 0.1\%-10\% of the mass of a solid planet causes the radius to increase 5-60\% above its gas-free value. A planet with a given mass and radius might have substantial water ice content (a so-called ocean planet) or alternatively a large rocky-iron core and some H and/or He. For example, a wide variety of compositions can explain the observed radius of GJ 436b, although all models require some H/He. We conclude that the identification of water worlds based on the mass-radius relationship alone is impossible unless a significant gas layer can be ruled out by other means. | Out of over 250 exoplanets known to date, over 20 are known to transit their stars. Transiting planets are important because we can derive the precise mass and radius, and can begin to determine other planetary properties, such as the bulk composition. Much attention has been given to ``ocean planets'' or ``water worlds'', planets composed mostly of solid water \citep{kuch2003, lege2004}. If a water world is found close to a star, it would be strong evidence for migration because insufficient volatiles exist near the star for {\it in situ} formation. The proposed identification of water worlds is through transits. From a measured mass and radius a low-density water planet could potentially be identified. We examine the possibility that water worlds cannot be uniquely identified based on the mass and radius of a transiting planet. An alternative interpretation could be a rocky planet with a thick hydrogen-rich atmosphere. Most authors have assumed that solid planets in the 5 to 10 $M_{\oplus}$ range have an insignificant amount of hydrogen \citep{vale2006, vale2007a, fort2007, seag2007, sels2007, soti2007}. Exoplanets have, however, contradicted our basic assumptions before. Notable examples include: the existence of hot Jupiters; the predominance of giant planets in eccentric orbits; and the gas-rock hybrid planet HD~149026b with its $\sim 60 M_{\oplus}$ core and $\sim 30 M_{\oplus}$ H/He envelope \citep{sato2005}. We adopt the idea that a wide range of atmospheric formation and loss mechanisms exist and can lead to a range of atmosphere masses on different exoplanets. We explore the mass-radius relationship for the lowest-mass exoplanets yet detected ($\sim$~5-20 $M_{\oplus}$) in order to identify potential ambiguities that result from the presence of a massive atmosphere. We explore atmospheres ranging from $\sim 10^{-3} M_{\oplus}$ (10 $\times$ Venus' atmospheric mass) to $\sim 1 M_{\oplus}$ (the estimated mass of Uranus' and Neptune's H/He \citep{guil2005, podo1995, hubb1991}), with a focus on the smaller mass range. We also explore potential compositions for the transiting Neptune-size planet GJ~436b \citep{butl2004, gill2007a}. | Our fiducial planet consists of a 30\% Fe core and a 70\% MgSiO$_3$ mantle, roughly analogous to Earth. We used a H/He mixture with helium mass fraction $Y=0.28$ (the He mass fraction of the solar nebula). We chose $T_{eq}=300$~K, based on the observation that a planet around an M-dwarf at an orbital distance of 0.1 AU has a similar equilibrium temperature to Earth (assuming similar albedos). We set $T_{eff}=30$~K, similar to Earth and Uranus\footnote{Earth has 44 $\times 10^{12}$~ W \citep{poll1993} and Uranus has $340 \times 10^{12}$~W of energy flow \citep{pear1990}.}. For the atmospheric parameters, we fixed $\mu_0 = \cos 60^{\circ}$ and $\gamma=0.1$ to represent radiation absorbed deep in the atmosphere\footnote{In comparison $\gamma=10$ would correspond to absorption high in the atmosphere.}. We later investigate variations on $Y$, $T_{eff}$, $T_{eq}$ and $\gamma$. Figure~1 shows a plot of the mass-radius relationship for fiducial planets of masses 5, 10, 15, and 20 $M_{\oplus}$. For each planet mass, we added atmospheres ranging in mass from 0.001-1 $M_{\oplus}$. A robust finding for all models is that a small amount of gas creates a large radius increase. While this result is expected, the radius increase is far more dramatic than anticipated. For example, an H/He atmosphere of $\sim 0.001$ by mass---only ten times greater than Venus' atmospheric mass fraction---is required for a noticeable radius increase. As seen in Figure~1, adding a hydrogen-helium atmosphere with just 0.1\% of the mass of a 10 $M_{\oplus}$ rocky planet results in a 5\% increase in the planetary radius---within a measurement precision that has been obtained for currently known transiting planets. \begin{figure} \plotone{f1} \caption{The increase in radius due to adding H/He to a solid planet. A H/He layer of 0.002-1 $M_{\oplus}$ is added to a solid planet of 5, 10, 15, or 20 $M_{\oplus}$, with fiducial model parameters (30\% Fe and 70\% MgSiO$_3$). The black points are for atmospheres at 0.01 $M_{\oplus}$ and every 0.1 $M_{\oplus}$ afterwards. The mass-radius relationship of solid planets with no gas is plotted for comparison. The water (blue), rock (red), and iron (green) curves are taken from \citet{seag2007} and represent homogeneous solid planets. Intermediate compositions for differentiated planets are, from top down: dashed-blue, 75\% H$_2$O, 22\% MgSiO$_3$, and 3\% Fe; dashed-dotted-blue, 48\% H$_2$O, 48.5\% MgSiO$_3$, and 6.5\% Fe; dotted-blue, 25\% H$_2$O, 52.5\% MgSiO$_3$, and 22.5\% Fe; dashed-red, 67.5\% MgSiO$_3$ and 32.5\% Fe; and dotted-red, 30\% MgSiO$_3$ and 70\% Fe. In general, the addition of a gas layer of up to $\sim 5$\% of the solid planet mass will inflate the radius of a rocky-iron planet through the range of radii corresponding to water planets with different water mass fractions. \label{fig:gas1}} \end{figure} As a second example, adding a gas layer of H/He equal to 1\% of the mass of our fiducial planets increases the radius by $\sim$ 20\% of the original planet radius, or by about 0.35 $R_{\oplus}$ for the planet masses we considered. Our major finding is that that exoplanets with a significant H/He layer cannot be distinguished from water worlds, based on $M_p$ and $R_p$ alone. For our fiducial solid exoplanets, adding up to 5\% H/He by mass (for 10 $M_{\oplus}$ planets) is sufficient to push the planet's radius through the entire range of radii corresponding to solid planets with no gas, including planets with up to 100 percent water composition. While we have not completed an exhaustive study of possible compositions, we find the non-uniqueness of water planets to be valid for any conditions we investigated. This generic finding holds for a wide range of assumptions of assumed temperatures. Taking our fiducial model, we vary $T_{eq}$, $T_{eff}$, and $\gamma$ individually. For a 10 $M_{\oplus}$ solid planet with an additional 0.1 $M_{\oplus}$ H/He atmosphere, increasing $T_{eq}$ from 300 to 500 K increases the radius by about ~1\% (Figure~2). For the same planet varying $T_{eff}$ from 10~K to 50~K results in an 8\% increase in radius. While large, this value is comparable to expected radius uncertanties for these planets \citep{gill2007a, gill2007b, demi2007}. Varying the altitude where radiation is absorbed (specified by $\gamma$) has a much smaller effect on the planet radius. Varying $\gamma$ from 0.1 to 10 causes the radius to decrease by 0.2\%. \begin{figure} \plotone{f2} \caption{The effects on the radius of varying the equilibrium temperature (top) and effective temperature (bottom) for a 10 $M_{\oplus}$ planet with otherwise fiducial parameters. $T_{eq}$ values of 300, 400, and 500~K are plotted to simulate the effect of uncertainty in the orbital parameters and albedo on the expected radius. $T_{eff}$ values of 10, 30, and 50 K are plotted on the same scale as the $T_{eq}$ plot, to show uncertainties in the planet's interior temperature. Uncertainties in the internal energy of a planet lead to large variations in radii for a given mass, showing how temperature is a large uncertainty in the interpretation of a planet's internal composition. \label{fig:temp}} \end{figure} A corollary of our main result is that when a planet has a significant H/He atmosphere there is a wide degeneracy in allowable internal composition. This is not just compositional, but also relates to the trade off of temperature and mass of H/He gas. It could be argued that specifying a planet's composition implies a particular internal thermal profile derived from a consistent cooling history. As addressed in \S2, the many unknowns and free input parameters for rocky planet interiors---such as the possible differences between atmospheric and interior compositions, equation of states, and the effect of tides on the planet's cooling history---prevent a self-consistent solution for the present time. How could a 5--20 $M_{\oplus}$ exoplanet get a substantial H/He layer? Two different scenarios may produce them: direct capture of gas from the protoplanetary disk (possibly modified by the escape of some fraction of the original gas) or outgassing during accretion. A planet may capture and retain up to 1 to 2 $M_{\oplus}$ of H/He if the planetary core did not grow quickly enough to capture more before the gas in the disk evaporated (as is the paradigm for Uranus and Neptune). Alternatively, for short-period exoplanets, a 1 to 2~$M_{\oplus}$ H/He envelope may result after substantial loss of an initially massive gas envelope from irradiated evaporation \citep[e.g.][]{bara2006}. \citet{alib2006} consider atmospheric evaporation during migration, and conclude that the $10 M_{\oplus}$ innermost planet in HD~69830, at 0.08 AU, kept $\sim$2 $M_{\oplus}$ of H/He over the 4 Gyr lifetime of the star. Little attention has been given to the mass and composition of exoplanet atmospheres from outgassing. Venus' atmosphere is $10^{-4} M_{\oplus}$; if Venus had a surface gravity high enough to prevent H escape, its atmosphere would be over $10^{-3} M_{\oplus}$. Even more massive H-rich atmospheres are possible. If a massive iron-silicate planet formed with enough water, the iron may react with the water during differentiation, liberating hydrogen gas \citep{ring1979, waen1994}. L. Elkins-Tanton et al. (in prep.) estimate that the maximum H component is about six percent by mass for a terrestrial-composition planet. For a $10 M_{\oplus}$ planet this would result in a $0.6 M_{\oplus}$ H envelope. For short-period, low-mass planets, theoretical arguments of atmospheric escape may be the best way to identify a water world based on the mass and radius measurements alone \citep{lege2004, sels2007}. Indeed, our assumption of H/He atmospheres for exoplanets relies on the condition that atmospheric mass loss has not evaporated all of the H/He. In the absence of hydrodynamic escape, the exospheric temperature (and not the atmospheric $T_{eff}$) drives the thermal Jeans escape of light gases. Earth and Jupiter both have exobase temperatures of 1000~K \citep{depa2001}, significantly above their $T_{eff}$ of 255~K and 124.4~K respectively \citep{cox2000}. Uranus and Neptune have exobase temperatures around 750~K \citep{depa2001}. We note that because GJ 436 must have at least 1~$M_{\oplus}$ of H/He, its exospheric temperature is not too high. On the other extreme, planets of $5 M_{\oplus}$ would require very low exospheric temperatures ($\sim 300$~K) to retain a massive atmosphere over the course of billions of years. Nevertheless, a young $5 M_{\oplus}$ Earth-mass planet with a captured atmosphere could still have a H/He atmosphere and an old $5 M_{\oplus}$ planet could retain a substantial He fraction, making its compositional identification ambiguous. \begin{figure} \plotone{f3} \caption{Density vs. radius for three different potential compositions of GJ436b. From top to bottom: solid (black) curve, 19.0 $M_{\oplus}$ core (30\% Fe, 70\% MgSiO$_3$) with 3.2 $M_{\oplus}$ H/He (Y=0.28); dotted (red) curve, 20.0 $M_{\oplus}$ core (100\% (Mg,Fe)SiO$_3$) with 2.2 $M_{\oplus}$ H (Y=0); dashed (blue) curve, 20.5 $M_{\oplus}$ core (90\% H$_2$O, 10\% MgSiO$_3$) with 1.7 $M_{\oplus}$ H/He (Y=0.28). All three planets have the same total radius (4.3 $R_{\oplus}$) and total mass (22.2 $M_{\oplus}$). \label{fig:GJinterior}} \end{figure} We now turn to a qualitative study of GJ436b, to show that the interpretation that GJ436b is a water world akin to Uranus and Neptune \citep{gill2007b} is not the only possibility. We consider the GJ~436b values $M_p = 22.2 M_{\oplus}$ and $R_p = 4.3 R_{\oplus}$ from \citet{demi2007}. The internal structure in Figure~3 shows how three planets with very different internal compositions can have the same total mass and radius. We first explore a planet similar to our fiducial model: a 22.2 $M_{\oplus}$ solid planet with Earth-like iron/rock mass ratio (30/70), $T_{eff}=30$~K, and $T_{eq}=600$~K in rough agreement with the orbital parameters (assuming an albedo of 0.1). By adding $\sim$3.2$M_{\oplus}$ of hydrogen-helium to the 19.0 $M_{\oplus}$ solid planet, we are able to reproduce GJ~436b's radius. We note that the mass of gas is 15\% of the solid mass, likely too much to have originated from outgassing, and so capture must be at least partially invoked to explain such a massive atmosphere. The second composition for GJ~436b we considered is for water worlds, one with a 50\% water mantle (by mass) and 50\% silicate core, and another with 90\% water mantle and a 10\% silicate core. These planets also need some H/He to match the known radius, 12\% and 8\% by mass, respectively. The third model approximates planets with atmospheres created from outgassing, considering an extreme scenario where all of the available water has oxidized iron, leaving a 100\% (Mg,Fe)SiO$_3$ solid planet core. To match the observed radius, a 22.2 $M_{\oplus}$ planet requires $\sim 2.2 M_{\oplus}$ of H alone, a case that assumes no initial trapping and subsequent outgassing of He. The pure-hydrogen atmosphere is 10\% of the mass of the solid planet. This is above the theoretical maximum of outgassing based on observed abundances of metallic iron in chondritic meteorites from our solar system (see Elkins-Tanton et al. in prep). Although not an exhaustive study, the range of interior compositions illustrates the variety of possibilities, though all models require some H/He. While our study is preliminary, we make the robust point that H-rich thick atmospheres will confuse the interpretation of planets based on a measured mass and radius. This point is independent of the uncertainties retained by our model including $T_{eq}$, $T_{eff}$, the mass fraction of H/He, and the mixing ratio of H and He. We find that the identification of water worlds based on the mass-radius relationship alone is impossible unless a significant gas layer can be ruled out by other means. Spectroscopy is the most likely means and may become routine with transit transmission and emission spectroscopy, and eventually with spectroscopy by direct imaging. | 7 | 10 | 0710.4941 |
0710 | 0710.4036_arXiv.txt | We present near-infrared (H- and K-band) integral-field observations of the circumnuclear star formation rings in five nearby spiral galaxies. The data, obtained at the {\it Very Large Telescope} with the SINFONI spectrograph, are used to construct maps of various emission lines that reveal the individual star forming regions ("hot spots") delineating the rings. We derive the morphological parameters of the rings, and construct velocity fields of the stars and the emission line gas. We propose a qualitative, but robust, diagnostic for relative hot spot ages based on the intensity ratios of the emission lines \brg , \hel , and \fetwo . Application of this diagnostic to the data presented here provides tentative support for a scenario in which star formation in the rings is triggered predominantly at two well-defined regions close to, and downstream from, the intersection of dust lanes along the bar with the inner Lindblad resonance. | In many spiral galaxies of early- and intermediate Hubble type (Sa-Sc), active star formation is organized in a ring-like structure. These star formation rings offer a unique opportunity to study massive star formation in external galaxies: they produce up to 2/3 of the bolometric luminosity of their host galaxies \citep[e.g. NGC\,7469; ][]{gen95}, and often contain a large fraction of the entire star formation activity of the galaxy. The general picture of why molecular gas assembles in a ring is well understood as a natural consequence of a non-axisymmetric gravitational potential, nearly always due to the presence of a stellar bar or oval distortion \citep[e.g.][]{com85,kna95,hel96,but96}. Because of its dissipative nature, molecular gas accumulates around the radii at which the stellar orbits experience dynamical resonances with the rotating bar potential. Depending on the pattern speed of the bar and the rotation curve of the galaxy, there can be one or multiple such resonances. The high gas densities found in the rings, combined with a variety of excitation mechanisms such as ultra-violet radiation from young stars and mechanical shocks due to energetic outflows from massive stars and/or an active galactic nucleus (AGN) help reveal the physical state of the interstellar matter (ISM), be it molecular, atomic, or ionized gas. Besides being fascinating laboratories in their own right, star formation rings are important also for the secular evolution of disk galaxies \citep{kor04}. This is particularly true for the innermost of the dynamical resonances which is called either the "inner Lindblad resonance" (ILR) or - in cases where a compact massive object leads to an additional dynamical resonance - the ``nuclear Lindblad resonance'' \citep*[NLR; ][]{fuk98}. Either of these resonances can produce gas rings with radii of a few hundred pc or less, depending on the enclosed mass, and the rotation speed of the galaxy disk. On these spatial scales, a number of processes can cause the gas to lose angular momentum and to subsequently flow towards the nucleus. Examples for such processes include torques due to the stellar potential \citep{gar05}, dynamical friction between giant molecular clouds that form within the ring due to self-gravity of the gas \citep*{fuk00}, the formation of spiral density waves \citep{eng00}, or mechanical energy released by star formation in the ring via stellar winds and/or supernova explosions. Understanding the gas dynamics and star formation processes of nuclear\footnote{Throughout this paper, the term ``nuclear ring'' is used to imply that it is the innermost (star forming) gas ring that can be resolved with the resolution of our data, typically a few hundred pc in diameter.} rings in disk galaxies is therefore crucial for developing models of gas accumulation at the very nucleus of a galaxy and the evolution of any compact massive object (CMO) which can exist in the form of a nuclear star cluster \citep[NC,][]{boe02} and/or a supermassive black hole (SMBH). While some theoretical models exist on the gas behavior around a CMO \citep[e.g.][]{fuk00}, observational data to constrain these models are rare, mostly because of the limited spatial resolution of mm-observations. Only recently has it become possible to study the molecular gas flows within a few tens of pc from the nucleus in a small number of nearby galaxies \citep[e.g.][]{sch03,sch06,sch07}. In order to increase the number of well-studied nuclear rings, we have begun a project to study the near-infrared (NIR) properties of five such objects, using the SINFONI integral-field spectrograph on the Very Large Telescope (VLT). This paper discusses the morphologies, star formation rates, and kinematic properties of the rings. In a subsequent paper, we will make use of the spectroscopic information contained in the SINFONI data to investigate in more detail the star formation process, stellar populations, and gas excitation mechanism(s) in individual galaxies, both in the rings and the galaxy nuclei. This paper is structured as follows. In \S\ref{sec:data}, we describe the sample selection, observational details, and data reduction techniques common to all galaxies. The continuum and emission line morphology of the rings as well as the velocity fields of stars and gas for the individual galaxies are presented in \S\ref{sec:diagnostics}. We discuss the results in the context of competing models for the propagation of star formation in the rings in \S\ref{sec:discuss}, and summarize our analysis in \S\ref{sec:summary}. | } Based on high-resolution NIR integral-field observations of five nuclear star formation rings, we have presented their emission line morphologies and velocity structure both in gas and stars. We have introduced a new method to derive relative hot spot ages along the rings using the relative strengths of the \hel , \brg , and \fetwo\ lines. We employ this method to investigate the plausibilty of two competing scenarios for the way star formation is induced in nuclear rings, namely i) the ``popcorn'' model in which hot spots appear stochastically around the ring, and ii) the ``pearls on a string'' scenario in which star formation is triggered predominantly at two overdensity regions on either side of the ring. Only the latter predicts a well-ordered age sequence of hot spots along either half of a ring. The data presented in this study provide tentative support for the ``pearls on a string'' scenario, in that three out of five sample galaxies show some evidence for an age gradient of hot spots along the ring, while the remaining two galaxies have incomplete information and thus are consistent with either model. Of course, it might well be that star formation proceeds differently in some rings than in others. However, given that nuclear rings appear to form via a common mechanism (i.e. the gas response to a bar-shaped potential), it is not unreasonable to expect that they also follow a common path for inducing star formation. The small number of objects described here is clearly insufficient to provide reliable statistics, and similar studies of larger galaxy samples are needed to decide whether a particular mechanism governs the star formation in the majority of nuclear rings. Be that as it may, the proposed method of using NIR line ratios to estimate relative ages of young star clusters has been demonstrated to be a powerful tool for studies along this line. | 7 | 10 | 0710.4036 |
0710 | 0710.5801_arXiv.txt | We discuss the use of Sloan Digital Sky Survey (SDSS) {\it ugriz} point-spread function (PSF) photometry for setting the zero points of {\it UBVRI} CCD images. From a comparison with the Landolt (1992) standards and our own photometry we find that there is a fairly abrupt change in {\it B, V, R, \& I} zero points around $g, r, i \sim 14.5$, and in the $U$ zero point at $u \sim 16$. These changes correspond to where there is significant interpolation due to saturation in the SDSS PSF fluxes. There also seems to be another, much smaller systematic effect for stars with $g, r \gtrsim 19.5$. The latter effect is consistent with a small Malmquist bias. Because of the difficulties with PSF fluxes of brighter stars, we recommend that comparisons of {\it ugriz} and {\it UBVRI} photometry should only be made for unsaturated stars with $g, r$ and $i$ in the range 14.5 -- 19.5, and $u$ in the range 16 -- 19.5. We give a prescription for setting the $UBVRI$ zero points for CCD images, and general equations for transforming from {\it ugriz} to {\it UBVRI}. | When CCD images of a field are taken it is necessary to determine the photometric zero points from stars of known magnitudes. It is, however, not unusual for there to be no stars with $UBVRI$ photometry available. Fortunately, the Sloan Digital Sky Survey (SDSS) now provides homogenous $ugriz$ photometry for stars in a large fraction of the northern sky out of the plane of the Milky Way. Technical details of the SDSS are given in \citep{york00} and \citep{stoughton02}. The $ugriz$ system \citep{fukugita96} is significantly different from the widely used $UBVRI$ Johnson-Cousins system \citep{cousins76}, so it is necessary to transform between the two systems. A number of papers \citep{fukugita96,smith02,karaali03,karaali05,bilir05,jordi06,rodgers06,ivezic07,davenport07,bilir07} have considered the transformations between $ugriz$ and $UBVRI$ (see Section 6 for a discussion of these transformations). During the course of using SDSS $ugriz$ photometry to establish the zero points for comparison stars for photometry of active galactic nuclei (AGN), we noticed that the zero points were different for the fainter stars in a field than for the brighter stars. The difference was in the sense that stars with $g \lesssim 14$ were systematically brighter than predicted from the SDSS magnitudes. The difference did not seem to depend on the color of the stars and a check of the CCD used showed no evidence for non-linearity. A subsequent comparison of magnitudes of Landolt standards \citep{landolt92} revealed a similar zero-point difference for stars brighter or fainter than $r \sim 14$. In this note we report results of our investigation of the limitations of using SDSS photometry for bright stars, and give a prescription for setting zero points in CCD images taken through {\it UBVRI} filters. | The transformations we give in equations (1) -- (5) are consistent with the range of previously published transformations. Because our linear transformation equations are derived for a practical purpose of calibrating $UBVRI$ photometry, and are available for each of the Johnson-Cousins filters individually, our transformations are different in nature from those previously published. We briefly summarize here previously published transformations and discuss how they differ from the ones given above. \cite{fukugita96} give synthetic transformations from $UBVRI$ to $u'g'r'i'z'$. \cite{smith02} gave transformations between $UBVRI$ and $u'g'r'i'z'$ magnitudes observed with the Photometric Telescope (PT) at Apache Point Observatory for some filters and for colors. \cite{rodgers06} give improved color transformations between $u'g'r'i'z'$ and $UBVRI$ for main-sequence stars. They also consider higher-order color terms. It is important to note the difference between $u'g'r'i'z'$ and $ugriz$. This is discussed in \cite{smith07}. Additional technical details concerning the difference between the two systems as well as transformations between them are discussed in \cite{tucker06}. \cite{jordi06} give color transformations between $ugriz$ as observed with the SDSS 2.5-m telescope (rather than the PT) and $UBVRI$. Additional transformations are given by \cite{jester05}, \cite{karaali03}, \cite{karaali05},\cite{bilir05}, \cite{davenport07}, and \cite{bilir07}. Some of the transformations including \cite{ivezic07} consider polynomials in the color terms, but we found no need for higher-order terms for the restricted range of colors we consider. Note that the above cited transformations consider only colors, transform from $UBVRI$ to $ugriz$, are derived for the $u'g'r'i'z'$ system, or give transformations only for select Johnson-Cousins filters. In this note, our aim has been to give a practical means of photometrically calibrating $UBVRI$ CCD images. Researchers who are interested in astrophysical applications of SDSS photometry (such as the determination of spectroscopic parallaxes or fitting theoretical isochrones to HR diagrams) are referred to the above mentioned papers because the $ugriz$ to $UBVRI$ transformations depend on the luminosity class and metallicity of the stars. We have minimized these effects for zero-point setting by using a fairly tight color selection. | 7 | 10 | 0710.5801 |
0710 | 0710.5030_arXiv.txt | Optical high-resolution spectra of the R Coronae Borealis star V CrA at light maximum and during minimum light are discussed. Abundance analysis confirms previous results showing that V CrA has the composition of the small subclass of R Coronae Borealis (RCB) stars know as `minority' RCBs, i.e., the Si/Fe and S/Fe ratios are 100 times their solar values. A notable novel result for RCBs is the detection of the 1-0 Swan system $^{12}$C$^{13}$C bandhead indicating that $^{13}$C is abundant: spectrum synthesis shows that $^{12}$C/$^{13}$C is about 3 to 4. Absorption line profiles are variable at maximum light with some lines showing evidence of splitting by about 10 km s$^{-1}$. A spectrum obtained as the star was recovering from a deep minimum shows the presence of cool C$_2$ molecules with a rotational temperature of about 1200K, a temperature suggestive of gas in which carbon is condensing into soot. The presence of rapidly outflowing gas is shown by blue-shifted absorption components of the Na\,{\sc i} D and K\,{\sc i} 7698 \AA\ resonance lines. | R Coronae Borealis stars (here, RCBs) are a rare class of peculiar variable stars. The two defining characteristics of RCBs are (i) a propensity to fade at unpredictable times by up to about 8 magnitudes as a result of obscuration by clouds of soot, and (ii) a supergiant-like atmosphere that is very H-deficient and He-rich. The subject of this paper, V Coronae Australis (V CrA) is even a peculiar member of this class of peculiar stars. It is a `minority' R CrB. The distinction between majority and minority members was made first by Lambert \& Rao (1994) on the basis of an abundance analysis of warm RCBs. The minority RCBs are quite severely deficient in iron relative to the majority RCBs and to the Sun but some elements, particularly Si and S, have near-solar abundances in the minority (and majority) RCBs. This combination results in some very unusual abundance ratios, for example, the Si/Fe and S/Fe ratios of minority RCBs are approximately 100 times the solar ratios. V CrA also seems to be an especially lively producer of dust (Feast et al. 1997). The realization that V CrA is an unusual RCB led us to occasional spectroscopic monitoring of its optical spectrum. The current paper discusses high-resolution spectroscopic observations obtained on seven occasions between 1989 and 2003 when the star was either at or near maximum light or in decline by 3 to 5 magnitudes. A suitable spectrum at maximum light is subjected here to an abundance analysis. The previous analysis (Asplund et al. 2000) was based on a spectrum obtained during a shallow light minimum. This spectrum may have been contaminated by phenomena expected at minimum light (i.e., some lines may have been filled in partially by emission). In addition, our new spectra cover a broader bandpass at higher resolution and signal-to-noise ratio than the earlier spectrum. Other spectra show for the first time for V CrA line splitting indicative of the presence of an atmospheric shock. Spectra taken at minimum light are discussed in the context of our detailed studies of 1995-96 and 2003 minima of R CrB (Rao et al 1999; Rao, Lambert \& Shetrone 2006). \begin{figure} \epsfxsize=8truecm \epsffile{lcvcra1.ps} \caption{The visual (red dots) light curve of V CrA showing the several light minima during 1988--1998 period. Dates on which four spectroscopic observations were obtained are indicated by a dashed line (also see Table 1). Upper limits to the brightness are shown by green unfilled circles. The observations are from the AAVSO database.} \end{figure} | Differences between the spectra of RCBs at minimum light encourage us to continue our monitoring of V CrA and other RCBs. Studies of the initial stages of a decline should reveal clues to the trigger that sets off a decline. In this era when photometric observations by amateur astronomers are reported on the internet almost instantaneously, spectroscopic and other follow-up observations are limited by access to suitable telescopes. The advent of queue scheduling is smoothing the path to obtaining the follow-up observations. Development of a consortium of observers would also ease the situation. Observations of the RCBs in the deepest of minima are likely to provide novel data on their extended atmospheres, especially on the enigmatic broad lines which, the Na D lines apart, have been studied in detail with high-quality spectra only in the case of R CrB. For RCBs, aside from the three brightest stars R CrB, RY Sgr, and V854 Cen, a large telescope will be needed to acquire quality spectra at the faint magnitudes expected to reveal the broad lines. | 7 | 10 | 0710.5030 |
0710 | 0710.2426_arXiv.txt | Sgr A* is a source of strongly variable emission in several energy bands. It is generally agreed that this emission comes from the material surrounding the black hole which is either falling in or flowing out. The activity must be driven by accretion but the character of accretion flow in this object is an open question. We suggest that the inflow is dominated by the relatively low angular momentum material originating in one of the nearby group of stars. Such material flows in directly towards the black hole up to the distance of order of ten Schwarzschild radii or less, where it hits the angular momentum barrier which leads naturally to a flow variability. We study both the analytical and the numerical solutions for the flow dynamics, and we analyze the radiation spectra in both cases using the Monte Carlo code to simulate the synchrotron, bremsstrahlung and the Compton scattering. Our model roughly reproduces the broad band spectrum of Sgr A* and its variability if we allow for a small fraction of energy to be converted to non-thermal population of electrons. It is also consistent (for a range of viewing angles) with the strong constraints on the amount of circumnuclear material imposed by the measurements of the Faraday rotation. | % The character of the accretion flow onto a black hole depends on the initial angular momentum of the material. This angular momentum is specified by the outer boundary conditions which depend on the relative motion of the donor with respect to the black hole. This angular momentum corresponds to a certain circularization radius, i.e. the radius where this angular momentum is equal to the local Keplerian value. In binary systems the material comes from the secondary star and in general is possess high angular momentum due to the orbital motion. In Low mass X-ray binaries the flow proceeds through an inner Lagrange point and the circularization radius is a significant fraction of a Roche radius around a black hole, of order of $10^4 R_g$ ($R_g = GM/c$). In high mass X-ray binaries the accretion flow comes from the intercepted focused wind, so the circularization radius is smaller but still large, of order of $10^3 R_g$. In such case the inflowing material form an accretion disk around a black hole and the inflow proceeds due to the angular momentum transfer. Apart from the outermost region and the region close to the ISCO (innermost stable circular orbit), the distribution of the angular momentum is relatively smooth and not much different from the Keplerian law. The exact departures from Keplerian motion depends on the disk temperature (or more exactly, on the pressure distribution). In active galactic nuclei (AGN) the source of material is less specified. The material comes either from the stars (in the form of stellar winds) or from the gaseous phase of the galactic material. Bright AGN (quasars, Seyfert 1 galaxies) show the presence of accretion disks similar to the disks in binary systems so we can conclude that the angular momentum reaching the galactic center is high. In sources showing water maser activity we observe the outer parts of the disk directly, and in most sources the motion of the disk material is Keplerian. However, in weakly active galaxies like Sgr A* or giant elliptical galaxies we see no direct evidence of a disk. In Sgr A* the presence of the cold disk is actually excluded by the lack of eclipses of the stars which move very close to the central black hole and are systematically monitored since several years. Since in weakly active galaxies there are no direct observational arguments for any value of the angular momentum of the donated material and the location of material sources, three types of models are being considered: \begin{itemize} \item high angular momentum flow, with circularization radius of order of hundreds-thousands of $R_g$ \item low angular momentum flow, with circularization radius of order of a few $R_g$ \item spherical and quasi-spherical accretion, without angular momentum barrier. \end{itemize} The high angular momentum flow solutions for weakly active galaxies generally belong to ADAF (advection dominated accretion flow) family \citep{1977ApJ...214..840I,1994ApJ...428L..13N}, with possibly additional effects like outflows \citep{1999MNRAS.303L...1B} and convection. In this case the flow is not exactly Keplerian since the pressure gradients are important, but the local ratio of the angular momentum to the Keplerian angular momentum in most part of the flow is not wildly different from unity, and the angular momentum transfer (through viscosity) or angular momentum loss (through magnetic wind) at all radii is essential. Stationary solutions usually exist, and asymptotically the density of the flow approaches zero at infinity. In spherical and quasi-spherical flow there is no angular momentum barrier so the loss of angular momentum is not the necessary condition for the accretion to occur. Examples of such solutions are: purely spherical Bondi flow or flows where the angular momentum density is below the minimum angular momentum at the circular orbit around a black hole which is given by \begin{equation} l_{min} = 3 \sqrt{3} GM/c \end{equation} in case of Schwarzschild black hole; more general formula for a Kerr black hole can be found in \citep{1972ApJ...178..347B}. In Bondi solution \citep{1952MNRAS.112..195B,2003ApJ...591..891B} the outer boundary condition are specified by the density and the temperature of the uniform medium surrounding black hole at large distances. The flow velocity is zero at infinity, the inflow becomes transonic at the Bondi radius, and the supersonic flow reaches the black hole horizon. The Bondi radius depends significantly on the gas properties (e.g. politropic index; Bondi radius is of order of thousands of $R_g$ for relativistic flow with $\gamma = 4/3$ but is approaches zero if $\gamma \rightarrow 5/3$, typical for perfect fluid non-relativistic solution), but the accretion rate is much less sensitive to those assumptions.Purely Bondi flow has generally very low radiative efficiency so it cannot reproduce the observed luminosity in most weakly active galaxies \citep{2006A&A...450...93M}. If the accreting material at the outer boundary condition has certain angular momentum $l < l_{min}$, the dynamics of the flow is slightly modified in comparison with Bondi flow and the flow is not spherically symmetric any more but the stationary solution for the flow always exists. The intermediate case of low angular momentum the situation is the most complex as initially the flow behaves as the Bondi flow but close to the black hole the flow starts suddenly to feel the angular momentum barrier \citep{1981ApJ...246..314A}. In this case analytical stationary solutions frequently do not exist. In numerical solutions the flow is variable and does not reach a stationary solution in the computing time. If the angular momentum of the donated material is also a subject of changes (e.g. the result of the stellar motion), a truly stationary solution indeed can never be reached for physical reasons. In the case of Sgr A* the available spatial resolution is the highest and we can have the best insight into the sources of material \citep{2007IAUS..238..173G}. Therefore, in the present paper we concentrate specifically on this source and we argue that the low angular momentum flow is an interesting and promising option for the flow description. | % Low angular momentum accretion flow is a promising scenario for the accretion onto Sgr A* due to its natural variability pattern. The flow is slightly more energetically efficient than the purely spherical Bondi flow and can reproduce both the required level of the luminosity and is consistent with the data on Faraday rotation measure. The overall broad band spectra are also roughly reproduced if a fraction of energy is allowed to be converted the non-thermal population of electrons. The current results are therefore encouraging, and the further work is in progress. \ack% The present work was supported by the Polish Grant 1P03D~008~29 and the Polish Astroparticle Network 621/E-78/SN-0068/2007. | 7 | 10 | 0710.2426 |
0710 | 0710.5354_arXiv.txt | We simulate the collisional formation of a ring galaxy and we integrate its evolution up to 1.5 Gyr after the interaction. About $100-200$ Myr after the collision, the simulated galaxy is very similar to observed ring galaxies (e.g. Cartwheel). After this stage, the ring keeps expanding and fades. Approximately $0.5-1$ Gyr after the interaction, the disc becomes very large ($\sim{}100$ kpc) and flat. Such extended discs have been observed only in giant low surface brightness galaxies (GLSBs). We compare various properties of our simulated galaxies (surface brightness profile, morphology, HI spectrum and rotation curve) with the observations of four well-known GLSBs (UGC6614, Malin 1, Malin 2 and NGC7589). The simulations match quite well the observations, suggesting that ring galaxies could be the progenitors of GLSBs. This result is crucial for the cold dark matter (CDM) model, as it was very difficult, so far, to explain the formation of GLSBs within the CDM scenario. | Ring galaxies are one of the most intriguing categories of peculiar galaxies. About 280 galaxies have been classified as ring-like in the Catalogue of Southern Peculiar Galaxies and Associations (CPGA; Arp \& Madore 1987). They have commonly been divided in two different classes, P-type and O-type ring galaxies (Few \& Madore 1986). The former consists of galaxies where the nucleus is often off-centre and the ring is quite knotty and irregular, while the latter includes objects where the nucleus is central and the ring is regular. Even if for some objects such classification is ambiguous, the number of P-type and of O-type ring galaxies are roughly comparable. The differences among the two classes are probably connected with the formation mechanism of such galaxies: for P-type ring galaxies a collisional origin has been proposed (Lynds \& Toomre 1976; Theys \& Spiegel 1976; Appleton \& Struck-Marcell 1987a, 1987b; Hernquist \& Weil 1993; Mihos \& Hernquist 1994; Appleton \& Struck-Marcell 1996; Struck 1997; Horellou \& Combes 2001), as most of them have at least one nearby companion (Few \& Madore 1986); whereas O-type ring galaxies can be resonant (R)S galaxies (de Vacouleurs 1959). Simulations of galaxy collisions leading to the formation of P-type ring galaxies show that the ring phase is quite short-lived (Hernquist \& Weil 1993; Mihos \& Hernquist 1994; Horellou \& Combes 2001; Mapelli et al. 2007, hereafter M07): the simulated disc galaxy develops a ring similar to the observed ones $\approx{}100$ Myr after the collision with the intruder; but the ring remains dense and clearly visible only for the first $\approx{}300$ Myr (M07). Thus, ring galaxies are expected to rapidly evolve into something else. Up to now, only analytic models for ring waves have been adopted to study the late stages of the ring galaxy life (Struck-Marcell \& Lotan 1990; Appleton \& Struck-Marcell 1996). Neither observations nor simulations have been carried on to investigate the fate of ring galaxies after the ring phase, leaving a lot of uncertainties. This paper aims at describing the subsequent stages of the evolution of a ring galaxy (up to $\sim{}1.5$ Gyr), by means of SPH/$N$-body simulations. Our simulations suggest a possible link between old ring galaxies and the so-called giant low surface brightness galaxies (GLSBs). The GLSBs are low surface brightness galaxies (LSBs) characterized by unusually large extension of the stellar and gaseous discs (up to $\sim{}100$ kpc; Bothun et al. 1987; Impey \& Bothun 1989; Bothun et al. 1990; Sprayberry et al. 1995; Pickering et al. 1997; Moore \& Parker 2007) and (often) by the presence of a normal stellar bulge (Sprayberry et al. 1995; Pickering et al. 1997). The existence of GLSBs has always been puzzling. Galaxy formation simulations within the cold dark matter (CDM) model have serious difficulties in producing realistic disc galaxies. In such simulations, too much angular momentum is lost during the assembly of objects, producing discs that tend to be too compact and dense. While a 'Milky Way-like' galaxy can be reproduced by high-resolution CDM simulations (provided that its merging history is fairly quiet and that heating by supernovae is properly accounted for; Governato et al. 2007), the extended discs of GLSBs require that much more angular momentum is preserved during the hierarchical build up. Such extended discs are thus beyond the reach of current galaxy formation simulations, and it is unclear whether improving the realism of such simulations will solve the problem. Hoffman, Silk \& Wyse (1992) proposed that GLSBs form from rare density peaks in voids. However, most of currently known GLSBs do not appear to be connected with voids and often show interacting companions (Pickering et al. 1997). A more promising scenario is the formation of GLSBs from massive disc galaxies due to a bar instability: a large scale bar can redistribute the disc matter and significantly increase the disc scale-length (Noguchi 2001; Mayer \& Wadsley 2004). In this case the redistribution of angular momentum by the bar instability could counteract the natural tendency of hierarchical assembly to remove angular momentum from the disc material (Kaufmann et al. 2007). Bar instabilities, however, normally do not increase the disc scale length by more than a factor of 2-2.5 (Debattista et al. 2006; Kaufmann et al. 2007). In this paper we show that also the propagation of the ring in an old collisional ring galaxy can lead to the redistribution of mass and angular momentum in both the stellar and gas component out to a distance of $\sim{}100-150$ kpc from the centre of the galaxy, producing features (e.g. the surface brightness profile, the star formation, the HI emission spectra and the rotation curve) which are typical of GLSBs. | In this paper we presented a numerical model of ring galaxy evolution. About $100-200$ Myr after the collision with the intruder, the target disc galaxy evolves into a ring galaxy, similar to Cartwheel. Afterward, the ring progressively expands and fades. After $\approx{}0.5-1.0$ Gyr the ring is no longer distinguishable from the disc, its surface density is more than 1 order of magnitude lower than in the Cartwheel phase, and the disc extends up to $\sim{}100$ kpc. We showed that simulated ring galaxies in the late stages of their dynamical evolution ($\gtrsim{}500$ Myr) are very similar to the observed GLSBs. The $R$-band surface brightness profile, the SFR, the HI spectra and also the rotation curves of four GLSBs are well reproduced by the simulations. This result is unlikely due to a simple coincidence. If all GLSBs were originated by the evolution of P-type ring galaxies, their current number density would be comparable to the observed number density of P-type ring galaxies, i.e. $\sim{}5.4\times{}10^{-6}\,{}h^3\,{}{\rm Mpc}^{-3}$ (where $h$ is the Hubble constant; Few \& Madore 1986). Since we know $\sim{}17$ galaxies which can be classified as GLSBs (Sprayberry et al. 1995; Pickering et al. 1997) and which are within $\sim{}340$ Mpc from our Galaxy, the lower limit of the current number density of GLSBs is $\sim{}2.7\times{}10^{-7}\,{}h^3\,{}{\rm Mpc}^{-3}$, i.e. about one order of magnitude lower than the density of ring galaxies. This can imply either that a large fraction of GLSBs have not been detected yet, or that only the $\sim{}$5 per cent of ring galaxies ends up into a GLSB. The former scenario is quite realistic, as magnitude-limited surveys are strongly biased against GLSBs (Bothun et al. 1997). The latter hypothesis is also likely, as ring galaxies can end their life in other ways. For example, if the intruder is not sufficiently massive, the density wave is not strong enough to produce a GLSB, or, if the relative velocity is not sufficiently high, the companion can come back and merge or disrupt the propagating wave. Furthermore, recycled dwarf galaxies might also form from the debris of old collisional ring galaxies (e.g. the case of NGC~5291; Bournaud et al. 2007). Furthermore, our model can explain the formation of GLSBs within the context of the CDM scenario, in which very extended discs are hardly explained as a result of normal hierarchical assembly. Hence it will be quite important to check the consistency of our model with future observations. For example, the available HI data (Pickering et al. 1997) show that most of GLSBs have strong non-circular motions. From the published data it is not possible to understand whether these motions are consistent with the expansion and/or the falling back of gas in the ring. Thus, it is very important to make new radio observations of GLSBs or, at least, re-examine the archival data. Another important point to address is how many GLSBs have a nearby companion which could be the intruder. Among the 4 GLSBs considered here, NGC~7589 has a well-known interacting companion (Pickering et al. 1997). Furthermore, UGC~6614 is surrounded by other galaxies with approximately the same redshift (e.g. KUG 1136+173, AGC~211143 and CGCG~097$-$034), but no studies have been done to establish whether they are in the same group as UGC6614. Finally, the surroundings of Malin~1 are populated by many companion candidates, but no redshift measurements are available at present (Sprayberry et al. 1995). A further issue raised by this paper is whether there are objects in the intermediate stage between ring galaxies and GLSBs. The ring galaxy Arp~10 has been thought to be a relatively old ring galaxy, because its rings are not as prominent as those of other ring galaxies (e.g. Cartwheel) and because its SFR, from H$\alpha{}$ observations, appears quite low (Charmandaris, Appleton \& Marston 1993). However, Bizyaev, Moiseev \& Vorobyov (2007) have shown that the SFR of Arp~10, from far-infrared observations , is $\sim{}10-21\,{}M_\odot{}$ yr$^{-1}$ (analogous to the one of Cartwheel; Mayya et al. 2005), and that the collision which produced the ring occurred only $\sim{}85$ Myr ago. Then, both current observational data and theoretical models are not sufficient to support the idea that Arp~10 is an old ring galaxy\footnote{Even our simulations cannot solve the problem, as they indicate that the surface brightness profile of Arp~10 can be matched by a $\sim{}80$ Myr old ring galaxy (in agreement with Bizyaev et al. 2007), as well as by a $\sim{}300$ Myr old ring galaxy. In particular, the surface brightness profile of Arp~10 is matched by a $80$ Myr old ring galaxy with the same initial conditions as run A but with $R_d=8.8$ kpc. Nice agreement with the observations is obtained also for a $300$ Myr old ring galaxy with the same initial conditions as runs B in table 1 of M07, but with $R_d=13.2$ kpc.}. On the other hand, one of the four GLSBs considered in this paper, UGC6614, shows (both in H$\alpha{}$, in HI and in optical) structures which look like the remnants of the inner ring. Thus, UGC6614 could be in an intermediate, connecting stage between GLSBs and ring galaxies. It would be crucial to study more deeply UGC6614, as well as to search for other galaxies which could represent the intermediate phase between GLSBs and ring galaxies. | 7 | 10 | 0710.5354 |
0710 | 0710.5162_arXiv.txt | We present recent results from a Keck study of the composition of the Galactic bulge, as well as results from the bulge Bulge Radial Velocity Assay (BRAVA). Culminating a 10 year investigation, Fulbright, McWilliam, \& Rich (2006, 2007) solved the problem of deriving the iron abundance in the Galactic bulge, and find enhanced alpha element abundances, consistent with the earlier work of McWilliam \& Rich (1994). We also report on a radial velocity survey of {\sl 2MASS}-selected M giant stars in the Galactic bulge, observed with the CTIO 4m Hydra multi-object spectrograph. This program is to test dynamical models of the bulge and to search for and map any dynamically cold substructure in the Galactic bulge. We show initial results on fields at $-10^{\circ} < l <+10^{\circ}$ and $b=-4^{\circ}$. We construct a longitude-velocity plot for the bulge stars and the model data, and find that contrary to previous studies, the bulge does not rotate as a solid body; from $-5^{\circ}<l<+5^{\circ}$ the rotation curve has a slope of $\approx 100\ km\ s^{-1}$ and flattens considerably at greater $l$ and reaches a maximum rotation of $45\ {km\ s^{-1}}$ (heliocentric) or $\sim 70\ {km\ s^{-1}}$ (Galactocentric). This rotation is slower than that predicted by the dynamical model of Zhao (1996). | 7 | 10 | 0710.5162 |
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0710 | 0710.1431_arXiv.txt | We present the first high spatial resolution near-infrared direct and polarimetric observations of \pars{}, obtained with the VLT/NACO instrument. We complemented these measurements with archival infrared observations, such as HST/WFPC2 imaging, HST/NICMOS polarimetry, Spitzer IRAC and MIPS photometry, Spitzer IRS spectroscopy as well as ISO photometry. Our main conclusions are the following: (1) we argue that \pars{} is probably an FU\,Orionis-type object; (2) \pars{} is not associated with any rich cluster of young stars; (3) our measurements reveal a circumstellar envelope, a polar cavity and an edge-on disc; the disc seems to be geometrically flat and extends from approximately 48 to 360 AU from the star; (4) the SED can be reproduced with a simple model of a circumstellar disc and an envelope; (5) within the framework of an evolutionary sequence of FUors proposed by \citet{green} and \citet{quanz}, \pars{} can be classified as an intermediate-aged object. | The most important process of low-mass star formation is the accretion of circumstellar material onto the young star. According to current models, the mass accumulation is time-dependent with alternating episodes of high and low accretion rates \citep{vorobyov,boley}. The high phase may correspond to FU\,Orionis objects (FUors), a unique class of low-mass pre-main sequence stars that have undergone a major outburst in optical light of $4\,$mag or more \citep{herbig77}. Currently about 20 objects have been classified as FU\,Ori-type \citep[for a list see][]{fuors}. Records of the outbursts are not available for all of these objects: some were classified as FUors since they share many spectral properties with the FUor prototypes. According to the most widely accepted model \citep{kh91}, the observed spectral energy distribution (SED) is reproduced by a combination of an accretion disc and an envelope. The energy of the outburst originates from the dramatic increase of the accretion rate. On the other hand, \citet{herbig2003} favour a model in which the outburst occurs in an unstable young star rotating near breakup velocity. \pars{} is a system consisting of a central object, HBC\,687 ($\alpha_{2000}$ = 19$^{\rm h}$ 29$^{\rm m}$ 0\fs87, $\delta_{2000}$ = 9\degr{} 38\arcmin{} 42\farcs{}69), and an extended nebula first listed in the catalogue of \citet{par}. Though no optical outburst was ever observed, \pars{} was classified as a FUor on the basis of its optical spectroscopic and far-infrared properties \citep{sn92}. According to \citet{henning} \pars{} is situated in a molecular cloud named ``Cloud\,A'', very close to the Galactic plane. Indeed, the distance of 400\,pc and the radial velocity of $v_{lsr} = +27 \pm 15$\,km s$^{-1}$ \citep{sn92} agree very well with the respective values for Cloud\,A \citep[$v_{lsr} = +27 \pm 4$\,km s$^{-1}$, $d = 500 \pm 100$\,pc,][]{dame85}. Recently \citet{quanz} questioned the FUor nature of \pars{} referring to mid-infrared spectral properties, which better resemble those of post-AGB stars (this issue will be discussed in Sect.~\ref{sec:discussion}). The elongated nebulosity around the central source extends ${\approx}\,1\,$arcmin to the north, while it is less developed to the south. It is visible in the optical and the near-infrared, though the source becomes unresolved at $10.8$ and $18.2\,\mu$m \citep{polom}. Polarimetric observations by \citet{draper} revealed a centrosymmetric polarization pattern and an elongated band of low polarization perpendicular to the axis of the bright nebula. \citet{bm90} interpreted the polarization map of \pars{} in terms of multiple scattering in flattened, optically thick structures and derived an inclination angle of 80-85\degr{} and a size of 30 $\times$ 8 arcsec for this disc-like structure. The existence of a disc is also supported by the discovery of a short bipolar outflow oriented along the polar axis of the nebula \citep{sn92}. Emission at submm wavelengths was detected by \citet{henning} and \citet{polom}. The surroundings of \pars{} have been searched for companions several times, but close-by stars seen in J, H and K-band images proved to be field stars \citep{li}, and no close-by sources have been found at longer wavelengths so far \citep{polom}. Recently several studies were published on the circumstellar environment of FUors using high-resolution infrared and interferometric techniques (e.g.~\citealt{green, malbet1998, malbet2005, millan-gabet, quanz}). Such observations could prove or disprove the basic assumptions of FUor models. In this paper we present the highest resolution imaging and polarimetric data yet on \pars{} from the VLT/NACO instrument, allowing the inspection of the circumstellar material and disc at spatial scales of ${\approx}\,0.07$ arcsec. The edge-on geometry of \pars{} represent an ideal configuration to separate the components of the circumstellar environment (disc, envelope) and understand their role in the FUor phenomenon. | \label{sec:discussion} \subsection{\pars{}: an FU\,Orionis-type star} Although no optical outburst was ever observed, \pars{} shares many properties characteristic of FUors. Its spectral type in the literature ranges from A5 to F8 supergiant \citep[e.g.~][]{sn92}, similar to the typical FUor spectral types (F--G supergiant). In addition, it shows strong infrared excess and drives a bipolar outflow, similarly to many FUors. However, in a recent paper by \citet{quanz} the pre-main sequence nature and the FUor status of \pars{} were questioned, mainly because its PAH emission bands are untypical for young stars. They suggest that \pars{} is either an intermediate mass FUor object, or an evolved star sharing typical properties with post-AGB stars. An important parameter which may discriminate between a pre-main sequence object and a post-AGB star is the luminosity. The typical luminosity of post-AGB stars is in the order of $10^3-10^4\,$L$_{\odot}$ \citep{kwok}, while FUors are typically fainter than a few $100\,$L$_{\odot}$. The luminosity of \pars{}, $10\,$L$_{\odot}$ (see Sect.~\ref{sec:sed}), is at least a factor of 100 lower than that of post-AGB stars. Even adopting the largest distance estimate mentioned in the literature (1800\,pc, \citealt{sn92}), the luminosity of \pars{} is still too low. The observed proper motion of the Herbig-Haro knots, however, prefers the lower distance (400\,pc), otherwise the outflow velocities (${\approx}\,2400\,$km s$^{-1}$ at 1800\,pc) would be unusually high for a young stellar object \citep[typically up to $600\,$km s$^{-1}$, see e.g.~][]{mundt, hessman}. Moreover, the existence of Herbig-Haro outflows are usually tracers of low-mass star formation \citep{hh}. In the following discussion, we consider \pars{} to be a FUor and we discuss its properties in the context of young eruptive stars. Nevertheless we note that most of our results on the morphology and circumstellar structure are valid regardless of the nature of the central object. \subsection{\pars{}: an isolated young star?} \label{sec:isolated} \begin{figure} \centering \includegraphics[angle=0,width=0.95\linewidth]{kospal_fig10.ps} \caption{2MASS colour-colour diagram of 188 sources in an area of $5.3\,{\times}\,5.3$ arcmin centred on \pars. The main-sequence and giant branch are marked by solid lines \citep{koornneef}, the reddening path with dashed lines \citep{cardelli} and the T\,Tauri locus with dotted line \citep{meyer}.} \label{fig:cc2mass} \end{figure} \begin{figure} \centering \includegraphics[angle=0,width=0.95\linewidth]{kospal_fig11.ps} \caption{IRAC colour-colour diagram of 100 sources located in the same area as in Fig.~\ref{fig:cc2mass}. The gray square in the upper right corner marks the approximate domain of Class II sources \citep{allen}.} \label{fig:ccirac} \end{figure} \begin{table} \caption{VLT/NACO photometry for the two closest stars. Uncertainties are about 0.1 mag.} \label{tab:twostars} \centering \begin{tabular}{c c c c c} \hline $\alpha_{2000}$ & $\delta_{2000}$ & H mag & K$_S$ mag & L$^{\prime}$ mag \\ \hline 19$^{\rm h}$ 29$^{\rm m}$ 0\fs96 & 9\degr{} 38\arcmin{} 42\farcs11 & 20.2 & 19.1 & $>$14.6 \\ 19$^{\rm h}$ 29$^{\rm m}$ 0\fs70 & 9\degr{} 38\arcmin{} 44\farcs78 & 20.6 & 19.5 & $>$14.4 \\ \hline \end{tabular} \end{table} \pars{} is situated close to the Galactic plane ($l=45.8^{\circ}$, $b=-3.8^{\circ}$), in a molecular cloud called Cloud\,A. This cloud was identified by \citet{dame85} in their CO survey of molecular clouds in the northern Milky Way. The cloud occupies an area of 8 square degrees in the sky, and the only known young star associated with it is the T\,Tauri star AS\,353 \citep{dame85}. In order to check whether FUors are usually associated with star forming regions, we searched the literature and found that most FUors are located in areas of active star formation \citep[e.g.][]{henning}. To find out whether there is star formation in the vicinity of \pars{}, we searched for pre-main sequence stars. For this purpose we constructed a 2MASS J$-$H vs.~H$-$K$_S$ and an IRAC [3.6]$-$[4.5] vs.~[5.8]$-$[8.0] colour-colour diagram for sources found in our $5\farcm3\,{\times}\,5\farcm3$ IRAC field of view (Fig.~\ref{fig:cc2mass}, \ref{fig:ccirac}). Our selection criteria in the case of 2MASS was S/N ${>}\,10$ and uncertainties ${<}\,0.1\,$mag in all J, H and K$_S$ bands, while in the case of IRAC S/N ${>}\,3$ and detectability at all four bands were required. The 2MASS diagram revealed that most of the nearby objects are reddened main sequence or giant stars. On the IRAC diagram, however, there are three objects (apart from \pars{} itself), which display infrared excess at $8\,\mu$m (marked by A, B and C in Fig.~\ref{fig:ccirac}). According to the classification of \citet{allen}, Class II sources exhibit colours of [3.6]$-$[4.5] ${>}\,0.0$ and [5.8]$-$[8.0] ${>}\,0.4$, thus one of these stars (A) might be a Class II source, while B and C are more likely Class III/main sequence sources. The nature of source A and its possible relationship to \pars{} is yet to be investigated. Nevertheless, \pars{} seems to be rather isolated compared to most FUors and certainly not associated with any rich cluster of young stellar objects. We also searched for possible close companions of \pars{} in the WFPC2 and NACO direct images. In order to establish a detection limit for source detection, we measured the sky brightness on the NACO images (before sky-subtraction), and estimated a limiting magnitude for each filter. The resulting values are $22.8$, $21.6$, and $15.2\,$mag in H, K$_S$ and L$^{\prime}$, respectively. In case of the HST/WFPC2 image, the larger ($80\,{\times}\,80$ arcsec) field of view made it possible to estimate a limiting magnitude using star counts; the resulting value is $23.5\,$mag. Due to the bright reflection nebula, the detection limit is somewhat lower close to the star. The two closest objects we found are the following: one star to the southeast, at a distance of $1.4$ arcsec (560\,AU at 400\,pc), and another one to the northwest, at a distance of $3.3$ arcsec (1320\,AU at 400\,pc). These sources are marked with arrows in Fig.~\ref{fig:images}. Neither of the stars are visible at $3.8$ or $0.814\,\mu$m, although we can give an upper limit for their L$^{\prime}$ brightness. Their positions and photometry are given in Table~\ref{tab:twostars}. As these sources are very red, they can equally be heavily reddened background stars, or stars with infrared excess (indicating that they might be associated with \pars{}). Supposing that they are reddened main sequence stars, one can estimate an extinction of A$_V\,{\approx}\,10\,{-}\,15\,$mag. Further multifilter observations may help to clarify the nature of these objects and their possible relationship to \pars{}. \subsection{The circumstellar environment of \pars{}} \begin{figure} \centering \includegraphics[angle=0,width=0.45\linewidth]{kospal_fig12.ps} \caption{Sketch of the morphology of circumstellar material around \pars{}, overlaid on the HST/WFPC2 image. The central star is surrounded by an edge-on disc. Perpendicular to the disc, the star drives a bipolar outflow that excavates an outflow cavity in the dense circumstellar material. Light from the central star illuminates the walls of the cavity.} \label{fig:morph} \end{figure} \begin{figure} \centering \includegraphics[angle=90,width=0.95\linewidth]{kospal_fig13.ps} \caption{Brightness profiles of \pars{} at $0.8\,\mu$m and in the H and K$_S$ bands. Dotted lines mark Hubble's relation.} \label{fig:metszet} \end{figure} \begin{figure} \centering \includegraphics[angle=0,width=0.95\linewidth]{kospal_fig14.ps} \caption{South-north cut at $0.6$ arcsec east from the star. {\it Solid line:} total intensity; {\it dashed line:} polarized intensity.} \label{fig:cutok} \end{figure} \begin{figure*} \centering \includegraphics[angle=90,width=0.95\linewidth]{kospal_fig15.ps} \caption{VLT/NACO polarization map of \pars{} overlaid on H-band total intensity contours. The circle in the upper right corner displays the FWHM of the polarization measurement. Polarization vectors are displayed at full resolution, only showing the central low polarization band, where the polarization vectors are aligned. Such arrangement is expected when multiple scattering occurs in an edge-on disc.} \label{fig:kiskep} \end{figure*} The appearance and polarization properties of the nebula around \pars{} can be understood in the following way: the star drives an approximately north-south oriented bipolar outflow, which had excavated a conical cavity in the dense circumstellar material (Fig.~\ref{fig:morph}). The star illuminates this cavity and the light is scattered towards us mainly from the walls of the cavity. The outflow direction is perpendicular to an almost edge-on dense circumstellar disc. This picture is supported by the following facts: (a) the centrosymmetric polarization pattern is characteristic of reflection nebulae with single scattering; (b) the morphology and limb brightening suggest a hollow cavity (as opposed to an ``outflow nebula'', where the lobes are composed of dense material ejected by the central source); and (c) the low-polarization lane across the star strongly suggests the presence of an edge-on circumstellar disc, where multiple scattering occurs. In general the \pars{} system shows similarities to the NGC\,2261 nebula associated with R\,Mon. This object also consists of a northern cometary nebula and a southern jetlike feature \citep{warren}. \subsubsection{Envelope/Cavity} \label{sec:env} We characterised the opening of the upper lobe by marking the ridge along the northeastern and northwestern arcs which we interpret as the walls of the cavity. As viewed from the star northwards, the cavity starts as a cone with an opening angle of ${\approx}\,60^{\circ}$, giving the nebula in Fig.~\ref{fig:images} a characteristic equilateral triangle-shape. Farther away from the star the cavity deviates from the conical shape, becomes narrower. The whole cavity occupies an area of $8\,000\,{\times}\,24\,000\,$AU (at a distance of $400\,$pc). The sharp outer boundary of the nebula implies a significant density contrast between the cavity and the surrounding envelope. In Fig.~\ref{fig:metszet} we plotted radial brightness profiles at $0.8\,\mu$m and in the H and K$_S$ bands, starting from the star northwards. The slope of the intensity profiles in the inner ${\sim}\,1.6''$ ($640\,$AU) follows closely Hubble's relation \citep{hubble}, i.e.~the brightness is proportional to r$^{-2}$. This trend can be clearly followed above the noise out to about $5\,$arcsec in our H and K$_S$ images, giving an estimate of $2000\,$AU for the outer size of the envelope. The fact that the brightness profiles follow Hubble's relation implies that the nebula is produced by isotropic single scattering, in accordance with the centrosymmetric polarization pattern and the high degree of polarization. Around ${\sim}\,1.6''$ the profiles becomes steeper, probably due to decreased density in the inside of the cavity. Similar steepening in the outer part of the nebula was already mentioned by \citet{li}. The profiles are similar at all observed wavelengths and do not show significant dependence on the position angle. \subsubsection{Disc} \label{sec:disc} As mentioned in Sect.~\ref{sec:pol}, both the NACO and NICMOS polarimetric images show a lane of low polarization oriented nearly east-west across the star (Fig.~\ref{fig:pol}). The drop in the degree of polarization can also be clearly seen in Fig.~\ref{fig:cutok}, where a north-south cut at $0\farcs6$ east to the star is plotted. Following \citet{bm90}, we interpret the low polarization by multiple scattering in an edge-on disc (possible other explanations for the origin of the low polarization areas are discussed in e.g.~\citealt{lr}). According to models of such circumstellar structures \citep[e.g.][]{whitney, fischer} the polarization vectors are oriented parallel to the disc plane. As can be seen in Fig.~\ref{fig:kiskep}, despite the low degree of polarization, the predicted alignment of the vectors can be clearly seen in the case of \pars{} too. The NICMOS polarization map shows a similar effect. The dark lane in Fig.~\ref{fig:pol} can be followed inwards to as close as $48\,$AU from the star. This is an upper limit for the inner radius of the circumstellar disc. An interesting feature of the mentioned models is two depolarized areas on either side of the central star, which mark the outer end points of an edge-on disc. The depolarization is due to a transition from the linearly aligned to the centro-symmetic polarization pattern. These depolarized areas can be seen in Fig.~\ref{fig:vector} bottom, on both sides at about $0.9''$ from the star. At the distance of \pars{} this corresponds to $360\,$AU, and can be adopted as the outer radius of the dense part of the disc, where multiple scattering at near-infrared wavelengths is dominant. It is interesting that the NACO map does not show clear depolarized areas at the same position, but seems to resolve the transition in vector orientation, while the larger beam of NICMOS averaged the differently oriented vectors, resulting in depolarized spots. The polarized intensity and degree of polarization images in Fig.~\ref{fig:pol} suggest a slight asymmetry in the dense disc: the eastern (left) side is straight, while the western (right) side shows a kink and also has a different position angle than that on the other side. The thickness of the disc can be measured on the area where the polarization vectors are aligned (Fig.~\ref{fig:kiskep}). This approximately corresponds to the area where the degree of polarization is below ${\approx}\,10\%$. The resulting thickness is approximately $0.1''$ ($40\,$AU) to the east and is somewhat larger, $0.2''$ ($80\,$AU), to the west. One should note that, since these values are close to the spatial resolution of the polarimetric images, these numbers should be considered as upper limits for the thickness of the circumstellar disc around \pars{}. They are upper limits also because if the inclination is not exactly 90$^{\circ}$, the thickness of the disc can be even less. The thickness does not show significant increase with radial distance, suggesting the picture of a flat, rather than a flared disc, at least considering the dense, multiple-scattering part. The smooth brightness and polarization distribution between the disc and the surrounding envelope (Fig.~\ref{fig:cutok}), however, implies that there is a continuous density transition between the two components. From the ratio of the horizontal to vertical sizes of the disc a lower limit of $84^{\circ}$ for the inclination of the system (the angle between the normal of the disc and the line of sight) can be derived. \citet{fischer} computed a grid of polarization maps of young stellar objects with the aim of helping the interpretation of polarimetric imaging observations. They consider five different models, four with massive, self-gravitating discs and one with a massless Keplerian disc. Since the mass of circumstellar material of \pars{} derived from submillimetre observations is relatively low (${<}\,0.3\,$M$_{\odot}$, \citealt{henning, polom, sw, hillen}), the most appropriate model for our case is a Keplerian disc. Indeed the polarization pattern as computed by \citet{fischer} for an inclination of 87$^{\circ}$ (their Fig.~1) looks remarkably similar to our Fig.~\ref{fig:vector} (the differences might be explained by the narrower cavity and flatter disc of \pars{}). Thus, the geometry and structure assumed by \citet{fischer} in their Keplerian model could be a good starting point for further radiative transfer modelling of \pars{}. \subsection{Modelling the circumstellar environment} \label{sec:modelling} \begin{table} \centering \caption{Model parameters. Parameters in italics are fixed, while the others were fitted.} \label{tab:par} \begin{tabular}{lcc} \hline Parameter & Variable & Value \\ \hline Inner disc radius & $R_1$ & 3.5\,$R_\odot$ \\ {\it Outer disc radius} & $R_2$ & {\it 360\,AU} \\ Temperature at 1 AU & $T_{\rm d,0}$ & 285\,K \\ {\it Power-law index for temperature} & $q_{\rm d}$ & {\it 0.75} \\ Power-law index for surface density & $p_{\rm d}$ & 1.6 \\ {\it Disc mass} & $M_{\rm d}$ & {\it 0.02\,M$_\odot$} \\ {\it Inclination} & $i$ & {\it 86$^\circ$} \\ Inner envelope radius & $R_3$ & 5.4\,AU \\ {\it Outer envelope radius} & $R_4$ & {\it 2000\,AU} \\ Temperature at 5 AU & $T_{\rm e,0}$ & 368\,K \\ {\it Power-law index for temperature} & $q_{\rm e}$ & {\it 0.4} \\ Power-law index for surface density & $p_{\rm e}$ & 0.4 \\ Envelope mass & $M_{\rm e}$ & 0.02\,M$_\odot$ \\ {\it Interstellar Extinction} & $A_{V}$ & {\it 2\,mag}\\ \hline \end{tabular} \end{table} Our observations provide some direct measurements of the geometry of the circumstellar structure (disc size and thickness, inclination, envelope size). In the following we discuss the consistency of this picture with the observed SED. Our approach is to construct a simple disc+envelope model, in which we fix those parameters whose values are known from our NACO observations or from other sources (outer disc radius from this work; power-law index for disc temperature from \citealt{shakura}; disc mass was set in order to ensure that the whole disc is optically thick; inclination from this work; outer envelope radius from this work, the power-law index for envelope temperature is a typical value for optically thin envelopes containing larger than interstellar grains, e.g.~\citealt{hartmannkonyv}, Eqn.~4.13; interstellar extinction from \citealt{hillen}). Then we check whether the SED can be fitted by tuning the remaining parameters. We adopted an analytical disk model \citep{adams}, which has been successfully used to model FUors \citep{quanz2, v1647ori_midi}. Our model consists of two components, an optically thick and geometrically thin accretion disc \citep{shakura} and an optically thin envelope (no cavity is assumed). No central star is included in the simulation, partly because in outbursting FUors the star's contribution is negligible compared to that of the inner disc \citep{hk96}, and partly because of the edge-on geometry where the star is obscured by the disc. This assumption is supported by the fact that the shape of the SED at optical wavelengths is broader than a stellar photosphere. The model also does not take into account internal extinction and light scattering, thus it cannot reproduce any of the near-IR imaging and polarimetric observations. The temperature and surface density distribution in the disc are described by power-laws: \begin{equation} T(r)=T_{{\rm d}, 0}\left(\frac{r}{1\,{\rm AU}}\right)^{-q_{\rm d}}, \end{equation} \begin{equation} \Sigma(r)=\Sigma_{{\rm d},0}\left(\frac{r}{1\,{\rm AU}}\right)^{-p_{\rm d}}. \end{equation} Similar power-laws were assumed for the envelope. The observed flux at a specific frequency is given by \begin{eqnarray} F_\nu &=& \frac{\cos{i}}{D^2}\int_{\rm R1}^{\rm R2}2\pi r(1-e^{\frac{-\Sigma_{{\rm d}}\kappa_\nu}{\cos{i}}})B_\nu(T_{\rm d}){\rm d}r + \\ & & \frac{1}{D^2}\int_{\rm R3}^{\rm R4}2\pi r(1-e^{-\Sigma_{{\rm e}}\kappa_\nu})B_\nu(T_{{\rm e}}){\rm d}r. \end{eqnarray} The first term describes the emission of the accretion disc, the second term describes the radiation of the optically thin envelope. For the dust opacity we used a constant value of $\kappa_{\nu}\,{=}$ 1 cm$^2$g$^{-1}$ at $\lambda\,{>}\,1300\,\mu$m, $\kappa_{\nu}\,{=}\,\kappa_{1300\mu\rm m}\left(\frac{\lambda}{1300\mu\rm m}\right)^{-1}$ between $1300$ and $100\,\mu$m and again a constant value of $\kappa_{\nu}\,{=}\,\kappa(100\,\mu\rm m)$ at $\lambda\,{<}\,100\,\mu$m. We fitted the SED via $\chi^2$ minimisation using a genetic optimization algorithm PIKAIA \citep{charbonneau}. This algorithm performs the maximization of a user defined function, for which purpose we used the inverse $\chi^2$. Since there are many photometric measurements in the mid-infrared domain, but just a few in the far-infrared, the mid-infrared region has a higher weight during the fit, compared to the far-infrared domain. Therefore, in order to ensure an equally good fit at all wavelengths, we divided the SED into four regions and weighted the $\chi^2$ of each domain with the inverse of number of photmetric points the region contained. Then the final $\chi^2$ was the sum of the $\chi^2$ of all regions. The regions we used were: $0.3-3$, $3-30$, $30-300$ and $300-3000\,\mu$m. The parameters of the best-fit model, which gives a weighted $\chi^2$ of 0.67, are listed in Table~\ref{tab:par}. The fitted model SED as well as the disc and envelope components are overplotted in Fig.~\ref{fig:sed}. The model SED is consistent with the observed fluxes. This shows that the picture of a thin accretion disc and an envelope is consistent with both the measured SED and the geometry and disc/envelope parameters inferred from our polarimetric observations. Detailed modelling of the silicate, PAH, and ice spectral features, as well as the correct treatment of internal extinction and scattering would require radiative transfer modelling, which will be the topic of a subsequent paper. \subsection{The evolutionary status of \pars{}} \label{sec:general} The geometry of FUor models discussed in the literature (e.g.~\citealt{hk96, tbb}) usually consist of a central star surrounded by an accretion disc and an infalling envelope with a wind-driven polar hole. These assumptions are supported by the fact that they fit well the SED \citep{green, quanz}, the interferometric visibilities \citep{mg,v1647ori_midi} and the temporal evolution of the SED \citep{fuors}. In this paper we present the first direct imaging of these circumstellar structures in a FUor. Our polarimetric measurements of \pars{} show the existence of a circumstellar disc which extends from at least 48 to 360 AU. The most striking feature of the disc is its flatness over the whole observed range. The short-wavelength part of the SED could be well reproduced using a radial temperature profile of $r^{-0.75}$ (Sect.~\ref{sec:modelling}). This profile is expected from both a geometrically thin accretion disc and a flat reprocessing disc. An envelope was also seen in the polarization maps of \pars{} and it was also a necessary component for the SED modelling. Envelopes are involved in many FUor models and in this paper we present a direct detection of this model component. Our images reveal that the envelope can be followed inwards as close to the star as the disc. FUor models often assume a polar cavity in the envelope, created by a strong outflow or disc wind. The direct images of \pars{} clearly show the presence of such a cavity and we also detected a bipolar outflow in the \pars{} system. In the recent years, as new interferometric and infrared spectroscopic observations were published for FUors, the group turned out to be more inhomogeneous in physical properties than earlier assumed, when mainly optical photometry and spectroscopy had been available. \citet{quanz} proposed that some differences might be understood as an evolutionary sequence. They suggest that FUors constitute the link between embedded Class I objects and the more evolved Class II objects. Members of the group exhibiting silicate absorption at $10\,\mu$m are younger and more embedded (Category 1, e.g.~V346\,Nor); while objects with pure silicate emission are more evolved (Category 2, e.g.~FU\,Ori and Bran\,76). There are objects showing a superposition of silicate absorption and emission, which are probably in an intermediary evolutionary stage (e.g.~RNO\,1B). \citet{green} also sorted FUors, based on the ratio of the far-infrared excess and the luminosity of the central accretion disc, $f_d$ (Equ.~7 in their paper). A large relative excess ($f_d\,{>}\,5\%$) indicates an envelope of large covering fraction (V1057\,Cyg and V1515\,Cyg), while low relative excess means a tenuous or completely missing envelope (Bran\,76 and FU\,Ori). This is also an evolutionary sequence, as young, more embedded objects have large envelopes, while around more evolved stars, the envelope has already dispersed. $f_d$ can also be used to calculate the opening angle of the envelope, thus a prediction of this scheme is that the opening angle is becoming wider during the evolution, probably due to strong outflows during the repeated FUor outbursts. The two classification schemes are not inconsistent and one can merge them into the following evolutionary sequence: (1) the {\it youngest objects} exhibit silicate absorption and large far-infrared excess (V346\,Nor, probably also OO\,Ser and L1551\,IRS\,5 belong here); (2) {\it intermediate-aged objects}, where the silicate feature is already in emission but there is still a significant far-infrared excess (V1057\,Cyg, V1515\,Cyg, probably also RNO\,1B and V1647\,Ori); (3) the most {\it evolved objects} show pure silicate emission and low far-infrared excess (FU\,Ori, Bran 76). We note, however, that this classification has some weak points. As \citet{quanz} already mentioned, an edge-on geometry in a more evolved system may appear as a younger one. Moreover, during an outburst and the subsequent fading phase, certain spectral features as well as the global shape of the SED may change. \pars{} can be placed in this evolutionary scheme, though one should keep in mind that because of the nearly edge-on geometry, the classification of this object is somewhat uncertain. \pars{} displays silicate emission (Fig.~\ref{fig:pah}). Integrating the flux of the two components in our simple model (Sect.~\ref{sec:modelling}), and correcting the apparent disc luminosity for inclination effect ($i=86^{\circ}$, Table~\ref{tab:par}) using Equ.~6 of \citet{green}, we obtained a large relative far infrared excess of $f_d\,{=}\,75\%$. These two properties place \pars{} into the intermediate-aged category (though because of its inclination, it may actually seem younger than it is). Following Equ.~7 of \citet{green}, from the $f_d$ value, we also computed the opening angle of the envelope. The resulting opening angle of $60^{\circ}$ agrees well with the angle measured in the direct NACO images (Sect.~\ref{sec:env}). \pars{} was placed into the evolutionary scheme using two parameters: the silicate feature and the relative far infrared excess. In the following we discuss whether its other physical characteristics match with those of other FUors. \begin{itemize} \item[{\it (i)}] Our observations revealed that the circumstellar disc of \pars{} is very flat. Due to lack of similar direct measurements for other FUors, we can only speculate that perhaps all FUors with envelopes have such flat discs. On the other hand, the most evolved FUor, FU\,Ori, seems to have no envelope but its disc is probably flared \citep{kh91, green, quanz2}. This might suggest that disc flaring develops at later stages, when illumination from the central source may heat the disc surface more directly. \item[{\it (ii)}] In a nearly edge-on system like \pars{}, one expects to see the $10\,\mu$m silicate feature in absorption. The fact that \pars{} has silicate emission indicates that the line of sight towards the central region is not completely obscured. Using the optical depth of the $15.2\,\mu$m CO$_2$ ice feature, we calculated an $A_V\,{=}\,8\,$mag ($A_V\,{=}\, 38.7 A_{15.2 \mu\rm{}m}$, \citealt{savage}). This value is surprisingly low compared to V1057\,Cyg ($A_V\,{\sim}\,50-100\,$mag, \citealt{kh91}). This indicates a much more tenuous envelope, which is also supported by the low envelope mass of $0.02\,\rm{}M_{\odot}$ in our modelling. \item[{\it (iii)}] Following \citet{quanz}, we analysed the profile of the $15.2\,\mu$m CO$_2$ ice feature of \pars{}. The inset in Fig.~\ref{fig:pah} shows that the feature has a characteristic double-peaked sub-structure, very similar to HH\,46\,IRS, an embedded young source \citep{boogert}. HH\,46\,IRS is a reference case for processed ice. The presence of processed ice in \pars{} indicates heating processes and the segregation of CO$_2$ and H$_2$O ice, already at this evolutionary stage. Other FUors exhibiting this kind of profile are L1551\,IRS\,5, RNO\,1B and RNO\,1C \citep{quanz}. \end{itemize} The evolutionary state of a young stellar object can also be estimated following the method proposed by \citet{chen}. According to their Equ.~(1) we calculated a bolometric temperature of T$_{\rm bol}=410\,$K for the measured SED. We compared this value with the distribution of corresponding values among young stellar objects in the Taurus and $\rho\,$Ophiuchus star forming regions (Chen et al. 1995). From this check we can conclude that \pars{} seems to be a class I object, and its age is ${\sim}\,10^5\,$yr. However, \citet{green} argued that the apparent SED of the disc component depends on the inclination. Thus we computed T$_{\rm bol}$ also for a face-on disc configuration and obtained T$_{\rm bol}=1160\,$K, corresponding to a Class II object. In fact, \pars{} is probably close to the Class I / Class II border, in accordance with the proposal of \citet{quanz}. | 7 | 10 | 0710.1431 |
0710 | 0710.3328_arXiv.txt | We analyze the electromotive force (EMF) terms and basic assumptions of the linear and nonlinear dynamo theories in our three-dimensional (3D) numerical model of the Parker instability with cosmic rays and shear in a galactic disk. We also apply the well known prescriptions of the EMF obtained by the nonlinear dynamo theory (Blackman \& Field 2002 and Kleeorin et al. 2003) to check if the EMF reconstructed from their prescriptions corresponds to the EMF obtained directly from our numerical models. We show that our modeled EMF is fully nonlinear and it is not possible to apply any of the considered nonlinear dynamo approximations due to the fact that the conditions for the scale separation are not fulfilled. | \label{sec:intro} It seems that the issue of the magnetic field amplification in galaxies may be well explained by the two main physical mechanisms: the Parker instability (PI), which takes into account the cosmic rays (CR) and the shear \cite[e.g.][and references therein]{hanasz03,hanasz04}, and the magneto-rotational instability \cite[MRI, e.g.][]{dziourkevitch04,kitchatinov04}. The possible scenario of the magnetic field evolution could be presented as follows: when the protogalaxy starts to rotate differentially, the MRI mechanism occurs, and this results in a very efficient magnetic field amplification even to the level of $\mu$G with the global e-folding time of 100~Myr or even less \citep{dziourkevitch04}. Simultaneously, the quadrupole symmetry of the large-scale magnetic field is being created. MRI also causes the turbulent motions in the galactic disks. The process of supernovae (SN) explosions, that arises in young galactic objects \cite[e.g.][and references therein]{widrow02} suppresses the MRI mechanism. In the same time, the cosmic rays produced in SN remnants may induce the Parker instability process \cite[e.g.][]{parker92,hanasz04}. Hence, we may conclude that during the early stage of galaxy evolution the MRI process is being replaced by the PI mechanism. The local simulations of the large-scale magnetic field took into account the turbulent dynamo theory and the magnetic back-reaction onto the turbulent motions \citep{piddington70,piddington72a,piddington75a,kulsrud95}. The authors drew the conclusion that it was difficult to obtain the amplification of the total magnetic field \cite[e.g.][and references therein]{widrow02}. It might be explained by the equipartition of the random magnetic field component with the random turbulent motions. Such process suppresses the dynamo action at later times. Moreover, the magnetic field amplification might be easily stopped even by the presence of the weak large-scale magnetic field \citep{cattaneo91,vainshtein92,cattaneo94,cattaneo96,ziegler96}. \cite{blackman99} explained that papers analyzing the analytically strong suppression of the dynamo coefficient $\alpha$ should distinguish between different state orders of the turbulent quantities. However, their analysis did not completely solve the problem of the quenching of the dynamo coefficients. In their next paper \citep{blackman00}, they proved that the results from the \cite{cattaneo96} model were based on the assumption about the periodicity of the boundary conditions. Nevertheless, the quenching effects could also appear even when the open boundary conditions were applied \cite[e.g.][]{brandenburg01b}. The classical dynamo theory does not conserve the total magnetic helicity \cite[see e.g.][BF02]{blackman02}. In media characterized by the high magnetic Reynolds number ($R_m\gg1$) the total helicity should remain constant in closed regions \cite[e.g.][]{berger84,brandenburg02b,brandenburg05b,subramanian02}. The permanent helicity is also an additional factor, which suppresses the dynamo activity \citep[e.g.][]{brandenburg02b}. The following papers \citep{blackman00,kleeorin00,blackman03,kleeorin99,kleeorin02,kleeorin00, kleeorin03,rogachevskii00,rogachevskii01} presented the two methods that allowed the modeling of the dynamo action evading the problem of the constant helicity. The first method is based on the ejection of the magnetic helicity through boundaries. The second one uses the creation of the negative and positive helicity at the large and small scales respectively \cite[see][]{brandenburg02b,kleeorin02,kleeorin03}. \cite{blackman02} in their next paper analyzed the nonlinear prescription of both dynamo coefficient $\alpha$ and $\beta$. Both factors were obtained without any linearization and took into account all terms in the equation of the evolution of the fluctuating part of the magnetic field (see Eq.~\ref{eqn:fluct_field_evol}). The results were similar to the dynamo coefficients obtained by \citep{pouquet76}, but the units were different (without the time integration). They also solved the small- and large-scale helicity dynamo equations numerically simultaneously with the equation for the EMF time evolution. That allowed them to obtain the growth of the large-scale magnetic field in the kinematic phase. The new form of the dynamo coefficients for anisotropic turbulent motions with the presence of the large-scale magnetic field was presented by \cite{rogachevskii01}. They calculated dynamo coefficients according to the \cite{raedler80} EMF prescription \cite[see also][]{kowal05}, which neglected all quadratic terms in the mean field in the EMF. \cite{kleeorin03} used those forms of the dynamo coefficients to solve numerically the dynamo equation in the local thin-disc approximation. The authors took into account the quenching of both coefficients, $\alpha$ and $\beta$, and helicity flux through the boundary. They found that it was possible to obtain the growth of the large-scale magnetic field when $\alpha$-quenching was only analyzed \citep{kleeorin02}. If the model included also $\beta$-quenching, no growing solution of the dynamo \citep{kleeorin03} could be obtained. Both in Blackman-Field and in Kleeorin-Rogachevskii approaches there is an $\alpha_{\rm m}$ term that quantifies the small scale helicity current. This term depends on the current helicity flux, and there are different theories for this flux. In Kleeorin et al. a heuristically motivated expression for the flux was used, in \cite{brandenburg05b} the \cite{vishniac01} flux was used, and in \cite{shukurov06} a simple advective flux was used. In all these cases the current helicity flux allows the field to saturate at high levels. The latest research on the turbulent enforcement in the solar convective zone calculated in the local cube with the shear has shown that the open boundaries help to obtain the amplification of the large-scale magnetic field even without the helicity of turbulent motions \cite[end references therein]{brandenburg05a,brandenburg05b}. We have to stress that their result, an increase of the total magnetic energy, was obtained without any assumption considering the additional EMF of dynamo. The authors applied isotropic and homogeneous turbulence with and without the helical forcing in their model. On the other hand, Brandenburg and his collaborators \citep{brandenburg05a} interpreted their results in terms of the mean field dynamo theory. The time evolution of $\alpha$ was obtained from the calculated electromotive force \citep{brandenburg04,kowal05,brandenburg02a,brandenburg01a}. \cite{brandenburg04} explained that thanks to the flux of the current helicity flowing out of the cube the process of $\alpha$-quenching tends to be not as disastrous as forseen. The only conditions are the intermediate level of $R_m$ and open boundaries. When the high value of $R_m$ is applied to the model, the comparatively lower value of the large-scale magnetic field strength are obtained \citep{brandenburg04}. However, in astrophysical objects $R_m$ is always high. That is why this result seems to be peculiar. On the other hand, the authors made it clear that the total magnetic energy grows mainly in the kinematic phase of the dynamo and their results do not depend strongly on $R_m$. Futhermore, they calculated the value of $\alpha$ coefficient based on the modeled EMF. The obtained factor was similar to the same coefficient calculated according to the \cite{kleeorin00,kleeorin02,kleeorin03} prescriptions. We believe that the fact that such similarity occurs results from the isotropic and homogeneous turbulence in both models \citep{brandenburg05a,brandenburg05b,brandenburg04}. It may also happen due to the fact that the authors applied standard dynamo approximations, which neglect the quadratic terms in the mean field in EMF. The previously mentioned results indicate that the realistic physical simulations are of the great importance when the MRI \citep{dziourkevitch04} and the PI \citep{hanasz04,hanasz03,kowal05} processes are considered. Both models, which meet enumerated requirements, showed that it was possible to amplify galactic magnetic field efficiently. \cite{hanasz03} and \cite{hanasz04,hanasz05} presented that the following two processes: the Parker instability driven by cosmic rays from supernovae and the shear from the differential rotation enable the magnetic field amplification (with the e-folding time scale of 250~Myr or even 140~Myr). The model also applied realistic gravity according to the \cite{ferriere98} prescriptions. The idea of obtaining the dynamo coefficients ($\alpha$ and $\beta$) from the electromotive force calculated in the local numerical simulations proved to be essential for many other authors too \cite[e.g.][etc]{ziegler96,brandenburg02a}. We included the calculations of the dynamo coefficients, which we obtained from the calculated EMF in our previous paper \citep{kowal05}. The application of the statistical methods provided us with the acceptable values of the dynamo $\alpha$-tensor. On the other hand, the values of $\beta$ coefficients were negative. Such values are inconsistent with the R\"adler prescription \citep{raedler80}. This may be caused either by the applied statistical method, which does not take into account physical differentiation, or by the linear EMF approximation \citep{kowal05}. For this reason we decided to analyze that matter in our present study. We search for the conditions, which should be fullfiled in order to make linearization of the electromotive force in the mean field dynamo theory possible. We would like to examine the following problems: the scale separation, the ratios of the terms in the equation for the small-scale magnetic field evolution \cite[e.g.][]{raedler80}, the magnitude of the turbulent kinetic energy in comparison to the large-scale magnetic one. Next, we plan to apply the estimations of the dynamo coefficients derived by \cite{blackman00,rogachevskii01} and \cite{brandenburg04} to our models. We would like to check if their approximations fit into the calculated electromotive force in our models \cite{hanasz04}. In this part of this work we do not include the explicit analysis involving the conservation of the magnetic helicity. The investigation of the magnetic helicity conservation in our models is already advanced, but it is complex enough to be described in separate paper, which is under preparation. We cannot apply the approximations of \cite{ruediger93} and \cite{kitchatinov94} derived from the EMF quenching by the usage of the Second Order Correlation Approximations. They assumed that the Strouhal number $S$ is essentially smaller than 1 ($S\ll1$, where $S = \tau_c \times \rm v/ \rm l_c$). This assumption is not fullfiled in our numerical experiments, where $S$ is about 1. Finally, we discuss our results. \begin{table*} \begin{center} \caption{Parameters of models examined in this paper. (*) The conversion rate, presented as 10\% in \cite{hanasz04}, was in fact equal to 100\%, due to a trivial calculation mistake. Therefore, the overall injection rate of cosmic ray energy was equivalent to a realistic one, corresponding to SN rate=~20~kpc$^{-2}$Myr$^{-1}$, with the energy conversion factor = 10\%. \label{table-params}} \begin{tabular}{|l|c|c|c|c|} \hline Model & A & B & C & D\\ \hline Domain sizes [kpc] & 0.5 $\times$ 1 $\times$ 1.2 & \multicolumn{3}{c}{0.5 $\times$ 1 $\times$ 4} \vline \\ Resolution & 50 $\times$ 100 $\times$ 120& \multicolumn{3}{c}{50 $\times$ 100 $\times$ 400} \vline \\ Vertical gravity at $R$ [kpc] = & 8.5 & \multicolumn{3}{c}{5} \vline\\ Gas column density [cm$^{-2}$] & ... & \multicolumn{3}{c}{27$\times$10$^{20}$} \vline\\ Angular velocity $\Omega$ [Myr$^{-1}$] & 0.05 & \multicolumn{3}{c}{0.05} \vline \\ SN rate [kpc$^{-2}$Myr$^{-1}$] & $2$ & \multicolumn{3}{c}{130} \vline \\ Initial $\alpha=e_{\rm mag}/e_{\rm gas}$ & $10^{-8}$ & \multicolumn{3}{c}{$10^{-4}$}\vline \\ Diffusion coefficients $K_\parallel$, $K_\perp$ [cm$^2$s$^{-1}$] & $3 \times 10^{27}$, $3 \times 10^{26}$ & \multicolumn{3}{c}{$3 \times 10^{27}$, $3 \times 10^{26}$}\vline \\ Conversion rate of SN kinetic to CR energy & 10\%$^{(*)}$ & \multicolumn{3}{c}{10\%}\vline \\ \hline Resistivity $\eta$ [cm$^2$s$^{-1}$] & $3\times 10^{24}$ & $0\times 10^{24}$ & $3\times 10^{24}$ & $30\times 10^{24}$ \\ \hline \end{tabular} \end{center} \end{table*} | \begin{enumerate} \item {Neither the velocity nor the magnetic field scale separation occurs in our model.} \item { The electromotive forces in the cosmic-ray driven dynamo model are nonlinear, but none of the two examined nonlinear approaches is capable of reproducing electromotive forces in the numerical experiments correctly.} \item {Various nonlinear prescriptions of the dynamo coefficients have been proposed by other outhors, however, they are not capable of reconstructing the electromotive force resulting from experiments of cosmic-ray driven dynamo. Moreover the reconstructions of the magnetic field produce too fast or too slow growth of the magnetic energy in comparison with the results of cosmic-ray driven dynamo numerical experiments.} \end{enumerate} Extension of the present work, including considerations of magnetic helicity conservation will be presented in the forthcoming paper. | 7 | 10 | 0710.3328 |
0710 | 0710.1094_arXiv.txt | We report a measurement of the supernova (SN) rates (Ia and core-collapse) in galaxy clusters based on the 136 SNe of the sample described in \citet{C99} and \citet{M05}. Early-type cluster galaxies show a type Ia SN rate (0.066 SNuM) similar to that obtained by \citet{sharon07} and more than 3 times larger than that in field early-type galaxies (0.019 SNuM). This difference has a 98\% statistical confidence level. We examine many possible observational biases which could affect the rate determination, and conclude that none of them is likely to significantly alter the results. We investigate how the rate is related to several properties of the parent galaxies, and find that cluster membership, morphology and radio power all affect the SN rate, while galaxy mass has no measurable effect. The increased rate may be due to galaxy interactions in clusters, inducing either the formation of young stars or a different evolution of the progenitor binary systems. We present the first measurement of the core-collapse SN rate in cluster late-type galaxies, which turns out to be comparable to the rate in field galaxies. This suggests that no large systematic difference in the initial mass function exists between the two environments. | \label{sec:intro} Type Ia Supernovae (SNe Ia) are believed to be the result of the thermonuclear explosion of a C/O white dwarf (WD) in a binary system due to mass exchange with the secondary star. This conclusion follows from a few fundamental arguments: the explosion requires a degenerate system, such as a white dwarf; the presence of SNe Ia in old stellar systems implies that at least some of their progenitors must come from old, low-mass stars; the lack of hydrogen in the SN spectra requires that the progenitor has lost its outer envelope; and, the released energy per unit mass is of the order of the energy output of the thermonuclear conversion of carbon or oxygen into iron. Considerable uncertainties about the explosion model remain within this broad framework, such as the structure and the composition of the exploding WD (He, C/O, or O/Ne), the mass at explosion (at, below, or above the Chandrasekhar mass) and the flame propagation (detonation, deflagration, or a combination of the two). The key observations constraining the explosion models are the light curve and the evolution of the spectra. Large uncertainties also remain regarding the nature of the progenitor binary system, its evolution through one or more common envelope phases, and its configuration (single or double-degenerate) at the moment of the explosion (see \citealt{yungelson05}, for a review). Solving the problem of the progenitor system is of great importance for modern cosmology as SNe dominate metal production, (e.g., \citealt{matteucci86}), are expected to be important producer of high-redshift dust \citep{maiolino01, maiolino04a,maiolino04b,bianchi07}, and are essential to understand the feedback process during galaxy formation (e.g., \citealt{scannapieco06}). The nature of the progenitor systems can be probed by studying the SN rate in different stellar populations, and constraining the delay time distribution (DTD) between star formation and SN explosion. \smallskip In 1983, Greggio \& Renzini computed the expected DTD for a single-degenerate system. The computation was later refined by many authors and extended to double-degenerate systems \citep{tornambe86,tornambe89,tutukov94,yungelson00,matteucci01,belczynski05,greggio05}. The DTD can be convolved with the star formation history (SFH) of each galaxy to obtain its SN rate. The observation of the SN rates per unit mass in galaxies of different types \citep{M05,sullivan06} and in radio-loud early-type galaxies \citep{dellavalle05} has proved to be an effective way to constrain the DTD. The SN rates per unit mass show that SNe Ia must come from both young and old progenitors \citep{M05,sullivan06}. The dependence of the SN rate on the radio power of the parent galaxy is well reproduced by a ``two channel'' model \citep{mannucci06}, in which about half of the SNe Ia, the so-called ``prompt'' population, explode soon after the formation of the progenitors, on time scales shorter than $10^8$ yr, while the other half (the ``tardy'' population) explode on a much longer time scale, of the order of $10^{10}$ yr. Several attempts to compare the evolution of SN rate with redshift with that of the SFR have also been presented (see, among many others, \citealt{galyam04,dahlen04,cappellaro05,neill06,barris06,botticella07} and \citealt{poznanski07}), but the large uncertainties on both quantities prevent strong conclusions (see, for example, \citealt{forster06}). \smallskip In principle, an accurate measurement of the DTD could identify the progenitor binary system. In practice, both the large number of free parameters involved in the theoretical computations of the DTD, and the complex SFHs of most of the galaxies make this identification much more uncertain. To solve the problem of the complexity of the SFH, it is interesting to measure the SN Ia rate in galaxy clusters. Most of the stellar mass of these systems is contained in elliptical galaxies, whose stellar populations are dominated by old stars. Despite the problem that even a small amount of new stars could give a significant contribution to the SN rate (see the discussion in sect.~\ref{sec:discussion}), the reduction in the uncertainty in the SFH is of great help to derive the DTD. \smallskip The cluster SN rate is also of great importance to study the metallicity evolution of the universe. The gravitational potential well of galaxy clusters is deep enough to retain in the intracluster medium (ICM) all the metals which are produced in galactic or intergalactic SNe. As a result, the metallicity of the ICM is a good measure of the integrated past history of cluster star formation and metal production. As discussed by \citet{renzini93}, the measured amount of iron is an order of magnitude too high to be produced by SNe Ia exploding at the current rate. Explanations of this effect include the presence of higher SN rates in the past \citep{matteucci06}, the importance of the intracluster stellar population \citep{zaritsky04}, or evolving properties of star formation processes \citep{maoz04,lowenstein06}. The observed abundance ratios in the ICM can be used to constrain the ratio between the total numbers of Ia and CC SNe, as recently done by \citet{deplaa07}. Constraints on the SN Ia models can also be derived from the radial distribution of metallicity \citep{dupke02}. \citet{calura07} used the observed cosmic evolution of iron abundances in \citet{balestra07} to constrain the history of SN explosion, iron formation and gas stripping in galaxy clusters. They found good agreement with the observations, especially when the ``two channel'' model of SNe Ia by \citet{mannucci06} is used. \smallskip There are strong motivations for measuring also the cluster rates of the other physical class of SNe, the core-collapse (CC) group. Type II and type Ib/c SNe are attributed to this group because there is a general consensus that these explosions are due to the collapse of the core of a massive (about 8--40~\msun) star. Thus, CC SNe are expected to be good tracers of star formation in moderately dusty environments (see \citealt{mannucci07}). Their rate per unit mass is also very sensitive to the initial mass function (IMF), because SN explosions are due to massive stars while most of the mass is locked in low-mass stars. As a consequence, studying the CC SN rate as a function of environment is a sensitive test for any systematic difference in IMF. \begin{table} \caption{ Measured type Ia SN rates in early-type cluster galaxies \label{tab:history} } \begin{tabular}{llcc} \hline \hline Reference & ~~$z$ ~~($z$ range) & N$_{SN}$& Rate\\ & & & (SNuB) \\ \hline This work & 0.02 (0.005--0.04) & 20 & $0.28^{+0.11}_{-0.08}$\\ \citet{crane77} & 0.023 (0.020--0.026)& 8 & $\sim$0.10 \\ \citet{barbon78} & 0.023 (0.020--0.026)& 5 & $\sim$0.16 \\ \citet{germany04} & 0.05 (0.02--0.08) & 23 & unpubl. \\ \citet{sharon07} & 0.15 (0.06--0.19) & 6 & 0.27$^{+0.16}_{-0.11}$\\ \citet{galyam02} & 0.25 (0.18--0.37) & 1 & 0.39$^{+1.65}_{-0.37}$\\ \citet{galyam02} & 0.90 (0.83--1.27) & 1 & 0.80$^{+0.92}_{-0.41}$\\ \hline \end{tabular} \end{table} \subsection{The observed cluster supernova rate} Prompted by all these motivations, several groups have measured the SN Ia rate in galaxy clusters, but the results are still quite sparse. The first published values are due to \citet{crane77} and \citet{barbon78} (see Table~\ref{tab:history} for a summary, including the results of our work, discussed below), before a clear distinction between type Ia and Ib/c had been introduced. They used a sample of 5--8 SNe in the Coma Cluster and constrained the SN rate to be of the order of 0.15 SNuB (SN per century per $10^{10}$ \lsun\ in the B band). The SN rate as a function of galaxy environment was also addressed by \citet{caldwell81} to derive information on SN progenitors. Modern searches for cluster SNe begin with \citet{norgaard89} who discovered a SN Ia in a cluster at $z=0.31$. Starting from the late '90s, the Mount Stromlo 1.3 m telescope was used to monitor a few tens of Abell Clusters \citep{reiss98}. Three years of monitoring resulted in the detection of 23 candidate SNe Ia in cluster galaxies \citep{germany04}, but a rate based on this sample was never published. The first rates for cluster galaxies based on modern searches were published by \citet{galyam02}. These authors used archive images from the Hubble Space Telescope (HST) of 9 galaxy clusters, and discovered 6 SNe, 2 of which are associated with the clusters, at $z=0.18$ and $z=0.83$. The derived rates were affected by large statistical uncertainties due to the small number of detected SNe, but were consistent with a moderate increase of the rate with redshift compared to the rate in local elliptical galaxies. A sample of 140 low-redshift Abell clusters were monitored by the Wise Observatory Optical Transient Search (WOOTS, \citealt{galyam07}) using the Wise 1m telescope. The seven detected cluster SNe were used to constrain the fraction of intergalactic stars and SNe \citep{galyam03} and to measure the cluster SN rate \citep{sharon07}. This latter work obtains a value of the SN rate per unit mass of $0.098^{+0.058}_{-0.039}$ SNuM (SN per century per $10^{10}$ \msun\ of stellar mass), which is larger than, but still consistent with, the value of $0.038^{+0.014}_{-0.012}$ SNuM, derived by \citet{M05} for local ellipticals. Finally, a SN search in clusters is ongoing at the Bok Telescope on Kitt Peak \citep{sand07}. All of the previous published SN Ia rates are based on a small number of SNe and, as a consequence, have large statistical errors. Also, a cluster rate for CC SNe has never been published because many of the cited samples only contain Ia SNe. In this work, we use the SNe in the \citet{C99} sample to study the SN rate as a function of galaxy environment. Throughout this paper we use the ``737'' values of the cosmological parameters: $(h_{100},\Omega_m,\Omega_\Lambda)=(0.7,0.3,0.7)$. \begin{figure} \includegraphics[width=9cm]{galdist.ps} \caption{ \label{fig:galdist} Surface density of galaxies (upper panel) and fraction of early-type galaxies (lower panel) as a function of the projected distance from the closest cluster. Above 3 Mpc, the average for field galaxies is shown. The vertical dotted lines show the two projected distances, 0.5 and 1.5 Mpc, used to define cluster galaxies (see text). } \end{figure} | \label{sec:discussion} The interpretation of the possible difference in SN Ia rate between cluster and field early-type galaxies is not straightforward. As the observed rate is the convolution of the SFH with the DTD, the differences could be due to either of these functions. \begin{enumerate} \item The first possibility is that the rate difference is due to differences in the stellar populations. \citet{M05}, \citet{sullivan06}, and \citet{aubourg07} have shown that the type Ia SN rate has a strong dependence on the parent stellar population, with younger stars producing more SNe. The difference in SN rate could be related to this effect, i.e., to a higher level of recent star formation in cluster ellipticals. Only a very small amount of younger stars is needed, because the amplitude of the DTD at short times can be hundreds of times larger than at long times. As an example, the \citet{greggio83} single-degenerate model has 300 times more amplitude at $10^8$ yr than at $10^{10}$ yr, and this means that a recently formed stellar population contributing 0.3\% of the mass can provide as many SNe as the remaining 99.7\% of old stars. For the ``two channel'' model by \citet{mannucci06}, the amount of young stars needed can be even lower, at the 0.1\% level, as this DTD amplitude ratio between $10^7$ and $10^{10}$ years is as large as 1000. \smallskip The presence of traces of star formation in early-type galaxies is not inconsistent with other observations. Many ellipticals show signs of recent interactions or star formation activity: faint emission lines \citep{sarzi06}, tidal tails \citep{vandokkum05}, dust lanes \citep{vandokkum95,colbert01}, HI gas \citep{morganti06}, molecular gas \citep{welch03}, and very blue UV colors \citep{kaviraj06,schawinski07,haines07,kaviraj07}. Even if the interpretation of most of these effects is matter of debate (for example, \citealt{serego07} have found only small amounts of HI gas in cluster ellipticals), the observations suggest a widespread, low-level presence of star formation. The dependence of this presence with environment is not settled yet. \citet{ferreras06} have found evidence for recent star formation, at the percent level, in ellipticals in compact groups, but not in field ellipticals. In contrast, \citet{verdugo07} and \citet{haines07} have found higher levels of present star formation in field rather then cluster early-type galaxies. \smallskip Some studies (see, for example, \citealt{sanchez06}, \citealt{bernardi06} and \citealt{collobert06}), have found younger ages in field early-type galaxies with respect to cluster galaxies (but \citealt{serego06} have found no difference). Taken at face value, this would seem to contradict the star formation interpretation of the SN rate, but this is not necessary the case. Field ellipticals could be younger that cluster ellipticals, but nevertheless they could show a lower level of {\em present} star formation. The difference in the age of the {\em dominant} stellar population of early-type galaxies, of the order of 1 Gyr for ages of about 12 Gyr, might not be directly related to the amount of star formation in the last few $10^8$ years. Such a contribution cannot be detected in the integrated colors of the galaxies. The expected differences are at the 0.05 mag level for the (B--K) color, assuming the younger stars are not associated with dust, and even smaller (0.02 mag for $A_V$=1), allowing for dust extinction. It is usually assumed that early-type galaxies can form new stars only after merging with a small, gas rich galaxy, because usually they do not host much interstellar gas. The average amount of stars formed is proportional to the merger (or encounter) rate, to the typical amount of gas in the accreted galaxy, and to the efficiency of star formation in the accreted gas. It is possible that one or more of these quantities are larger for cluster galaxies than for field galaxies because of the different galaxy volume density and galaxy-galaxy encounter velocity. \smallskip If this is the correct interpretation, the ``prompt'' population of SNe Ia would be associated with the explosion of CC SNe from the same young stellar populations. If a SN Ia is to explode within $10^8$ yr of the formation of its progenitor, the primary star of the progenitor binary system must have a mass above 5.5 \msun\ to allow for the formation of a white dwarf in such a short time. \citet{mannucci06} have shown that reproducing the observed SN rates by using the ``bimodal'' DTD in that paper implies that about 7\% of all stars between 5.5 and 8 \msun\ explode as ``prompt'' SNe Ia, while the ``tardy'' population corresponds to a lower explosion efficiency, about 2\%, and on a much longer timescale (see also \citealt{maoz07} for various estimates of these efficiencies). For a Salpeter IMF and assuming that 100\% of the stars between 8 and 40 \msun\ end up at CC SNe, we expect 1.3 CC SNe for each ``prompt'' type Ia. Assuming that the difference between cluster and field early-type galaxies is due to the ``prompt'' SNe Ia, the rate of this population is of the order of 0.066-0.019=0.047 SNuM (see Table~\ref{tab:massrate}). Converting this rate to an observed number, about 2 CC SNe are expected in the cluster early-type galaxies of our sample, consistent with our null detection at about 1.3$\sigma$ level. We conclude that the non detection of CC in the early-type galaxies belonging to our sample and the corresponding upper limits to the CC rate are consistent with the hypothesis of a ``prompt'' Ia component. We also note that some CC SNe have been discovered in the recent past in prototypical early-type galaxies. \citep{pastorello07}. \item A second possible interpretation is that the higher rate in cluster early-type galaxies is related to differences in the DTD. If the stars in ellipticals are 9-12 Gyrs old (see, for example, \citealt{mannucci01}), the SN rate is dominated by the tail of the DTD at long times. Differences in the environments could produce small differences in the shape of this function, for example because of the higher numbers of encounters. A interesting possibility is also that the changes in the DTD are related to differences in metallicity between cluster and field early-type galaxies, as discussed by \cite{sanchez06,bernardi06,collobert06} and \cite{prieto07}. The differences between cluster and field galaxy metallicity presented by these papers are neither large nor always in the same direction. Nevertheless systematic, although not large, differences in metallicity could be present and produce significant changes in the DTD, for example, by affecting the efficiency of mass loss during the complex life of a binary system. \end{enumerate} Table~\ref{tab:history} lists the different measurements of the SN rate in early-type cluster galaxies. The evolution of this rate can be compared with the history of star formation of the parent galaxies to derive the DTD. Currently published cluster SN rates at z$>$0.2 are too uncertain to permit any strong conclusions. However, current and future searches for SNe are expected to change this situation and allow for the derivation of meaningful constraints (see, for example, \citealt{sharon06}). \smallskip To summarise, we have used a sample of 136 SNe in the local universe to measure the SN rate as a function of environment. For the first time, we measure the CC SN rate in clusters. We find it is very similar to the CC SN rate in field galaxies, suggesting that the IMF is not a strong function of the environment. For Ia SNe, the rates in clusters and in the field are similar for all galaxy types except for the early-type systems, where we detect a significant excess in clusters. This excess is not related to other properties of those galaxies, such as mass, morphology, or radio loudness. Environments itself appears to be important. We interpret this effect as possibly due to galaxy-galaxy interaction in clusters, either producing a small amount of young stars (of the order of the percent in mass over one Hubble time), or affecting the evolution of the properties of the binary systems. \bigskip {\bf Acknowledgments} We thank Sperello di Serego and the MEGA group (Arcetri Extragalactic Meeting) for useful discussions about the properties of elliptical galaxies. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. DM, MD, and AG thank the Kavli Institute for Theoretical Physics for its hospitality. This research was supported in part by the National Science Foundation under Grant No. PHY05-51164. | 7 | 10 | 0710.1094 |
0710 | 0710.1788_arXiv.txt | {The {\it ROSAT} X-ray source \cal\ has recently been identified as a likely compact object whose properties suggest it could be a very nearby radio millisecond pulsar at $d = 80 - 260$\,pc.} {We investigated this hypothesis by searching for radio pulsations using the Westerbork Synthesis Radio Telescope.} {We observed \cal\ at 385 and 1380\,MHz, recording at high time and frequency resolution in order to maintain sensitivity to millisecond pulsations. These data were searched both for dispersed single pulses and using Fourier techniques sensitive to constant and orbitally modulated periodicities.} {No radio pulsations were detected in these observations, resulting in pulsed radio luminosity limits of $L_{400}^{\rm max} \approx 0.3 (d/250 {\rm pc})^2$\,mJy kpc$^2$ and $L_{1400}^{\rm max} \approx 0.03 (d/250 {\rm pc})^2$\,mJy kpc$^2$ at 400 and 1400\,MHz respectively.} {The lack of detectable radio pulsations from \cal\ brings into question its identification as a nearby radio pulsar, though, because the pulsar could be beamed away from us, this hypothesis cannot be strictly ruled out.} | Recently, \citet*{rfs07}, hereafter RFS07, have identified the X-ray source \cal\ (from the {\it ROSAT} All-Sky Survey Bright Source Catalog) as having an X-ray to optical flux ratio $F_{\rm X}(0.1-2.4 {\rm keV})/F_{\rm V} > 8700$ ($3\sigma$). Such a high ratio is strong evidence that \cal, dubbed ``Calvera'' by RFS07, is a compact object. RFS07 consider several specific source classes to explain \cal's X-ray spectrum and luminosity: an X-ray dim isolated neutron star (INS), an anomalous X-ray pulsar (AXP), a compact central object (CCO), or a nearby radio pulsar. Based on careful comparison of \cal's properties with the canonical features displayed by these different classes of neutron star, RFS07 conclude that the most likely explanation is that \cal\ is a very nearby radio pulsar ($d = 80-260$\,pc) similar to the radio millisecond pulsars (MSPs) residing in the globular cluster 47~Tuc. We have investigated this interpretation by conducting sensitive searches for radio pulsations using the Westerbork Synthesis Radio Telescope (WSRT) in the Netherlands. No pulsations were found by these searches, and we use these non-detections to place strong limits on the pulsed radio luminosity of \cal. These limits bring into question the interpretation of this object as a nearby radio pulsar. | No plausible astronomical radio pulsations or bright, dispersed single pulses were detected in any of our observations. Using the radiometer equation modified for pulsar signals \citep{dtws85}, we can place limits on \cal's flux density in these observations (Table~\ref{obs.tab}). We find that \cal\ has a maximum flux density at 400\,MHz $S_{400}^{\rm max} \approx 4$\,mJy and a maximum flux density at 1400\,MHz $S_{1400}^{\rm max} \approx 0.3$\,mJy, where the fractional uncertainty on these limits is roughly 50\%. RFS07 argue that \cal\ may be a radio pulsar with a nearby distance $d = 80-260$\,pc. Assuming the pulsar is isolated, we were sensitive to spin periods encompassing the observed range for radio pulsars ($P_{\rm spin} \sim 1$\,ms$-10$\,s), with a factor of roughly $2-5$ degredation in sensitivity due to red-noise at the longest periods. Using an assumed distance $d = 250$\,pc, we can convert our flux density limits to pseudo luminosity ($L \equiv Sd^2$) limits. We find $L_{400}^{\rm max} \approx 0.3 (d/250 {\rm pc})^2$\,mJy kpc$^2$ and $L_{1400}^{\rm max} \approx 0.03 (d/250 {\rm pc})^2$\,mJy kpc$^2$. A simple check of the ATNF pulsar catalog\footnote{Available at http://www.atnf.csiro.au/research/pulsar/psrcat} \citep{mhth05} shows that $\lesssim 1$\% of the known pulsars have a luminosity below these limits. This suggests that if \cal\ is a radio pulsar, then it is either especially weak, not beamed towards the Earth, or significantly further away than 250\,pc. Assuming \cal\ is a radio MSP, what is the probability that its radio beam will pass the Earth (i.e. what is the beaming fraction of such pulsars)? A period-dependent beaming fraction proportional to $P_{\rm spin}^{-0.5}$ has been shown for normal, un-recycled pulsars \citep[see e.g.][]{ran93}. When extrapolated to millisecond spin periods, this relation implies a $\sim 100$\% beaming fraction for MSPs. However, \citet{kxl+98} find observational evidence that this relationship does not apply directly to MSPs and that the beaming fraction of MSPs is more like $50-90$\%. These estimates are consistent with considerations of the MSP population in 47~Tuc. \citet{hge+05} used a deep {\it Chandra} observation of the cluster to perform a population analysis which concluded that there are likely $\sim 25$ radio MSPs ($< 60$ at 95\% confidence) residing in the cluster, independent of beaming. Comparing this with optical and radio studies of the MSP population by \citet{egh+03} and \citet{mdca04} respectively, who both conclude that the total MSP population in 47~Tuc is $\sim 30$, \citet{hge+05} conclude that the beaming fraction is $\gtrsim 37$\%. While these various studies suggest that the beaming fraction of MSPs is significantly larger than for normal pulsars, and could be quite high, we cannot strongly rule out beaming as the cause of our non-detection of \cal. It is also possible, though we feel unlikely, that we have not seen \cal\ because it is in a compact binary orbit and/or is highly accelerated by a binary companion. Our luminosity limits are for a coherent search and do not include these effects. However, given that we have employed acceleration searches, which are generally sensitive in cases where the total integration time $T_{\rm int}$ is less than roughly 1/10 of the orbital period $P_{\rm orb}$, \cal\ would have to be in a $\lesssim 2$-hr orbit {\it and} fairly weak (or eclipsed) not to have been detected in our searches of 15-min data sections. The shortest orbital period of any known radio MSP is 1.6\,hr \citep[][PSR~J0024$-$7204R in the GC 47~Tuc]{clf+00}. In conclusion, given the caveats we have discussed, primarily pertaining to extreme orbital or spin parameters, we believe that these searches were sensitive enough to detect any nearby radio pulsar coincident with \cal, assuming it is beamed towards the Earth. We have checked for catalogued radio sources in the FIRST, NVSS, and WENSS surveys and find no counterpart to \cal. Future, deeper observations detecting a radio, optical, or X-ray pulsar wind nebula could address whether this source is a nearby radio pulsar beamed away from Earth. Alternate scenarios for \cal's nature, for instance that it might be the first unhosted CCO (RFS07), should continue to be investigated. | 7 | 10 | 0710.1788 |
0710 | 0710.0351_arXiv.txt | {Modern observations and models of various astrophysical objects suggest that many of their physical parameters fluctuate substantially at different spatial scales. The rich variety of the emission processes, including Transition Radiation but not limited to it, arising in such turbulent media constitutes the scope of Stochastic Theory of Radiation. We review general approaches applied in the stochastic theory of radiation and specific methods used to calculate the transition radiation produced by fast particles in the magnetized randomly inhomogeneous plasma. The importance of the theory of transition radiation for astrophysics is illustrated by one example of its detailed application to a solar radio burst, including specially designed algorithms of the spectral forward fitting.} \def\gsim{\ \raise 3pt \hbox{$>$} \kern -8.5pt \raise -2pt \hbox{$\sim$}\ } \def\lsim{\ \raise 3pt \hbox{$<$} \kern -8.5pt \raise -2pt \hbox{$\sim$}\ } | The phenomenon of transition radiation was discovered theoretically by two Nobel Prize winning (2003 and 1958 respectively) physicists \cite{Gin_Fr}. Ginzburg and Frank (1946) considered a simplest case when a charged particle passed through a boundary between two dielectrically different media and so generated waves due to a variation of the dielectric constant at the boundary. Remarkably, no acceleration of the particle is necessary to produce the emission due to transition through the boundary. It is easy to understand that a similar effect of electromagnetic emission will take place if a medium is uniformly filled by turbulence that produces fluctuations of the dielectric constant throughout the whole volume rather than at an isolated boundary. Many astrophysical sources, especially those under strong energy release, are believed to be filled by turbulent, randomly inhomogeneous plasma and fast, nonthermal particles. In this situation, an efficient contribution of the transition radiation to the overall electromagnetic emission should be produced. Therefore, distinguishing this contribution from competing mechanisms is important. Below we describe the fundamentals of the transition radiation produced in a magnetized turbulent plasma, and demonstrate its high potential for astrophysical applications. | Nita et al. (2005) proved that the dm continuum component of this solar radio burst is produced by RTR and derived the level of the microturbulence in the plasma to be $\left<\Delta n^2\right>/n^2 = 10^{-5}$. This finding is potentially very important for other cosmic objects. Indeed, the obtained microturbulence level is not particularly strong and much stronger turbulence is expected in many cases, especially, when there is a strong release of the energy at the source. Sometimes, such energy release gives rise to a relativistic expansion of the source, so the emission spectrum is Doppler-boosted and RTR produced at the local plasma frequency can be observed at the Earth even from relatively tenuous sources with low plasma frequency. In this study we present more evidence in favor of RTR generation at the dm continuum solar bursts and use this emission component to derive additional plasma parameters. In particular, we determine the mean plasma frequency and its dispersion at the source in the course of time. Interestingly, these two parameters do not change much during the time of the dm burst. We note that these parameters are obtained from the total power spectra recorded without spatial resolution. Fig. 2 demonstrates that the radio sources at various frequencies do not coincide exactly. Therefore, in cases where a sequence of spatially resolved spectra are available we would be able to study the structure of the flaring plasma density in much greater detail as well as the distribution of the microturbulence over the source. Generally speaking, the RTR contribution is also informative about the fast electrons producing it. In the example presented in Fig. 3 we show the emission decay constants, which can be associated with the fast electron life times. In our case we used exponential fragments of the light curves at the late decay phase of the emission, since no exponential phase was found in the early decay phase. We found the life time to be within 10-40 sec, which corresponds to the electrons of 300 keV or larger in the case of dense flare plasma available in this event. On the other hand, we can expect that most of the RTR emission (around the peak of the burst) is produced by the electrons with E=100-200 keV (Nita et al. 2005). This apparent contradiction can be easily resolved if we recall that the lower energy electrons have the life time of only a few seconds in the given dense plasma, so they die even before the light curves reach the exponential decay stage, and we observe the RTR contribution from preferentially higher energy electrons late in the event. | 7 | 10 | 0710.0351 |
0710 | 0710.5574_arXiv.txt | We study the effect of quasar feedback on distributions of baryons and properties of intracluster medium in galaxy groups using high-resolution numerical simulations. We use the entropy-conserving Gadget code that includes gas cooling and star formation, modified to include a physically-based model of quasar feedback. For a sample of ten galaxy group-sized dark matter halos with masses in the range of $1$ to $5\times 10^{13} M_{\odot}/h$, star formation is suppressed by more than 50\% in the inner regions due to the additional pressure support by quasar feedback, while gas is driven from the inner region towards the outer region of the halos. As a result, the average gas density is 50\% lower in the inner region and 10\% higher in the outer region in the simulation, compared to a similar simulation with no quasar feedback. Gas pressure is lowered by about 40\% in the inner region and higher in the outer region, while temperature and entropy are enhanced in the inner region by about 20-40\%. The total group gas fraction in the two simulations generally differs by less than 10\%. We also find a small change of the total thermal Sunyaev-Zeldovich distortion, leading to 10\% changes in the microwave angular power spectrum at angular scales below two arcminutes. % | Galaxy clusters and groups are the largest gravitationally bound objects in the universe, and they dominate the total baryon content of the universe. Their spatial distribution and mass function contain information about the formation and evolution of large-scale structure, which in turn constrain a variety of fundamental cosmological properties including normalization of the matter power spectrum, the cosmic baryon density, and dark matter properties. However, in order to use them as a cosmological probe, it is necessary to understand their astrophysical properties, and in particular their baryon physics. This issue is of particular current interest due to upcoming arcminute-resolution microwave sky surveys like ACT \citep{kosowsky06,fowler07} and SPT \citep{ruhl05}, which will image galaxy clusters via the Sunyaev-Zeldovich distortions to the cosmic microwave blackbody spectrum from the hot electrons in the cluster gas \citep{sz80}. The majority of baryons in clusters and groups are in the form of hot intracluster gas rather than than individual galaxies. Properties of the Intracluster Medium (ICM) have been studied through a combination of X-ray and radio observations \citep{nulsen05, heinz02, fabian2000}. Although the dark matter distribution in galaxy clusters follow a self-similar relation \citep{pointecouteu05,vikhilin06}, the hot gas does not \citep{sanderson03,popesso05}. Additional non-gravitational sources of heating are required to explain the observations. One interesting and plausible possibility is the energy radiated from quasars or Active Galactic Nuclei (AGN) and deposited into the ICM \citep{kaiser91,valageassilk99,nath02, evan05, evan06}, which we study in this work. The best arena in which to study the impact of various feedback mechanisms is galaxy groups. Massive clusters with deeper gravitational potential wells are likely to have their global thermodynamic and morphological properties less affected by feedback. In comparison, galaxy groups have shallower potential wells while still having enough gas to display the effect of feedback on the ICM. Galaxy groups have recently been observed in X-rays at redshifts as large as $z=0.6$ \citep{willis05}. In the optical band, \cite{sdss07} have compiled group catalogs from the SDSS Data Release 5 catalog. Evidence for heating by a central AGN or radio source in galaxy groups and clusters has been the subject of several recent papers \citep{croston05, jetha06,sanderson05}. These observations show excess entropy in cluster cores, which suggests that some heating process must act to offset cooling. In recent years, cosmological simulations including dark matter and gas have been able to follow the evolution of individual galaxy groups and clusters. A number of studies have investigated the cluster baryon fraction and its evolution in numerical simulations. Adiabatic simulations that do not include radiative cooling find cluster baryon fractions around $0.85$ of the universal baryon fraction \cite{evrard90,metzler_evrard94,navarro_etal95,lubin_etal96,eke_etal98,frenk_etal99,mohr_etal99,bialek_etal01}. Preheating the gas reduces the fraction further \citep{bialek_etal01,borgani_etal02,muanwong_etal02,kay_etal03}. When cooling, star formation and other feedback processes are included, the baryon fraction is higher than that obtained from adiabatic simulations \citep{muanwong_etal02,kay_etal03,valdarnini03,ettori04,nagai07}. This leads to an ``overcooling'' problem and indicates an additional feedback mechanism. In the current study, we analyze the effect of quasar feedback on the baryon distribution and thermodynamics of hot gas in galaxy groups. We also study its implication for the Sunyaev-Zeldovich angular power spectrum, which receives a dominant contribution from high-redshift halos. Ref.~\cite{ks02} showed that the thermal SZ angular power spectrum provides a strong constraint on the normalization of the matter power spectrum, $\sigma_8$. Upcoming SZ surveys like ACT or SPT will have sufficient sensitivity to determine $\sigma_8$ with an accuracy limited by uncertainty in the theoretical model. Also, the kinematic SZ effect is a measure of bulk motions in the universe and may be a competitive probe for studying cosmology \citep{sehgal04, bk06, penn06, dedeo05, maturi07, roncarelli07}. But one of the major sources of uncertainty in modeling the kSZ effect is the gas fraction and its evolution. So understanding both the thermal and kinematic SZ signals requires detailed understanding of feedback mechanisms in galaxy clusters and groups. The mechanisms and effects of feedback are also a long-standing question in astrophysics, with particular bearing on the process of galaxy formation. To this end, we have analyzed a sample of ten galaxy groups at $z=1$ from numerical cosmological simulations of gas and dark matter which have been extended to include a self-consistent model for the evolution of massive black holes and their baryon feedback. At redshift $z>1$, the quasar mode of black hole accretion is expected to be the dominant feedback mechanism, compared to the radio-loud accretion mode which becomes important at lower redshifts \citep{sijacki07}. The size of our simulations prevents studying feedback in galaxy clusters, but rather restricts us to less massive galaxy groups. But as already mentioned, galaxy groups with shallow potential wells provide the best place to study non-gravitational heating and its implications for the properties of hot gas. High-redshift galaxy groups are also a major contributor to the thermal SZ power spectrum, which peaks around $z\approx 1$, when galaxy groups are more numerous than massive clusters \cite{ks02}. Following this introduction, Section II describes our simulation and its implementation of quasar feedback. In Section III we study the effect of numerical resolution on our results; in Section IV we describe our results and compare them with a simulation that do not include quasar feedback. Finally, in Section V we summarize our results and discuss directions for future work, including motivations and prospects for studying more massive galaxy clusters and more realistic feedback modeling for quasars and AGN. | 1. Compared to the no-feedback case, star formation is suppressed by 30-40\% in the inner regions of the halos because of the additional pressure support provided by quasar feedback. 2. Quasar feedback redistributes hot gas, driving it from the inner region towards the outer part of the halos. As a result, gas density is 20\% less in the inner part and 10\% to 15\% greater in the outer region when compared to the simulation without feedback. However, the gas fraction in the two simulation differs by only 5\% to 10\%, and gas fractions tends to increase mildly with increasing halo mass. 3. The ratio of gas mass to stellar mass increases by a factor of 3.5 in the simulation including quasar feedback and a factor of 2.5 in the simulation without quasar feedback in the region $0.2 R_{200m}<R<0.5 R_{200m}$. This contradicts the common assumption that this ratio is constant at all radii. 4. Both temperature and entropy increase by 30\% to 50\% in the halo core region because of the additional thermal energy radiated by quasars. 5. Pressure decreases by 30\% in the inner region and increases by 15\% to 20\% at radii larger than 0.4 $R_{200m}$ due to the increased gas density in this region. This leads to a change of about 6\% in the mean Sunyaev-Zeldovich $Y$-distortion. The resulting SZ angular power spectrum will be larger by around 10\% for $l>5000$. We find little dependence of the SZ enhancement with halo mass. The effects of quasar feedback on the intracluster medium will be most evident in the group-sized haloes considered here, with their relatively shallow gravitational potential wells. Observationally, the most interesting haloes are larger in mass by a factor of ten, galaxy clusters: these are the haloes which are most readily detected via their SZ, X-ray, or optical signals. The gas fractions do not show any particular trend with increasing halo mass, and star fractions increase very weakly with mass over the halo mass range studied here. So it is reasonable to expect that the results of this work will hold for cluster-sized halos as well. Nevertheless, given the substantial impact of quasar feedback on various properties of the intracluster medium which the current study suggests, it is imperative to study cluster-sized halos as well. This requires larger-volume simulations, as the number density of clusters decreases with cluster mass. To this end, we are currently running a simulation of box size $50\,{\rm Mpc}/h$; results will be reported elsewhere. This is the first attempt to study the impact of quasar feedback on the baryon fraction and thermodynamics of the intracluster medium in a cosmological hydrodynamic simulation. Both the gas and star fractions in our simulation are consistent with current observational limits \citep{allen_etal04, ettori99}. Note that we have studied only quasar feedback at redshifts greater than unity. However, active galactic nuclei also inject energy into the ICM via a ``radio mode'' which is believed to be the dominant feedback mechanism at lower redshift \citep{sijacki05, sijacki07}. Thus our results should be treated as a conservative estimate of the total impact of AGN feedback for galaxy groups at low redshifts. Gas pressure in cosmological halos, particularly those with masses ranging from galaxy groups to galaxy clusters, determines the important thermal Sunyaev-Zeldovich signal which will soon be measured with high precision. The gas fraction is important for connecting kinematic SZ signals of cluster gas momentum with theoretical predictions about cluster velocity or total momentum. This paper takes the first step towards quantifying the impact of quasars on these quantities, which turns out to be significant but not dominating. Much work remains to be done, both through larger simulations which contain many galaxy-cluster-sized haloes, and in enhancing the realism of the quasar feedback models. We hope the results here plus the exciting observational prospects in the near future will open the door to further advances in this area. | 7 | 10 | 0710.5574 |
0710 | 0710.2481_arXiv.txt | The stellar populations in the bulges of S0s, together with the galaxies' dynamics, masses and globular clusters, contain very interesting clues about their formation. I present here recent evidence suggesting that S0s are the descendants of fading spirals whose star formation ceased. | Combining published data with high-quality VLT/FORS spectroscopy of sample of Fornax S0s (Bedregal et al.\ 2006a) we have carried out a combined study of the Tully-Fisher relation and the stellar populations of these galaxies. Despite the relatively small sample and the considerable technical challenges involved in determining the true rotation velocity $V_{\rm rot}$ from absorption line spectra of galaxies with significant non-rotational support (see Mathieu et al.\ 2002), some very interesting results arise. S0s lie systematically below the spiral galaxy Tully-Fisher relation in both the optical and near-infrared (Figure~1). If S0s are the descendants of spiral galaxies, this offset can be naturally interpreted as arising from the luminosity evolution of spiral galaxies that have faded since ceasing star formation. Moreover, the amount of fading implied by the offset of individual S0s from the spiral relation seems to correlate with the luminosity-weighted age of their stellar population, particularly at their centres (Figure~2). This correlation suggests a scenario in which the star formation clock stopped when gas was stripped out from a spiral galaxy and it began to fade into an S0. The stronger correlation at small radii indicates a final last-gasp burst of star formation in this region. See Bedregal, Arag\'on-Salamanca \& Merrifield (2006b) for details. \begin{figure} \begin{center} \includegraphics[height=3.3in,width=4.0in,angle=0]{Aragon-Salamanca_fig1.ps} \end{center} \caption{$B$-band Tully-Fisher relation (TFR) for S0 galaxies using different samples from the literature (open symbols) and our VLT Fornax data (filled circles). The solid and dashed lines show two independent determinations of the TFR relation for local spirals. On average (dotted line), S0s are $\sim3$ times fainter than spirals at similar rotation velocities (Bedregal, Arag\'on-Salamanca \& Merrifield 2006b). } \label{fig:fig1} \end{figure} \begin{figure} \begin{center} \includegraphics[height=1.65in,angle=0]{Aragon-Salamanca_fig2.eps} \end{center} \caption{ For our VLT Fornax data we plot the shift in magnitudes from the $B$-band spiral TFR versus the stellar population age at the galaxy centre (left panel), at $1\,R_e$ (middle panel) and at $2\,R_e$ (right panel). The lines show models for fading spirals. Note that the correlation is strongest for the central stellar populations of the galaxies, suggesting that the last episode of star formation took place there (Bedregal, Arag\'on-Salamanca \& Merrifield 2006b). } \label{fig:fig2} \end{figure} | The stellar populations, dynamics and globular clusters of S0s provide evidence consistent with these galaxies being the descendants of fading spirals whose star formation ceased. However, caution is needed since significant problems could still exist with this picture (see, e.g., Christlein \& Zabludoff 2004; Boselli \& Gavazzi 2006). Moreover, the number of galaxies studied here is still small, and it would be highly desirable to extend this kind of studies to much larger samples covering a broad range of galaxy masses and environments. \begin{figure} \begin{center} \includegraphics[height=2.9in,angle=0]{Aragon-Salamanca_fig3.eps} \end{center} \caption{ Log$_{10}$ of the luminosity-weighted ages is Gyr vs.\ the globular cluster specific frequency ($S_N$) of S0s. The line shows the evolution expected for a fading galaxy according to the stellar population models of Bruzual \& Charlot (2003). The correlation between the fading of the galaxies (or increase in $S_N$) and the spectroscopically-determined age of their stellar populations is clearly consistent with the predictions of a simple fading model. Note that the $S_N$ value for NGC3115B is very unreliable and almost certainly severely overestimated due to contamination from the GC systems of neighbouring galaxies. See Barr et al.\ (2007) for details. } \label{fig:fig3} \end{figure} | 7 | 10 | 0710.2481 |
0710 | 0710.0484_arXiv.txt | We study the propagation of relativistic jets originating from AGNs within the Interstellar/Intergalactic Medium of their host galaxies, and use it to build a model for the suppression of stellar formation within the expanding cocoon. | \label{sec:jetcocoon} Relativistic jets from Supermassive Black Holes hosted within the bulges of spirals and nuclei of early-type galaxies inject significant amounts of energy into the medium within which they propagate, \begin{figure} \centering \includegraphics[scale=0.6,angle=90]{antonucciodelogu_fig1.eps} \caption{Expansion of the cocoon within the ISM. The jet has an input mechanical power $\textrm{P}_{jet} = 10^{46} \textrm{ergs}\cdot\textrm{cm}^{-3}\cdot\textrm{sec}^{-1}$, and the ISM density is: $\, n_{e}^{ism} = 1 e^{-} cm^{-3}$.}\label{fig:cocoon} \end{figure} creating an extended, underdense and hot cocoon. We have performed a series of simulations of these jets and cocoons using an AMR code, FLASH 2.5: a typical output is shown in fig.~\ref{fig:cocoon}.\\ We find that a slight modification of the exact models of \cite{1991MNRAS.250..581F} and \cite{1997MNRAS.286..215K}, approximates well the simulations. | 7 | 10 | 0710.0484 |
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0710 | 0710.2815_arXiv.txt | We present evolutionary tracks of binary systems with high mass companion stars and stellar-through-intermediate mass BHs. Using Eggleton's stellar evolution code, we compute the luminosity produced by accretion from the donor during its entire evolution. We compute also the evolution of the optical spectrum of the binary system taking the disc contribution and irradiation effects into account. The calculations presented here can be used to constrain the properties of the donor stars in Ultraluminous X-ray Sources by comparing their position on the HR or color-magnitude diagrams with the evolutionary tracks of massive BH binaries. This approach may actually provide interesting clues also on the properties of the binary system itself, including the BH mass. We found that, on the basis of their position on the color-magnitude diagram, some of the candidate counterparts considered can be ruled out and more stringent constraints can be applied to the donor masses. | Point like off nuclear X-ray sources with luminosities well in excess of the Eddington limit for a stellar mass black hole, have been discovered in a large number of nearby galaxies (e.g.~\citealt{2002ApJS..143...25C}; \citealt{sgtw04}; \citealt{2005ApJS..157...59L}). These Ultraluminous X-ray Sources (ULXs) are too dim to be low luminosity AGNs and too bright to be normal X-ray binaries (XRBs) emitting below the Eddington limit. In this paper we define a ULX as a source with bolometric luminosity in excess of the Eddington limit for a stellar mass black hole of $20\msun$ and less luminous than a few $10^{41}\rm erg/s$. This definition implies that a Galactic XRB radiating isotropically at or below the Eddington limit can not be a ULX. In spiral or starburst galaxies ULXs turn out to be associated with star forming regions, emission nebulae and stellar clusters (\citealt{zfrm02}, \citealt{pm02}). These facts along with, in some case, the detection of stellar optical counterparts (\citealt{2001MNRAS.325L...7R}; \citealt{2002MNRAS.335L..67G}; \citealt{2002ApJ...580L..31L}; \citealt{2004ApJ...602..249L};\citealt{kwz04};\citealt{zamp1}; \citealt{2005ApJ...629..233K}; \citealt{mucetal05, mucetal07}; \citealt{scpmw04}) strongly indicate an association between ULXs and young massive stars, although the nature of the accreting compact object remains unclear. A certain number of ULXs, however, appear as isolated X-ray sources with no obvious counterpart at any wavelength (\citealt{sm04}) and without a clear association with star forming regions or emission nebulae. Some low-luminosity ULXs may possibly be present also in elliptical galaxies (\citealt{jcbg03}), but their nature and the actual evidence for their existence is not well established (see e.g. \citealt{iba04}, \citealt{aglc04}, \citealt{2005ApJ...622L..89G}). The nature of the ULXs in spiral galaxies is less controversial. Different models have been proposed to explain the large luminosities reached by these sources. One of the favored models consists of an intermediate mass black hole (IMBH) with a mass in the range $10^{2}-10^{3}M_{\rm \odot}$, accreting from an high mass donor star. The presence of an IMBH can account for most of the observational properties of ULXs in a rather straightforward way. For instance, the observed cool disc spectra of some ULX can be explained with the fact that the innermost stable circular orbit of an IMBH is larger than that of a stellar mass black hole (e.g. \citealt{mfm04}). The detection of a $\sim$50-160 mHz quasi periodic oscillations in the power density spectrum of M82 X-1 and NGC5408 X-1 (\citealt{2003ApJ...586L..61S}, \citealt{2004ApJ...614L.113F}, \citealt{2006MNRAS.365.1123M}; \citealt{stroh07}), the very high luminosity of some ULXs ($\sim 10^{41} \ergs$) along with their cool discs, and the energy content and morphology of the nebulae around some of them (\citealt{pm02}) all suggest an IMBH interpretation. The main problem with this interpretation resides in the formation mechanism of such an extreme object. In fact, if IMBHs with masses in excess of $\sim 100 M_\odot$ exist, they will require a new formation root with respect to the stellar black holes in our Galaxy and to the supermassive black holes in Active Galactic Nuclei. Until now two scenarios have been proposed to form a black hole in the intermediate mass range: the runaway collision of massive stars in dense open clusters (\citealt{2004Natur.428..724P}, \citealt{2004ApJ...604..632G}) and the primordial collapse of a very high mass star with zero metallicity (\citealt{2000ApJ...540...39A}, \citealt{2001ApJ...551L..27M}). Both the mechanisms however suffer of a certain degree of uncertainty related to the incomplete knowledge of the behavior of very massive stars. Therefore we have no final evidence that an IMBH can really form. Furthermore, the interpretation of the soft components observed in some ULXs in terms of cool accretion discs is not univocal (e.g. \citealt{2005ApJ...631L.109C}, \citealt{2005ApJ...635..198D}, \citealt{2005MNRAS.357.1363R}, \citealt{fk06}, \citealt{gonc06}, \citealt{2006MNRAS.368..397S}). Other interpretations in terms of stellar (or quasi-stellar) mass black holes have been proposed. A mechanical (\citealt{kdwfe01}) or a relativistic beaming (\citealt{kfm02}) can reproduce the observed luminosities of ULXs up to a few $10^{40}\ergs$ with a beaming factor around $\sim 10.$ At most, as suggested by \citealt{2005MNRAS.357..275K}, accretion from helium rich matter from a geometrically thick disc can generate luminosities up to $\sim 5\times10^{40}\ergs$. However, luminosities in excess of $5\times 10^{40}\ergs$ ($\sim 5$\% of the ULX population) and the isotropy of the ionized nebulae around some ULXs can not be easily explained in terms of beaming models. On the other hand, photon bubbles disc instabilities (\citealt{b02}, \citealt{b06}) and emission from a slim disc (\citealt{wmm01}, \citealt{ezkmw03}) can produce genuine isotropic super-Eddington luminosities around 10 times the Eddington limit. As shown by \cite{2005MNRAS.356..401R}, a normal binary with a stellar mass black hole and a donor star with initial mass $\simgreat 10\mdot$, can in principle explain a large sample of ULXs if a super Eddington luminosity around $10$ is allowed (see also \citealt{2003MNRAS.341..385P}). However, the typical temperature of a slim disc in this regime ($\sim$ 1--2 keV) is inconsistent with the observation of some cool disc sources which are, at the same time, the most luminous ULXs. Furthermore, luminosities in excess of a few $10^{40}\ergs$ are not attainable with slim discs around stellar mass black holes. On the other hand, the photon bubble instability model makes no predictions on the observable accretion disc temperature and therefore cannot be compared directly with observations. Thanks to the high precision astrometry of the {\it Chandra} X-ray observatory, we know that some ULXs have optical counterparts which match with the X-ray source position (\citealt{2002ApJ...580L..31L}, \citealt{2004ApJ...602..249L}, \citealt{zamp1}, \citealt{kwz04}, \citealt{sm04}, \citealt{2005ApJ...629..233K}, \citealt{mucetal05, mucetal07}, \citealt{Liu07}). Most of them are suspected to be main sequence (MS) high mass stars or supergiant stars of uncertain spectral type. The ambiguity arises because, until now, optical spectra of these stars either are not available, or they are too noisy and with peculiar spectral features (\citealt{2004ApJ...602..249L}), or there is more than an optical counterpart in the X-ray error box (\citealt{sm04}, \citealt{mucetal05}). In this paper, we present the evolutionary tracks of binary systems with high mass companion stars and stellar-through-intermediate mass BHs and compare them with the properties of the donor stars in ULX binary systems. In \S 2 we present the model adopted to evolve the binary system. In \S 3, we summarize the properties of the four ULXs considered and their optical counterparts. Our results are presented in \S 4 and compared with the properties of ULX counterparts in \S 5. Conclusions follow in \S 6. | In this paper we have shown that the effects of the binary evolution on the optical properties of a donor star in a ULX binary system can significantly alter its colors and evolution. We included also the contribution of the optical emission of the accretion disc and the X-ray irradiation of the donor and disc surfaces. The evolutionary track of a donor in an accreting binary on the CM diagram is very different with respect to that of a single isolated star. These important differences can not be overlooked when trying to identify the donor mass of a ULX. We calculated tracks for stellar and intermediate mass black holes, and demonstrate the brighter nature of the IMBH as an intrinsic phenomenon produced by the low mass ratio in the binary. The photometric data of four ULX counterparts have been compared with the evolutionary tracks of massive donors on the CM diagram. We find that the counterparts 2, 3 and 4 of NGC 4559 X-7 are consistent only with donors of $10$--$15\msun$ undergoing a case B or C mass transfer episode. The only possibility for objects 5 and 8 is a very massive ($\sim50\msun$) MS donor or a H-shell burning donor with mass between 10-15 and 30$\msun$. Finally, object 1 is compatible only with a very massive ($M\simgreat 50\msun$) donor in a H-shell burning phase. Our results are quite different with respect to what obtained in previous analysis. \cite{scpmw04} compared theoretical evolutionary tracks of single stars between 9 and 25$\msun$ with the position in the CM diagram of the counterparts of NGC4559 X-7. They find that the colors and luminosity of six out of seven stars are consistent with the tracks of main sequence, blue or red supergiant with masses in the range 10-15 $M_\odot$ and ages $\sim 20$ Myr. Their favored counterpart, object 1, was identified with a main sequence or blue supergiant of $\sim 20 M_\odot$ and an age of $\sim 10$ Myr, although \cite{scpmw04} recognize that X-ray irradiation may affect its colors and hence its classification. Irradiation effects were taken into account by \cite{cpw07}, giving a donor mass of 5-20$M_{\odot}$ for object $1$. In their work, however, \cite{cpw07} considered irradiation on single stars and did not take into account the effects of the markedly different evolution of a donor star in a binary system. As far as NGC1313 X-2 is concerned, we definitely rule out the candidate counterpart C2 and identify C1 as a H-shell burning donor of 10$\msun<$M$<$15$\msun$ around a stellar mass black hole or a $\simless 15 \msun$ donor undergoing RLOF during MS (or the H-shell burning phase) in an IMBH binary system. On statistical grounds alone, it is then more likely that NGC1313 X-2 hosts a $\sim 100 \msun$ BH rather than a stellar mass BH as the H-shell burning phase is much shorter than the MS phase. Our result for object C2 leaves only C1 as the likely counterpart of NGC 1313 X-2, in agreement with the evidence coming from the refined X-ray astrometry of the field recently reported by \cite{Liu07}. On the basis of single star isochrone fitting of the parent stellar population, \cite{pak06} and \cite{ramsey06} estimate a maximum MS mass of 8--9$\msun$ for object C1. \cite{Liu07} find consistency with either a $\sim 8\msun$ star of very low metallicity or an O spectral type, solar metellicity star of $\sim 30\msun$. Our estimated range of donor masses for object C1 appears more in agreement with the value (10-18$\msun$) reported by \cite{mucetal07}, probably because they also included the contribution of the disc and irradiation effects. In Holmberg II X-1 there are several possibilities, and we can only rule out stars with M$<$10$\msun$ (\citealt{cpw07} give a lower bound of 5$\msun$), while for M81 X-9 we are left with the unique possibility that the donor is a $\sim 10\msun$ star in the H-shell burning phase. We conclude therefore that in all previous work where the binary evolution and/or irradiation effects ware not taken into account, the donor mass has been systematically overestimated. We find that mass transfer occurring at late stages is very unlikely, not only for the short timescale of the contact phase ($10^{3}$--$10^{5}$yrs) but also for the incompatibility of the optical colors with the majority of the observed counterparts. If mass transfer sets in during MS, two contact phases occur (segments A--B and C--D). The mass transfer timescales are $t_{A-B}\sim 10^{6}-10^{7}\rm\,yrs$ and $t_{C-D}\sim 10^{3}-10^{5}\rm\,yrs$ (depending on the donor mass). Therefore assuming a flat distribution of periods between 1 and $\sim$10 days, as during MS the contact phase is reached for $P_{orb}\simless 2$days, we expect a factor $\sim \frac{10}{2}\times \frac{t_{A-B}}{t_{C-D}}\simeq 500$--$5000$ more main sequence systems than H-shell burning donors. | 7 | 10 | 0710.2815 |
0710 | 0710.4328_arXiv.txt | The initial conditions and relevant physics for the formation of the earliest galaxies are well specified in the concordance cosmology. Using ab initio cosmological Eulerian adaptive mesh refinement radiation hydrodynamical calculations, we discuss how very massive stars start the process of cosmological reionization. The models include non-equilibrium primordial gas chemistry and cooling processes and accurate radiation transport in the Case B approximation using adaptively ray traced photon packages, retaining the time derivative in the transport equation. Supernova feedback is modeled by thermal explosions triggered at parsec scales. All calculations resolve the local Jeans length by at least 16 grid cells at all times and as such cover a spatial dynamic range of $\sim$10$^6$. These first sources of reionization are highly intermittent and anisotropic and first photoionize the small scales voids surrounding the halos they form in, rather than the dense filaments they are embedded in. As the merging objects form larger, dwarf sized galaxies, the escape fraction of UV radiation decreases and the \ion{H}{2} regions only break out on some sides of the galaxies making them even more anisotropic. In three cases, SN blast waves induce star formation in overdense regions that were formed earlier from ionization front instabilities. These stars form tens of parsecs away from the center of their parent DM halo. Approximately 5 ionizing photons are needed per sustained ionization when star formation in 10$^6$ \Ms~halos are dominant in the calculation. As the halos become larger than $\sim$$10^7 \Ms$, the ionizing photon escape fraction decreases, which in turn increases the number of photons per ionization to 15--50, in calculations with stellar feedback only. Radiative feedback decreases clumping factors by 25 per cent when compared to simulations without star formation and increases the average temperature of ionized gas to values between 3,000 and 10,000 K. | It is clear that quasars are not responsible to keep the universe ionized at redshift 6. The very brightest galaxies at those redshifts alone also provide few photons. The dominant sources of reionization so far are observationally unknown despite remarkable advances in finding sources at high redshift \citep[e.g.][]{Shapiro86, Bouwens04, Fan06, Thompson07, Eyles06} and hints for a large number of unresolved sources at very high redshifts \citep{Spergel07, Kashlinsky07} which is still a topic of debate \citep{Cooray07, Thompson07}. At the same time, ab initio numerical simulations of structure formation in the concordance model of structure formation have found that the first luminous objects in the universe are formed inside of cold dark matter (CDM) dominated halos of total masses $2 \times 10^5 - 10^6 \Ms$ \citep{Haiman96, Tegmark97, Abel98}. Fully cosmological ab initio calculations of \citet{Abel00, Abel02} and more recently \citet{Yoshida06} clearly show that these objects will form isolated very massive stars. Such stars will be copious emitters of ultraviolet (UV) radiation and are as such prime suspect to get the process of cosmological reionization started. In fact, one dimensional calculations of \citet{Whalen04} and \citet{Kitayama04} have already argued that the earliest \ion{H}{2} regions will evaporate the gas from the host halos and that in fact most of the UV radiation of such stars would escape into the intergalactic medium. Recently, \citet{Yoshida07a} and \citet{Abel07} demonstrated with full three-dimensional radiation hydrodynamical simulations that indeed the first \ion{H}{2} regions break out of their host halos quickly and fully disrupt the gaseous component of the cosmological parent halo. All of this gas finds itself radially moving away from the star at $\sim30\kms$ at a distance of $\sim100$ pc at the end of the stars life. At this time, the photo-ionized regions have now high electron fractions and little destructive Lyman-Werner band radiation fields creating ideal conditions for molecular hydrogen formation which may in fact stimulate further star formation above levels that would have occurred without the pre-ionization. Such conclusion have been obtained in calculations with approximations to multi dimensional radiative transfer or one dimensional numerical models \citep{Ricotti02a, Nagakura05, OShea05, Yoshida06, Ahn07, Johnson07}. These early stars may also explode in supernovae and rapidly enrich the surrounding material with heavy elements, deposit kinetic energy and entropy to the gas out of which subsequent structure is to form. This illustrates some of the complex interplay of star formation, primordial gas chemistry, radiative and supernova feedback and readily explains why any reliable results will only be obtained using full ab initio three dimensional hydrodynamical simulations. In this paper, we present the most detailed such calculations yet carried out to date and discuss issues important to the understanding of the process of cosmological reionization. It is timely to develop direct numerical models of early structure formation and cosmological reionization as considerable efforts are underway to \begin{enumerate} \item Observationally find the earliest galaxies with the James Webb Space Telescope \citep[JWST;][]{Gardner06} and the Atacama Large Millimeter Array \citep[ALMA;][]{Wilson05}, \item Further constrain the amount and spatial non-uniformity of the polarization of the cosmic microwave background radiation \citep{Page07}, \item Measure the surface of reionization with LOFAR \citep{Rottgering06}, MWA \citep{Bowman07}, GMRT \citep{Swarup91} and the Square Kilometer Array \citep[SKA;][]{Schilizzi04}, and \item Find high redshift gamma ray bursts with SWIFT \citep{Gehrels04} and their infrared follow up observations. \end{enumerate} We begin by describing the cosmological simulations that include primordial star formation and accurate radiative transfer. In \S\ref{sec:SF}, we report the details of the star formation environments and host halos in our calculations. Then in \S\ref{sec:reion}, we describe the resulting start of cosmological reionization, and investigate the environments in which these primordial stars form and the evolution of the clumping factor. We compare our results to previous calculations and further describe the nature of the primordial star formation and feedback in \S\ref{sec:discussion}. Finally we summarize our results in the last section. \begin{deluxetable*}{lccccccc} \tablecolumns{8} \tabletypesize{} \tablewidth{\textwidth} \tablecaption{Simulation Parameters\label{tab:sims}} \tablehead{ \colhead{Name} & \colhead{$l$} & \colhead{Cooling model} & \colhead{SF} & \colhead{SNe} & \colhead{N$_{\rm{part}}$} & \colhead{N$_{\rm{grid}}$} & \colhead{N$_{\rm{cell}}$} \\ \colhead{} & \colhead{[Mpc]} & \colhead{} & \colhead{} & \colhead{} & \colhead{} & \colhead{} & \colhead{} } \startdata SimA-Adb & 1.0 & Adiabatic & No & No & 2.22 $\times$ 10$^7$ & 30230 & 9.31 $\times$ 10$^7$ (453$^3$) \\ SimA-HHe & 1.0 & H, He & No & No & 2.22 $\times$ 10$^7$ & 40601 & 1.20 $\times$ 10$^8$ (494$^3$) \\ SimA-RT & 1.0 & H, He, \hh & Yes & No & 2.22 $\times$ 10$^7$ & 44664 & 1.19 $\times$ 10$^8$ (493$^3$) \\ SimB-Adb & 1.5 & Adiabatic & No & No & 1.26 $\times$ 10$^7$ & 23227 & 6.47 $\times$ 10$^7$ (402$^3$) \\ SimB-HHe & 1.5 & H, He & No & No & 1.26 $\times$ 10$^7$ & 21409 & 6.51 $\times$ 10$^7$ (402$^3$) \\ SimB-RT & 1.5 & H, He, \hh & Yes & No & 1.26 $\times$ 10$^7$ & 24013 & 6.54 $\times$ 10$^7$ (403$^3$) \\ SimB-SN & 1.5 & H, He, \hh & Yes & Yes & 1.26 $\times$ 10$^7$ & 24996 & 6.39 $\times$ 10$^7$ (400$^3$) \enddata \tablecomments{Col. (1): Simulation name. Col. (2): Box size. Col. (3): Cooling model. Col. (4): Star formation. Col. (5): Supernova feedback. Col. (6): Number of dark matter particles. Col. (7): Number of AMR grids. Col. (8): Number of unique grid cells.} \end{deluxetable*} | \label{sec:discussion} We have studied the details of massive metal-free star formation and its role in the start of cosmological reionization. We have treated star formation and radiation in a self-consistent manner, allowing for an accurate investigation of the evolution of cosmic structure under the influence of early Pop III stars. Stellar radiation from these stars provides thermal, dynamical, and ionizing feedback to the host halos and IGM. Although Pop III stars are not thought to provide the majority of ionizing photons needed for cosmological reionization, they play a key role in the early universe because early galaxies that form in these relic \ion{H}{2} regions are significantly affected by Pop III feedback. Hence it is important to consider primordial stellar feedback while studying early galaxy formation. In this section, we compare our results to previous numerical simulations and semi-analytic models of reionization and then discuss any potential caveats of our methods and possible future directions of this line of research. \begin{figure}[t] \begin{center} \epsscale{1.15} \plotone{f14_color} \caption{\label{fig:filtering} The Jeans mass $M_J$ and filtering mass $M_F$ that can form bound objects. The squares denote the total mass of star forming halos in all three simulations.} \end{center} \end{figure} \subsection{Comparison to Previous Models} \subsubsection{Filtering Mass} \label{sec:filtering} One source of negative feedback is the suppression of gas accretion into potential wells when the IGM is preheated. The lower limit of the mass of a star forming halo is the Jeans filtering mass \begin{equation} \label{eqn:filtering} M_F^{2/3}(a) = \frac{3}{a} \int_0^a \> da^\prime M_J^{2/3}(a^\prime) \left[ 1 - \left(\frac{a^\prime}{a}\right) \right], \end{equation} where $a$ and $M_J$ are the scale factor and time dependent Jeans mass in the \ion{H}{2} region \citep{Gnedin98, Gnedin00b}. Additionally, the virial shocks are weakened if the accreting gas is preheated and will reduce the collisional ionization in halos with $\tvir \gsim 10^4$ K. To illustrate the effect of Jeans smoothing, we take the large \ion{H}{2} region of SimB-SN because it has the largest ionized filling fraction, which is constantly being heated after $z = 21$. Temperatures in this region fluctuates between 1,000~K and 30,000~K, depending on the proximity of the currently living stars. In Figure \ref{fig:filtering}, we show the resulting filtering mass of regions with an ionization fraction greater than $10^{-3}$ along with the total mass of star forming halos. \citet{Gnedin00b} found the minimum mass of a star forming halo is better described by $M_F$ instead of $M_J$. Our simulations are in excellent agreement for halos that are experiencing star formation after reincorporation of their previously expelled gas. The filtering mass is the appropriate choice for a minimum mass in this case as the halo forms from preheated gas. However for halos that have already assembled before they become embedded in a relic \ion{H}{2} region, the appropriate minimum mass $M_{\rm{min}}$ is one that is regulated by the LW background \citep{Machacek01, Wise05} and photo-evaporation \citep[e.g.][]{Efstathiou92, Barkana99, Haiman01, Mesigner06}. This is evident in the multitude of star forming halos below $M_F$. With the exception of star formation induced by SN blast waves or I-fronts, this verifies the justification of using $M_{\rm{min}}$ and $M_F$ for Pop III and galaxy formation, respectively, as a criterion for star forming halos in semi-analytic models. \begin{figure*}[t] \begin{center} \epsscale{1.15} \plotone{f15} \caption{\label{fig:aniso} Density (left) and temperature (right) slices of an anisotropic \ion{H}{2} region in the most massive halo of SimB-RT. The star has lived for 2.5 Myr out of its 2.7 Myr lifetime. The field of view is 900 proper parsecs.} \end{center} \end{figure*} \subsubsection{Star Formation Efficiency} \label{sec:SFeff} Semi-analytic models rely on a star formation efficiency $f_\star$, which is the fraction of collapsed gas that forms stars, to calculate quantities such as emissivities, chemical enrichment, and IGM temperatures. Low-mass halos that form a central star have $f_\star \sim 10^{-3}$ whose value originates from a single 100~\Ms~star forming in a dark matter halo of mass 10$^6$~\Ms~\citep{Abel02, Bromm02, Yoshida06}. Pop II star forming halos are usually calibrated with star formation efficiencies from local dwarf and high-redshift starburst galaxies and are usually on the order of a few percent \citep[e.g.][]{Taylor99, Gnedin00a}. This leads to the question: how efficient is star formation in these high-redshift halos while explicitly considering feedback? This is especially important when halos start to form multiple massive stars and when metallicities are not sufficient to induce Pop II star formation. The critical metallicity for a transition to Pop II is still unclear. Recently, \citet{Jappsen07a} showed that metal line cooling is dynamically unimportant in diffuse gas until metallicities of $10^{-2} \> Z_\odot$. On the other hand, dust that is produced in SNe can generate efficient cooling down in dense gas with $10^{-6} \> Z_\odot$ \citep{Schneider06}. If the progenitors of the more massive halos did not result in a pair-instability SN, massive star formation can continue until it becomes sufficiently enriched. Hence our simulations can probe the efficiency of this scenario of massive metal-free star formation. It has also been suggested that the cosmological conditions that lead to the collapse of a metal-poor molecular cloud ($Z/Z_\odot \approx 10^{-3.5}$) may be more important than some critical metallicity in determining the initial mass function of a given stellar system \citep{Jappsen07b}. We calculate $f_\star$ with the ratio of the sum of the stellar masses to the total gas mass of unique star-forming halos. For example at the final redshift of 15.9 in SimA-RT, the most massive halo and its progenitors had hosted 11 stars and the gas mass of this halo is $1.8 \times 10^6 \Ms$, which results in $f_\star = 6.1 \times 10^{-4}$ for this particular halo. Expanding this quantity to all star forming halos, $f_\star/10^{-4} = 5.6, 6.7, 7.4$ for SimA-RT, SimB-RT, and SimB-SN, respectively. We note that our choice of $M_\star = 170 \Ms$ in SimB-SN increases $f_\star$ by 70\%. Our efficiencies are smaller than the isolated Pop III case because halos cannot form any stars once the first star expels the gas, and 40 -- 75 million years must pass until star can form again when the gas is reincorporated into the halo. By regarding the feedback created by Pop III stars and associated complexities during the assembly of these halos, the $f_\star$ values of $\sim$$6 \times 10^{-4}$ that are explicitly determined from our radiation hydrodynamical simulations provide a more accurate estimate on the early star formation efficiencies. \subsubsection{Intermittent \& Anisotropic Sources} Our treatment of star formation and feedback produces intermittent star formation, especially in low-mass halos. If one does not account for this, star formation rates might be overestimated in this phase of star formation. Kinetic energy feedback is the main cause of this behavior. As discussed in sections \ref{sec:haloMass} and \ref{sec:kinetic}, shock waves created by D-type I-fronts and SN explosions expel most of the gas in halos with masses $\lsim 10^7$ \Ms. A period of quiescence follows these instances of star formation. Then stars are able to form after enough material has accreted back into the halo. Only when the halo becomes massive enough to retain most of the outflows and cool efficiently through \lya~and \hh~radiative processes, star formation becomes more regular with successive stars forming. The central gas structures in the host halo are usually anisotropic as it is acquiring material through accretion along filaments and mergers. At scales smaller than 10 pc, the most optically thick regions produce shadows where the gas radially behind the dense clump is not photo-ionized or photo-heated by the source. This produces cometary and so-called elephant trunk structures that are also seen in local star forming regions and have been discussed in detail since \citet{Pottasch58}. At a larger distance, the surrounding cosmic structure is composed of intersecting or adjacent filaments and satellite halos that breaks spherical symmetry. The filaments and nearby halos are optically thick and remain cool and thus the density structures are largely unchanged. The entropy of dense regions are not increased by stellar radiation and will feel little negative feedback from an entropy floor that only exists in the ionized IGM \citep[cf.][]{Oh03}. Ray-tracing allows for accurate tracking of I-fronts in this inhomogeneous medium. Radiation propagates through the least optically thick path and generates champagne flows that have been studied extensively in the context of present day star formation \citep[e.g.][]{Franco90, Churchwell02, Shu02, Arthur06}. In the context of massive primordial stars, these champagne flows spread into the voids and are impeded by the inflowing filaments. The resulting \ion{H}{2} regions have ``butterfly'' morphologies \citep{Abel99, Abel07, Alvarez06a, Mellema06, Yoshida07a}. We also point out that sources embedded in relic \ion{H}{2} largely maintain or increase the ionization fraction. Here the already low optical depth of the recently ionized medium (within a recombination time) allows the radiation to travel to greater distances than a halo embedded in a completely neutral IGM. The \ion{H}{2} regions become increasingly anisotropic in higher mass halos. We show an example of the morphology of a \ion{H}{2} region near the end of the star's lifetime in a dark matter halo with mass $1.4 \times 10^7 \Ms$ in Figure \ref{fig:aniso}. \subsection{Potential Caveats and Future Directions} Although we have simulated the first generations of stars with radiation hydrodynamic simulations, our methods have neglected some potentially important processes and made an assumption about the Pop III stellar masses. One clear shortcoming of our simulations is the small volume and limited statistics of the objects studied here. However, it was our intention to focus on the effects of Pop III star formation on cosmological reionization and on the formation of an early dwarf galaxy instead of global statistics. The star formation only simulations (SimA-RT and SimB-RT) converge to the similar averaged quantities, e.g. ionized fraction, temperatures, star formation rates, at the final redshift. The evolution of these quantities differ because of the limited number of stars that form in the simulations, which then causes the evolution to depend on individual star formation times. This variance should be expected in the small volumes that we simulate and should not diminish the significance of our results. We have verified even in a 2.5-$\sigma$ peak that Pop III stars cannot fully reionize the universe, which verified previous conclusions that low-luminosity galaxies provide the majority of ionizing photons. Furthermore, it is beneficial to study Pop III stellar feedback because it regulates the nature of star formation in these galaxies that form from pre-heated material. Further radiation hydrodynamics simulations of primordial star and galaxy formation with larger volumes while still resolving the first star forming halos of mass $\sim$$3 \times 10^5 \Ms$ will improve the statistics of early star formation, especially in more typical overdensities, i.e. 1-$\sigma$ peaks, some of which could survive to become dwarf spheroidal galaxies at $z = 0$. In this work, we treated the LW radiation field as optically thin, but in reality, \hh~produces a non-zero optical depth above column densities of $10^{14} \> \mathrm{cm}^{-2}$ \citep{Draine96}. Conversely, Doppler shifts of the LW lines arising from large velocity anisotropies and gradient may render \hh~self-shielding unimportant up to column densities of $10^{20} - 10^{21} \> \mathrm{cm}^{-2}$ \citep{Glover01}. If self-shielding is important, it will lead to increased star formation in low-mass halos even when a nearby source is shining. Moreover, \hh~production can also be catalyzed ahead of I-fronts \citep{Ricotti01, Ahn07}. In these halos, LW radiation will be absorbed before it can dissociate the central \hh~core. On the same topic, we neglect any type of soft UV or LW background that is created by sources that are cosmologically nearby ($\Delta z / z \sim 0.1$). A soft UV background either creates positive or negative feedback, depending on its strength \citep{Mesigner06}, and a LW background increases the minimum halo mass of a star-forming halo \citep{Machacek01, Yoshida03, OShea07, Wise07b}. However in our calculations, the lack of self-shielding, which suppresses star formation in low-mass halos, and the neglect of a LW background, which allows star formation in these halos, may partially cancel each other. Hence one may expect no significant deviations in the SFRs and reionization history if one treats these processes explicitly. To address the incident radiation and the resulting UV background from more rare density fluctuations outside of our simulation volume, it will be useful to bridge the gap between the start of reionization on Mpc scales to larger scale (10 -- 100 Mpc) simulations of reionization, such as the work of \citet{Sokasian03}, \citet{Iliev06}, \citet{Zahn07}, and \citet{Kohler07}. Radiation characteristics from a volume that has similar overdensities as our Mpc-scale simulations can be sampled from such larger volumes to create a radiation background that inflicts the structures in our Mpc scale simulations. Inversely, perhaps the small-scale evolution of the clumping factor, filtering mass, and average temperature and ionization states can be used to create an accurate subgrid model in large volume reionization simulations. Another potential caveat is the continued use of primordial gas chemistry in metal enriched regions in the SN runs. Our simulations with SNe give excellent initial conditions to self-consistently treating the transition to low-mass star formation. In future work, we plan to introduce metal-line and dust cooling models \citep[e.g. from][] {Glover07, Smith07} to study this transition. The one main assumption about Pop III stars in our calculations is the fixed, user-defined stellar mass. The initial mass function (IMF) of these stars is largely unknown, therefore we did not want to introduce an uncertainty by choosing a fiducial IMF. It is possible to calculate a rough estimate of the stellar mass by comparing the accretion rates and Kelvin-Helmholtz time of the contracting molecular cloud \citep{Abel02, OShea05}. Protostellar models of primordial stars have also shown that the zero-age main sequence (ZAMS) is reached at 100 \Ms~for typical accretion histories after the star halts its adiabatic contraction \citep{Omukai03, Yoshida06}. Furthermore, we have neglected HD cooling, which may become important in halos embedded in relic \ion{H}{2} regions and result in lower mass ($\sim$$30 \Ms$) metal-free stars \citep{OShea05, Greif06, Yoshida07b}. Based on accretion histories of star forming halos, one can estimate the ZAMS stellar mass for each halo and create a more self-consistent and ab initio treatment of Pop III star formation and feedback. We conducted three radiation hydrodynamical, adaptive mesh refinement simulations that supplement our previous cosmological simulations that focused on the hydrodynamics and cooling during early galaxy formation. These new simulations concentrated on the formation and feedback of massive, metal-free stars. We used adaptive ray tracing to accurately track the resulting \ion{H}{2} regions and followed the evolution of the photo-ionized and photo-heated IGM. We also explored on the details of early star formation in these simulations. Theories of early galaxy formation and reionization and large scale reionization simulations can benefit from the useful quantities and characteristics of the high redshift universe, such as SFR and IGM temperatures and ionization states, calculated in our simulations. The key results from this work are listed below. \medskip 1. SFRs increase from $5 \times 10^{-4}$ at redshift 30 to $6 \times 10^{-3}$ \sfr~at redshift 20 in our simulations. Afterwards the SFR begins to have a bursting nature in halos more massive than $10^7 \Ms$ and fluctuates around $10^{-2}$ \sfr. These rates are larger than the ones calculated in \citet{Hernquist03} because our simulation volume samples a highly biased region that contains a 2.5-$\sigma$ density fluctuation. The associated emissivity from these stars increase from 1 to $\sim$100 ionizing photons per baryon per Hubble time between redshifts 15 and 30. 2. In order to provide a comparison to semi-analytic models, we calculate the star formation efficiency to be $\sim$$6 \times 10^{-4}$ averaged over all redshifts and the simulation volume. For Pop III star formation, this is a factor of two lower than stars that are not affected by feedback \citep{Abel02, Bromm02, Yoshida06, OShea07}. 3. Shock waves created by D-type I-fronts expel most of the gas in the host halos below $\sim$$5 \times 10^6 \Ms$. Above this mass, significant outflows that are still bound to the halo are generated. This feedback creates a dynamical picture of early structure formation, where star formation is suppressed in halos because of this baryon depletion, which is more effective than UV heating or the radiative dissociation of \hh. 4. We see three instances of induced star formation in halos with masses $\sim 3 \times 10^6 \Ms$. Here a star forms as a SN blast wave overtakes an overdensity created by an ionization front instability. \hh~formation is catalyzed by additional free electrons in the relic \ion{H}{2} region and in the SN blast wave \citep{Ferrara98}. 5. As star formation occurs regularly in the simulation after redshift 25, four (six) ionizing photons are needed per sustained hydrogen ionization. As the most massive halo becomes larger than $\sim$$10^7 \Ms$ in the simulations without SNe, \ion{H}{2} regions become trapped and ionizing radiation only escapes into the IGM in small solid angles. Hence the number of photons per effective ionization increases to 15 (50). In SimB-SN, stellar radiation from induced star formation have an escape fraction of nearly unity, which occur four times in the calculation. This allows the IGM to remain ionized at a volume fraction 3 times higher than without SNe. Similarly, the ionizing photon to ionization ratio also stays elevated at 10:1 instead of decreasing in the calculations with star formation only. 6. Our simulations that include star formation and \hh~formation capture the entire evolution of the clumping factor that is used in semi-analytic models to calculate the effective enhancement of recombinations in the IGM. We showed that clumping factors in the ionized medium fluctuate around the 75\% of the values found in adiabatic simulations. They evolve from unity at high redshifts and steadily increase to $\sim$4 and 3.5 with and without SNe at $z = 17$, respectively. Photo-evaporation from stellar feedback causes the decrease of the clumping factor. 7. We calculated the Jeans filtering mass with the volume-averaged temperature only in fully and partially ionized regions, which yields a better estimate than the temperature averaged over both ionized and neutral regions. The filtering mass depends on the thermal history of the IGM, which mainly cools through Compton cooling. It increases by two orders of magnitude to $\sim$$3 \times 10^7 \Ms$ at $z \sim 15$. It describes the minimum mass a halo requires to collapse after hosting a Pop III star. For halos forming their first star, the minimum halo mass is regulated by the LW background \citep{Machacek01} and photo-evaporation \citep[e.g.][]{Haiman01}. \medskip Pop III stellar feedback plays a key role in early star formation and the beginning of cosmological reionization. The shallow potential wells of their host halos only amplify their radiative feedback. Our understanding of the formation of the oldest galaxies and the characteristics of isolated dwarf galaxies may benefit from including the earliest stars and their feedback in galaxy formation models. Although these massive stars only partially reionized the universe, their feedback on the IGM and galaxies is crucial to include since it affects the characteristics of low-mass galaxies that are thought to be primarily responsible for cosmological reionization. Harnessing observational clues about reionization, observations of local dwarf spheroidal galaxies, and numerical simulations that accurately handle star formation and feedback may provide great insight on the formation of the first galaxies, their properties, and how they completed cosmological reionization. | 7 | 10 | 0710.4328 |
0710 | 0710.4602_arXiv.txt | We use semi-analytic techniques to evaluate the burst sensitivity of designs for the {\it EXIST} hard X-ray survey mission. Applying these techniques to the mission design proposed for the Beyond Einstein program, we find that with its very large field-of-view and faint gamma-ray burst detection threshold, {\it EXIST} will detect and localize approximately two bursts per day, a large fraction of which may be at high redshift. We estimate that {\it EXIST}'s maximum sensitivity will be $\sim 4$~times greater than that of {\it Swift}'s Burst Alert Telescope. Bursts will be localized to better than 40~arcsec at threshold, with a burst position as good as a few arcsec for strong bursts. {\it EXIST}'s combination of three different detector systems will provide spectra from 3~keV to more than 10~MeV. Thus, EXIST will enable a major leap in the understanding of bursts, their evolution, environment, and utility as cosmological probes. | In its quest to find black holes throughout the universe, the {\it Energetic X-ray Imaging Survey Telescope (EXIST)} will detect, localize and study a large number of gamma-ray bursts, events thought to result from the birth of stellar-mass black holes. We present the methods used to calculate {\it EXIST}'s capabilities as a gamma-ray burst detector; we use the {\it EXIST} design evaluated by the National Research Council's `Committee on NASA's Beyond Einstein Program: An Architecture for Implementation' (2007; see also Grindlay 2007). The combination of large detector area, broad energy coverage, and wide field-of-view (FOV) will result in the detection of a substantial number of bursts with a flux distribution extending to fainter fluxes than that of previous missions. Thus {\it EXIST} should detect high redshift bursts, perhaps even bursts resulting from the death of Pop~III stars. {\it EXIST}'s imaging detectors will localize the bursts, while the combination of detectors, both imaging and non-imaging, will result in well-determined spectra from 3~keV to well over 10~MeV. In this paper we first describe the {\it EXIST} mission design (\S 2), emphasizing aspects relevant to burst detection. Then we present the sensitivity methodology (\S 3), which we apply to the individual coded mask sub-telescopes (\S 4). {\it EXIST} will consist of arrays of these detectors with overlapping FOVs, and the overall mission sensitivity results from adding the sensitivity of the individual sub-telescopes (\S 5). Imaging using counts accumulated over different timescales increases the sensitivity (\S 6). Finally, we combine these different calculations to evaluate {\it EXIST}'s overall capabilities to study bursts (\S 7). | From the preceding analysis, we can draw several conclusions on {\it EXIST}'s impact on the study of gamma-ray bursts. First we estimate the {\it EXIST} burst detection rate. The BATSE observations provide the cumulative burst rate as a function of the peak flux value $\psi_B$ averaged over $\Delta t=1$~s in the $\Delta E=$50--300~keV band (Band 2002): \begin{equation} N_B \sim 550 \left[ {{\psi_B}\over \hbox{0.3 ph cm$^{-2}$ s$^{-1}$}} \right]^{-0.8} \hbox{ bursts yr$^{-1}$ sky$^{-1}$ } \quad . \end{equation} The HET threshold sensitivity for a single sub-telescope on-axis is $\psi_B\sim 0.12$~ph~cm$^{-2}$~s$^{-1}$ for $E_p > 100$~keV. Using the BATSE rate in eq.~10 and integrating over the solid angle distribution in Figure~6 gives a burst detection rate for the HET of $\sim 400$ bursts per year. Note that this rate is over the BATSE-specific values of $\Delta E$ and $\Delta t$, and {\it EXIST} will use at least two different values of $\Delta E$ (see \S 4.1) and a variety of $\Delta t$ values (see \S 6). Consequently this rate should be increased by approximately 50\% to account for the soft, faint, long duration bursts to which BATSE was less sensitive than {\it EXIST}'s HET will be; we therefore expect the HET array to detect $\sim 600$~bursts per year. The value of $\psi_B$ for an LET varies more with the burst spectral parameters than for an HET, and therefore estimates of the LET burst detection rate based on the BATSE rate are much more uncertain. For a single LET $\psi_B\sim 0.3$~ph~cm$^{-2}$~s$^{-1}$ on axis at $E_p=100$~keV, which gives a burst detection rate of $\sim$180 bursts per year using eq.~10 and the LET distribution in Figure~6. This rate should be increased by a factor of 2 to account for the different energy band $\Delta E$ and accumulation times $\Delta t$. We use a larger adjustment factor for the LETs than for the HETs because the LETs' energy band will overlap less with BATSE's than the HETs'. We therefore expect the LET array to detect $\sim350$~bursts per year. Next we simulate the spectra that the {\it EXIST} suite of detectors will observe. Figure~9 shows a count spectrum (counts s$^{-1}$ keV$^{-1}$) for a moderately strong burst as it might be observed by the LETs (lefthand set of curves), HETs (middle set) and the CsI active shields for the HETs (righthand set; based on Garson et al. 2006a). The solid curves show the signal count rate, while the dashed curves provide the estimated background. Thus {\it EXIST} will facilitate spectral-temporal studies. Particularly important to physical burst emission models is determining $E_p$, which is typically of order 250~keV (Kaneko et al. 2006). In addition, correlations of $E_p$ with other burst properties, such as the `isotropic' energy (the Amati relation---Amati 2006) or total energy (the Ghirlanda relation---Ghirlanda, Ghisellini \& Lazzati 2004), have been proposed. `Pseudo-redshifts' calculated from the observables related to the burst-frame parameters in these relations can be used in burst studies when spectroscopic redshifts are not available, and can guide ground observers in allocating telescope time to observing potential high redshift bursts. The recently proposed Firmani relation (Firmani et al. 2006) correlates $E_p$, the peak luminosity, and a measure of the burst duration, all of which are related to observables in the gamma ray band. Thus pseudo-redshifts will be estimated using the Firmani relation based on {\it EXIST} data alone, independent of observations by other facilities. With well determined broadband spectra down to 3~keV, {\it EXIST} will be capable of determining whether the Band function (Band et al. 1993) suffices to describe burst spectra. For example, Preece et al. (1996) found evidence in the BATSE data for the presence of additional emission below 10~keV. By scaling from the EXIST survey's source localization (Grindlay 2007), we find that bursts should be localized at threshold by the HETs and LETs to better than 40~arcsec and 8~arcsec, respectively; this localization should scale as $\sim(\sigma-1)^{-1}$. Because the HETs are more sensitive to the LETs, the HET localization is relevant to the faintest bursts EXIST will detect. {\it EXIST}'s burst capabilities calculated above will constitute a major leap beyond current detectors, and should increase the number of high redshift bursts detected. On average, high redshift bursts should be fainter, softer and longer than low redshift bursts (although the broad burst luminosity function and great variety in burst lightcurves and spectra obscure this trend). Figure~10 compares the detector sensitivities of the HET (solid curve) and LET (dashed curve) arrays to the BAT on {\it Swift} (dot-dashed curve) and BATSE's Large Area Detector (LAD---dot-dot-dashed curve). As discussed above, the sensitivity is the threshold peak flux $F_T$ integrated over the 1--1000~keV band as a function of the spectrum's $E_p$; $\alpha=-1$ and $\beta=-2$ are assumed. In addition, the figure shows families of identical bursts at different redshifts (the curves with the points marked by `+'). Each family is defined by the value of $E_p$ in the burst frame; here again $\alpha=-1$ and $\beta=-2$ are assumed. In each family the burst would be observed to have $F_T$=7.5 ph cm$^{-2}$ s$^{-1}$ if it were at $z=1$. The points marked by `+' are spaced every $\Delta z=1/2$; thus the uppermost points are at $z=1$ and the lowermost points are at $z=10$. The pulses in burst lightcurves become narrower (shorter) at higher energy, an effect that is generally proportional to $E^{0.4}$ (Fenimore et al. 1995). Since the observed lightcurve originated in a higher energy band, pulses should become narrower with redshift, reducing the peak flux when integrated over a fixed accumulation time; the plotted families include this effect. Finally, in \S 6 we showed that forming images on long timescales increases the sensitivity to long duration bursts, as might result from cosmological time dilation. | 7 | 10 | 0710.4602 |
0710 | 0710.1840_arXiv.txt | {Detailed studies of Be stars in environments with different metallicities like the Magellanic Clouds or the Galactic bulge are necessary to understand the formation and evolution mechanisms of the circumstellar disks. However, a detailed study of Be stars in the direction of the bulge of our own galaxy has not been performed until now.} {The aim of this work is to report the first systematic search for Be star candidates in the direction of the Galactic Bulge. We present the full catalogue, give a brief description of the stellar variability seen, and show some light curve examples.} {We searched for stars matching specific criteria of magnitude, color and variability in the $I$ band. Our search was conducted on the 48 OGLE II fields of the Galactic Bulge. } {This search has resulted in 29053 Be star candidates, 198 of them showing periodic light variations. Nearly 1500 stars in this final sample are almost certainly Be stars, providing an ideal sample for spectroscopic multiobject follow-up studies. } {} | Be stars are non-supergiant fast rotator B stars whose spectra have, or had at some time, one or more Balmer lines in emission (Collins 1987). This emission is originated from a flattened circumstellar disk and can come and go episodically on time scales of days to decades. The responsible mechanisms for the production and dynamics of the circumstellar gas are still not constrained. Possible mechanisms include non-radial pulsations, wind-compressed disk model, magnetic activity and binarity (Porter $\&$ Rivinius 2003, and references therein).\\ Be stars are variable in brightness on three time scales that are often superimposed. Many of them (especially early type Be stars) show short-term photometric variability on time-scales of 0.2 to 2 days and amplitudes up 0.1 magnitudes, caused by non-radial pulsation or rotation (Percy et al. 2002, 2004). Some have mid-term variations on times scales form weeks to months, probably due to density waves within the disk (Sterken et al. 1996). Their amplitudes go to up to 0.2 magnitudes. They show also long-term variations from years to decades, with amplitudes up 0.8 magnitudes (Mennickent, Vogt, \& Sterken 1994; Pavlovski et al. 1997; Hubert $\&$ Floquet 1998; Percy $\&$ Bakos 2001). Stagg (1987) found that this type of variability occurres in at least half of the Be stars. A few of Be stars are close binaries and other present ejection process due to magnetic activity resulting in outbursts (Hubert et al. 1997).\\ Many Galactic Be stars have been surveyed for photometric variability in order to detect and confirm short-term or long-term variations and to find correlations between them and get clues on the physical processes in Be stars. For instance, Hubert $\&$ Floquet (1998) investigated the short-period variability of Be stars using analysis based on the Fourier and CLEAN algorithms on the Hipparcos photometry. Percy et al. (2002, 2004) analyzed a large sample of stars with Hipparcos photometry using a form of autocorrelation function. Typical problems in these studies are the gaps in the time distribution of the measurements and the limitations of the used algorithms.\\ In recent years, many Be-star like variables have been discovered in the Magellanic Clouds, showing a big variety of light curves, some of them reminding those of the Galactic Be stars and others never observed in that type of stars. Keller et al. (2002) concluded that most of these blue variables should be Be stars. Searches for Be stars in the Magellanic Clouds were performed by Mennickent et al. (2002), Keller et al. (1999, 2002) and Sabogal et al. (2005) on the basis of selection criteria applied to different photometric databases (OGLE-II and MACHO), and took into account amplitude of variability and ranges of color-magnitudes in the selection process. De Wit et al. (2006) investigated a subsample of the blue variables found by Mennickent et al. (2002) in the Small Magellanic Cloud and found that the photometric variability of these Be stars is due to variations in the amount of Brehmstrahlung due to the evolution of the circumstellar gas from a disk-shaped envelope towards a ring-like structure. \\ The study of Be stars is relevant to make contributions to several important branches of stellar physics. In particular, detailed studies of Be stars in environments with different metallicities like the Magellanic Clouds or the Galactic bulge is crucial to understand the formation and evolution mechanisms of the circumstellar disks. However, a detailed study of Be stars in the direction of the bulge of our own galaxy has not been performed until now.\\ A very large number of stars was observed in the region of the Galactic Bulge during the course of the second phase of the Optical Gravitational Lensing Experiment (OGLE II) (Udalski, Kubiak \& Szymanski 1997). We have performed a search for Be star candidates into this database. Here we present the results of this search. | In this paper we have provided a catalogue of 29053 Be star candidates, 198 of them periodic, in the direction of the Galactic bulge that were selected basically on photometric criteria. Most of these Be star candidates are probably members of the Galactic disk and trace the gradient of metallicity towards the Galactic centre. They are ideal targets for future observing programs based on multiobject spectroscopy, narrow band photometry or H$\alpha$ imaging surveys. These programs could eventually establish their true nature and broke the residual degeneracy with variable red giants in the red part of the (V-I) color distribution. | 7 | 10 | 0710.1840 |
0710 | 0710.4524_arXiv.txt | It is well know that the coronagraphic observations of halo CMEs are subject to projection effects. Viewing in the plane of the sky does not allow us to determine the crucial parameters defining geoeffectivness of CMEs, such as the velocity, width or source location. We assume that halo CMEs at the beginning phase of propagation have constant velocities, are symmetric and propagate with constant angular widths. Using these approximations and determining projected velocities and difference between times when CME appears on the opposite sides of the occulting disk we are able to get the necessary parameters. We present consideration for the whole halo CMEs from SOHO/LASCO catalog until the end of 2000. We show that the halo CMEs are in average much more faster and wider than the all CMEs from the SOHO/LASCO catalog. | Space Weather is significantly controlled by coronal mass ejections (CMEs) which can affect the Earth in a different way. CMEs originating close to the central meridian, directed toward the Earth, excite the biggest scientific concern. In coronagraphic observations they appear as enhancement surrounding the entire occulting disk and they were called `halo CME'. Since the first identification by Howard et al. (1982) plenty of them were detected and now they are routinely recorded by the high sensitive SOHO/LASCO coronagraphs. In spite of large advantage over previous instruments, the SOHO/LASCO observations are still affected by a projection effect (Gopalswamy et al. 2000b). Viewing in the plane of the sky does not allow us to determine the crucial parameters defining geoeffectivness of CMEs, such as the velocity, width or source location. Prediction of the arrival of CME in the vicinity of Earth is critically important in space weather investigations. Basing on interplanetary shocks detected by Wind and the corresponding CMEs detected by SOHO, Gopalswamy et al. (2000a) developed and next (Gopalswamy, 2001) improved an empirical model to predict the arrival of CMEs at 1AU. The critical element affecting this model is the initial CME speed. The better prediction could be achieved if real initial velocities are used instead projected velocities determined from LASCO observations. Similarly, attempts made to estimate the projection effect based on the location of the solar source employ ad hoc assumptions on parameter such as the width of CMEs (Sheeley et al. 1999, Leblanc and Dulk, 2000). In the present paper we try to determine these crucial parameters defining geoeffectivness of CMEs, such as the velocity, width or source location. We assume that halo CMEs at the beginning phase of propagation have a constant velocities, are symmetric and propagate with constant angular widths. Using these approximations and determining projected velocities and difference between times when CME appears on the opposite sides of the occulting disk we are able to get necessary parameters. We present results for the whole halo CMEs from SOHO/LASCO catalog until the and of 2000. | In this paper we present possibility to estimate the crucial parameters determining geoeffectiveness of the halo CMEs. The clue of this method is based on the difference between times when the halo CME appears at the opposite sides (first and finally appearance) of the occulting disk. We considered the whole events form SOHO/LASCO catalog until the end of 2000. We were able to determine the real velocity, width and source location for 73th CMEs from our sample. Unfortunately, 58 events were symmetric or too faint to do necessary considerations. Results are listed in the three successive tables. This list could be use for further statistical examination or to prediction of the arrival of CME in the vicinity of Earth. Presented results suggest that the halo CMEs represent a specific class of CMEs which are very wide and fast. Using our results we have to remember that the simple model has several shortcomings: (i) CMEs may be accelerating, moving with constant speed or decelerating at the beginning phase of propagation. This means the constant velocity we assumed may not hold. (ii) CMEs may expand in addition to radial motion. Then the measured sky-plane speed is a sum of the expansion speed and the projected radial speed. This also would imply that the CMEs may not be a rigid cone as we assumed (Gopalswamy et al. 2001) (iii) The cone symmetry also may not hold. CME originating from loop structure could be elongated. All these limits can be overcome by stereoscopic observations only. Unfortunately, at the present time they are not available yet. It is necessary to develop the model to get the better fit to observations. The first step to improve our model could be achieved by consideration of acceleration and expansion of CMEs. \\ \\ {\small \bf Acknowledgments}\\ {\scriptsize This paper was done during work of Grzegorz Michalek at Center for Solar and Space Weather, Catholic University of America in Washington.\\ In this paper we used data from SOHO/LASCO CME catalog. This CME catalog is generated and maintained by the Center for Solar Physics and Space Weather, The Catholic University of America in cooperation with the Naval Research Laboratory and NASA. SOHO is a project of international cooperation between ESA and NASA.\\ Work done by Grzegorz Michalek was partly supported by {\it Komitet Bada\'{n} Naukowych} through the grant PB 258/P03/99/17.} \vspace{10mm} | 7 | 10 | 0710.4524 |
0710 | 0710.4238_arXiv.txt | {Since most high- and intermediate-mass protostars are at great distance and form in clusters, high linear resolution observations are needed to investigate their physical properties.} {To study the gas in the innermost region around the protostars in the proto-cluster IRAS\,05358+3543, we observed the source in several transitions of methanol and other molecular species with the Plateau de Bure Interferometer and the Submillimeter Array, reaching a linear resolution of 1100~AU.} {We determine the kinetic temperature of the gas around the protostars through an LVG and LTE analysis of their molecular emission; the column densities of CH$_3$OH, CH$_3$CN and SO$_2$ are also derived. Constrains on the density of the gas are estimated for two of the protostellar cores.} {We find that the dust condensations are in various evolutionary stages. The powerhouse of the cluster, mm1a, harbours a hot core with $T\sim 220~(75<T<330)$~K. A double-peaked profile is detected in several transitions toward mm1a, and we found a velocity gradient along a linear structure which could be perpendicular to one of the outflows from the vicinity of mm1a. Since the size of the double-peaked emission is less than 1100~AU, we suggest that mm1a might host a massive circumstellar disk. The other sources are in earlier stages of star formation. The least active source, mm3, could be a starless massive core, since it is cold ($T<20$~K), with a large reservoir of accreting material ($M\sim 19~M_\odot$), but no molecular emission peaks on it.} {} | The last decade has seen significant progress in the understanding of how high-mass stars form. Large samples of massive young stellar objects (YSOs) were studied with single-dish telescopes, to investigate their physical properties through the analysis of their (sub)mm continuum and molecular emission \citep[e.g.,][]{1996A&A...308..573M,1998A&A...336..339M,2000A&A...355..617M,1997MNRAS.291..261W,1998MNRAS.301..640W,1999MNRAS.309..905W,2000A&A...357..637H,2001ApJ...552L.167Z,2005ApJ...625..864Z,2002ApJ...566..931S,2002ApJ...566..945B,2002A&A...383..892B,2004A&A...426...97F,2004A&A...417..115W,2005A&A...434..257W}. However, an intrinsic feature of high-mass stars is that they form in clusters, and that most of them are at large (several kpc) distances. Therefore, single-dish studies, as valuable as they are, lack the necessary spatial resolution to resolve single protostars and study the inner regions where high-mass star formation takes place. Interferometric observations started shedding light into the complex nature of high-mass star forming regions with the adequate spatial resolution \citep[e.g.,][]{1997A&A...325..725C,1999A&A...345..949C,2005A&A...434.1039C,1999ApJ...514L..43W,2002A&A...387..931B,2005ApJ...628..800B,2006ApJ...649..888H,2000ApJ...535..833S}. However, the number of massive YSOs studied at high resolution is still too small to establish the general properties of the dense cores where massive stars form on a statistical base. In this paper, we present an interferometric analysis of the high-mass star forming region \object{IRAS\,05358+3543} at (sub)mm wavelengths in several molecular transitions. IRAS\,05358+3543 (also known in literature as \object{S233IR}) is part of a sample of 69 high-mass protostellar objects studied in great detail in recent years \citep{2002ApJ...566..931S,2002ApJ...566..945B,2002A&A...383..892B,2002A&A...390..289B,2004A&A...417..115W,2005A&A...434..257W,2005A&A...442..949F}. At a distance of 1.8~kpc \citep{1990ApJ...352..139S}, IRAS 05358+3543 has a bolometric luminosity of 6300~L$_{\sun}$; strong high-mass star formation activity is evidenced by maser emission \citep[see][]{1991ApJ...380L..75M,1995A&AS...112..299T,2000A&A...362.1093M} and outflow activity \citep{1990ApJ...352..139S,2002A&A...387..931B}. Previous interferometric observations by \citet{2002A&A...387..931B} resolved three dust condensations (mm1, mm2 and mm3) within an area of $9'' \times 4''$ ($\sim17\,100\times7\,200$~AU), and revealed at least three outflows in CO and SiO, the most prominent of which is more than a parsec in length, and massive $(M >10~M_ {\sun})$. Two of the three identified outflows originate from the vicinity of mm1, which is probably the main powerhouse in the region. To zoom in on the innermost region around the protostars, and study the physical properties of the individual potentially star-forming cores, we carried out a comprehensive program to observe the region at high spatial resolution with the Plateau de Bure Interferometer\footnote{IRAM is supported by INSU/CNRS (France), MPG (Germany) and IGN (Spain).} at 97~GHz and 241~GHz, and 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.} at 338~GHz. The new observations reach a resolution down to ~$0.6''$, corresponding to $\sim 1100$~AU at the distance of the source. \citet{05358-beuther} studied the continuum emission of this dataset. They identified four compact protostellar sources in the region; mm1 is resolved into two continuum peaks, mm1a and mm1b, with a projected linear separation of $\sim 1700$~AU. A mid-infrared source \citep{2006MNRAS.369.1196L}, and a compact 3.6~cm continuum source \citep{05358-beuther} coincide with mm1a, which is also associated with the class II methanol masers detected by \citet{2000A&A...362.1093M}. The previously identified source mm2 resolves into several sub-sources; however, only one of them (mm2a, according to the nomenclature used by \citealt{05358-beuther}) is a protostellar source, while the others are probably caused by the outflows in the region. The third source mm3 remains a single compact core even at the highest spatial resolution. In Table~\ref{cores}, we report the positions of the four sources identified in the continuum emission by \citet{05358-beuther}. In this paper, we discuss the spectral line observations complementing the continuum data discussed by \citet{05358-beuther}. In section \S\ref{observations}, the different observations are presented. In section \S\ref{obs-res}, we discuss our results, and analyse the extended emission of low excitation molecular transitions (\S\ref{extended}), as well as the molecular spectra at the positions of the dust condensations(\S\ref{proto}). Finally, in section \S\ref{analysis} we derive the physical parameters of the gas around the protostars from the analysis of their spectra. In the following sections, we use the term protostar for young massive stellar objects which are still accreting material from the surroundings, independently whether they already started burning hydrogen or not. \begin{table} \centering \caption{Positions of the four dust condensations in IRAS\,05358+3543 \citep[from][]{05358-beuther}.\label{cores}} \begin{tabular}{lcc} \multicolumn{1}{c}{Source} &\multicolumn{1}{c}{R.A. [J2000]}&\multicolumn{1}{c}{Dec. [J2000]}\\ \hline \hline mm1a&05:39:13.08&35:45:51.3\\ mm1b&05:39:13.13&35:45:50.8\\ mm2a&05:39:12.76&35:45:51.3\\ mm3&05:39:12.50&35:45:54.9\\ \hline \hline \end{tabular} \end{table} | Our new interferometric data resolve at least four cores in the high-mass protocluster IRAS\,05358+3543. By analysing the molecular spectrum of each condensations, we characterised the properties of the gas surrounding it. Our main results are summarised in the following: \begin{itemize} \item the main powerhouse of the region, mm1a, harbours a hot core with $T\sim 220$~K, and the central heating source has a chemical timescale of $10^{4.6}$~yr. Our data suggest that mm1a might host a massive circumstellar disk; \item although the properties of the mm continuum emission of mm1b are very similar to the one of mm1a, the two sources differ significantly in the cm and mid-infrared spectrum. This core could be in an earlier stage of star formation than mm1a, since no molecular emission is detected toward it, with the only exception of the $5 \to 4$~C$^{34}$S transition; \item given the low abundance of methanol, mm2 could be a low--intermediate mass protostar; \item strong emission is detected in several molecular species to the north-west of mm2, at a position where no continuum emission is detected. We suggest that this is caused by the interaction of the outflows with the ambient molecular cloud; \item the least active source, mm3, could be a starless massive core, since it is cold ($T<20$~K), with a large reservoir of accreting material ($M\sim 19~M_\odot$), but no molecular emission peaks on it. \end{itemize} | 7 | 10 | 0710.4238 |
0710 | 0710.5714_arXiv.txt | Fusion reactions in the crust of an accreting neutron star are an important source of heat, and the depth at which these reactions occur is important for determining the temperature profile of the star. Fusion reactions depend strongly on the nuclear charge $Z$. Nuclei with $Z\le 6$ can fuse at low densities in a liquid ocean. However, nuclei with $Z=8$ or 10 may not burn until higher densities where the crust is solid and electron capture has made the nuclei neutron rich. We calculate the $S$ factor for fusion reactions of neutron rich nuclei including $^{24}$O + $^{24}$O and $^{28}$Ne + $^{28}$Ne. We use a simple barrier penetration model. The $S$ factor could be further enhanced by dynamical effects involving the neutron rich skin. This possible enhancement in $S$ should be studied in the laboratory with neutron rich radioactive beams. We model the structure of the crust with molecular dynamics simulations. We find that the crust of accreting neutron stars may contain micro-crystals or regions of phase separation. Nevertheless, the screening factors that we determine for the enhancement of the rate of thermonuclear reactions are insensitive to these features. Finally, we calculate the rate of thermonuclear $^{24}$O + $^{24}$O fusion and find that $^{24}$O should burn at densities near $10^{11}$ g/cm$^3$. The energy released from this and similar reactions may be important for the temperature profile of the star. | Nuclei accreting onto a neutron star undergo a variety of reactions. First at low densities, conventional thermonuclear fusion takes place, see for example \cite{rpash}. Next as nuclei are buried to higher densities, the rising electron Fermi energy induces a series of electron captures \cite{gupta}. Finally at very high densities, nuclei can fuse via pycnonuclear reactions. These reactions are induced by the quantum zero point motion \cite{pycno}. The energy released, and the densities at which reactions occur, are important for determining the temperature profile of neutron star crusts. Superbursts are very energetic X-ray bursts from accreting neutron stars that are thought to involve the unstable thermonuclear burning of carbon \cite{superbursts, superbursts2}. However, some simulations do not reproduce the conditions needed for carbon ignition because they have too low temperatures \cite{superignition}. An additional heat source, from fusion or other reactions, could raise the temperature and allow carbon ignition at densities that reproduce observed burst frequencies. Recently the cooling of two neutron stars has been observed after extended outbursts \cite{Wijnands, cackett}. These outbursts heated the crusts out of equilibrium and then the cooling time was measured as the crusts returned to equilibrium. The surface temperature of the neutron star in KS 1731-260 decreased with an exponential time scale of 325 $\pm$ 100 days while MXB 1659-29 has a time scale of 505 $\pm$ 59 days \cite{cackett}. These cooling times depend on the thermal conductivity of the crust and the initial temperature profile. Comparing these observations of relatively rapid cooling to calculations by Rutledge et al. \cite{rutledge} and Shternin et al. \cite{shternin} suggests that the crust has a high thermal conductivity. However, if the initial temperature profile of the crust is peaked near the surface, then this peak could quickly diffuse to the surface and lead to rapid cooling. Therefore, cooling time scales are also sensitive to the initial temperature profile, and this depends on heating from nuclear reactions at moderate densities in the crust. Gupta et al. have calculated heating from electron capture reactions in the outer crust \cite{gupta}. While they find more heating than previous works, they still find no more than 0.4 MeV per nucleon total heating from all of the electron captures on any mass number $A$ system. Haensel and Zdunik have calculated pycnonuclear fusion reactions at great densities in the inner crust \cite{haensel}. However, if reactions occur deep in the inner crust, most of the heat may flow in to the core instead of out towards the surface. As a result, there may be a smaller impact on the temperature profile of the outer crust. A low crust thermal conductivity, for example from an amorphous solid, could help explain superburst ignition. This could better insulate the outer crust and allow higher carbon ignition temperatures. However, a low thermal conductivity appears to be directly contradicted by the observed short crust cooling times. Furthermore, our molecular dynamics simulations in ref. \cite{horowitz} and further results we present in Section \ref{MD} find a regular crystal structure, even when the system has a complex composition with many impurities. We do not find an amorphous phase. These results will be discussed further in a later publication. We conclude that the thermal conductivity of the crust is high. If the thermal conductivity is high, one may need additional heat sources, at moderate densities, in order to explain superburst ignition. Although Gupta et al. find additional heating from electron captures to excited nuclear states, simple nuclear structure properties may provide a natural limit to the total heating from electron captures \cite{brownprivate}. Haensel and Zdunik \cite{haensel,haensel2007} consider heating from pycnonuclear reactions using a simple one component plasma model. They find that fusion reactions may not take place until relatively high densities above $10^{12}$ g/cm$^3$. However, their use of a one component plasma could be a significant limitation. Fusion reactions depend strongly on the nuclear charge $Z$. Therefore, the reaction rate may be highest for the rare impurities that have the lowest $Z$, instead of for nuclei of average charge. In this paper, we go beyond Haensel and Zdunik and consider a full mixture of complex composition instead of assuming one average charge and mass. We focus on thermonuclear and pycnonuclear reactions at densities around $10^{11}$ g/cm$^3$. This is near the base of the outer crust. Heat released at this density could be important for superburst ignition and for crust cooling times. Nuclei at this density are expected to be neutron rich. Furthermore, the other nearby ions strongly screen the Coulomb barrier and greatly enhance the rate of thermonuclear reactions. We begin by describing the initial composition. This includes neutron rich light nuclei such as $^{24}$O and $^{28}$Ne. We calculate cross sections and $S$ factors for $^{24}$O + $^{24}$O and $^{28}$Ne + $^{28}$Ne fusion using a simple barrier penetration model. Note that the dynamics of the neutron rich skins of these nuclei can enhance the cross section over that predicted by our simple barrier penetration model. This is a very interesting and open nuclear structure question, see for example \cite{subbarrier}. Next, we use classical molecular dynamics simulations to determine the structure of the crust and screening factors for the enhancement of thermonuclear reactions. There are many previous calculations of screening factors for the one component plasma \cite{ocpscreening} and for binary ion mixtures, see for example \cite{bimscreening}. However, we are not aware of any previous calculations for a crystal of a complex multicomponent composition. Finally, we calculate reaction rates and conclude that $^{24}$O is expected to fuse at densities near $10^{11}$ g/cm$^3$ while $^{28}$Ne should react at densities near $10^{12}$ g/cm$^3$. Heat from these reactions may be important for determining the temperature profile of accreting neutron stars. | \label{summary} Fusion reactions in the crust of an accreting neutron star are an important source of heat, and the depth at which these reactions occur is important for determining the temperature profile of the star. Fusion reactions depend strongly on the nuclear charge $Z$. Nuclei with $Z\le 6$ can fuse at low densities in a liquid ocean. However, nuclei with $Z=8$ or 10 may not burn until higher densities where the crust is solid and electron capture has made the nuclei neutron rich. In Section \ref{crosssections} we calculated the $S$ factor for fusion reactions of neutron rich nuclei including $^{24}$O + $^{24}$O and $^{28}$Ne + $^{28}$Ne. We used a simple barrier penetration model. We find that $S$ for $^{24}$O+$^{24}$O is over eight orders of magnitude larger than that for $^{16}$O+$^{16}$O. The $S$ factor could be further enhanced by dynamical effects involving the neutron rich skin of $^{24}$O. For example, the skins of the two nuclei could deform to form a neck that would reduce the Coulomb barrier. This possible enhancement in $S$ should be studied in the laboratory with neutron rich radioactive beams. In Section \ref{MD} we modeled the structure of the crust with molecular dynamics simulations. We find that the crust of accreting neutron stars may contain micro-crystals or regions of phase separation. Nevertheless, the screening factors that we determined for the enhancement of the rate of thermonuclear reactions are insensitive to these features. Finally, we calculated in Section \ref{reaction rates} the rate of thermonuclear $^{24}$O + $^{24}$O fusion and find that $^{24}$O should burn at densities near $10^{11}$ g/cm$^3$. This is a lower density than some previous estimates. The 0.52 MeV per nucleon energy released may be important for the temperature profile of the star. In future work, we will use our molecular dynamics results to study other properties of the crust such as its thermal conductivity. In addition, we will use these MD results to calculate pycnonuclear reaction rates for the fusion of $^{28}$Ne and other heavier nuclei. | 7 | 10 | 0710.5714 |
0710 | 0710.2628_arXiv.txt | We observed four southern AXPs in 1999 near 1400 MHz with the Parkes 64-m radio telescope to search for periodic radio emission. No Fourier candidates were discovered in the initial analysis, but the recent radio activity observed for the AXP XTE J1810$-$197 has prompted us to revisit these data to search for single radio pulses and bursts. The data were searched for both persistent and bursting radio emission at a wide range of dispersion measures, but no detections of either kind were made. These results further weaken the proposed link between rotating radio transient sources and magnetars. However, continued radio searches of these and other AXPs at different epochs are warranted given the transient nature of the radio emission seen from XTE J1810$-$197, which until very recently was the only known radio-emitting AXP. | The detection of pulsed radio emission from the anomalous X-ray pulsar (AXP) XTE J1810$-$197 in 2006 \cite{crh+06}, and more recently from a second AXP, 1E 1547.0$-$5408 \cite{crh+07}, has renewed interest in searching for radio emission from these objects. In both of these cases, the radio activity is believed to be connected to the X-ray variability of the sources and is transient in nature (or at least highly variable). Given this transient behavior and that both persistent periodic emission and single pulses were detected from both AXPs, renewed searches of archival radio search data of AXPs at different epochs may reveal previously undetected radio signals from these sources. | We found no convincing radio signals in either the folding or single pulse searches. The derived upper limits on the radio emission from our AXP targets are presented in Table \ref{tbl-1}. These are the most stringent radio upper limits to date for these sources. The estimated 1400 MHz luminosity limits on the periodic radio emission ($\approxlt 1$ mJy kpc$^{2}$) are 2-3 times lower than those established for XTE J1810$-$197 prior to outburst. However, it is still conceivable that weak radio pulses are being emitted, but that they are below our detection threshold. The luminosity limits presented here for the periodic emission are lower by about two orders of magnitude than the 1400 MHz luminosity of the periodic pulsed radio emission from XTE J1810$-$197 soon after the radio emission was first detected ($\sim 80$ mJy kpc$^{2}$) \citep{crh+06}. Thus we would expect to be able to easily detect comparably strong radio emission if it were beamed toward us. Our luminosity limits on single radio pulses from our sources range from 22 to 69 Jy kpc$^{2}$ in the most conservative case, which is below the 1400 MHz luminosity of $\approxgt 100$ Jy kpc$^{2}$ derived from the pulse strengths reported for XTE J1810$-$197 in its radio discovery paper \citep{crh+06}. Since single pulses were detected from almost every rotation of XTE J1810$-$197, it is likely that we would have detected a large number of comparable pulses during our observations if such pulses were beamed toward us. Our non-detection of single pulses further weakens the hypothesis that rotating radio transients (RRATs) \citep{mlm+06} and magnetars are linked. This has been weakened by two other recent results. First, the X-ray detection of the RRAT J1819$-$1458 shows that its emission is more typical of middle-aged pulsars than it is of magnetars \citep{rbg+06}. Second, the nearby, rotation-powered pulsar PSR B0656+54 would probably have been identified as an RRAT if it were farther away \citep{wsr+06}. We conclude from our results that any periodic or bursting radio emission from the four target AXPs is either very weak (below our detection thresholds), not beamed toward us, or non-existent or sporadic at the epoch of observation. This last possibility is suggested by the connection between the X-ray and radio activity observed for the two known radio-emitting magnetars to date. Continued radio searches of AXPs are therefore warranted given the apparent transient nature of the radio emission. Further details of this work and a more complete discussion of the results are presented in a recent journal article \cite{chk07}. \begin{table} \begin{tabular}{lcccc} \hline & \tablehead{1}{r}{b}{1E 1048.1$-$5937} & \tablehead{1}{r}{b}{AX J1845$-$0258\tablenote{AXP candidate only}} & \tablehead{1}{r}{b}{1E 1841$-$045} & \tablehead{1}{r}{b}{1RXS J170849.0$-$400910} \\ \hline Spin period (s) & 6.45 & 6.97 & 11.77 & 11.00 \\ Ephemeris reference & \cite{kgc+01} & \cite{tkk+98}\tablenote{No period derivative available} & \cite{gvd99} & \cite{gk02} \\ Galactic longitude, latitude (deg) & 288.26, $-$0.52 & 29.52, 0.07 & 27.39, $-$0.01 & 346.47, 0.03 \\ $T_{\rm sky}$ (K)\tablenote{1374 MHz sky temperature estimated from \cite{hss+82} assuming a spectral index of $-2.6$} & 9.1 & 12.3 & 13.2 & 16.3 \\ Observation MJD & 51378 & 51391 & 51382 & 51379 \\ Observation date & 1999 Jul 19 & 1999 Aug 1 & 1999 Jul 23 & 1999 Jul 20 \\ $S_{1400}$ (mJy)\tablenote{1400 MHz flux density limit on pulsed emission estimated using the modified radiometer equation and an assumed duty cycle of 2.7\%} & $\approxlt 0.02$ & $\approxlt 0.02$ & $\approxlt 0.02$ & $\approxlt 0.02$ \\ $S_{1400}$ single (mJy)\tablenote{Range of single-pulse 1400 MHz flux limits for pulse time-scales 0.25-240 ms} & $\approxlt 875$-50 & $\approxlt 975$-60 & $\approxlt 1000$-60 & $\approxlt 1085$-65 \\ Distance (kpc)\tablenote{Taken from \cite{bri+06}. Question marks indicate significant uncertainty in the value} & $\sim 5$? & $\sim 8$? & $\sim 7$ & $\sim 8$? \\ $L_{1400}$ (mJy kpc$^{2}$)\tablenote{1400 MHz luminosity limit on pulsed emission, assuming a 1 sr beaming fraction} & $\approxlt 0.5$ & $\approxlt 1.3$ & $\approxlt 1.0$ & $\approxlt 1.3$ \\ $L_{1400}$ single (Jy kpc$^{2}$)\tablenote{Range of 1400 MHz luminosity limits on single pulses for pulse time-scales 0.25-240 ms} & $\approxlt 22$-1.3 & $\approxlt 62$-3.7 & $\approxlt 49$-2.9 & $\approxlt 69$-4.1 \\ \hline \end{tabular} \caption{Radio Search Parameters and Results} \label{tbl-1} \end{table} | 7 | 10 | 0710.2628 |
0710 | 0710.0288_arXiv.txt | The channeling effect of low energy ions along the crystallographic axes and planes of NaI(Tl) crystals is discussed in the framework of corollary investigations on WIMP Dark Matter candidates. In fact, the modeling of this existing effect implies a more complex evaluation of the luminosity yield for low energy recoiling Na and I ions. In the present paper related phenomenological arguments are developed and possible implications are discussed at some extent. | It is known that ions (and, thus, also recoiling nuclei) move in a crystal in a different way than in amorphous materials. In particular, ions moving (quasi-) parallel to crystallographic axes or planes feel the so-called ``channeling effect'' and show an anomalous deep penetration into the lattice of the crystal \cite{chan,nel63,oth2}; see Fig. \ref{fg:schema_chan}. For example, already on 1957, a penetration of $^{134}$Cs$^+$ ions into a Ge crystal was observed to a depth of about 1000 \AA \, \cite{Bre56}, larger than that expected in the case the ions would cross amorphous Ge ($\simeq 50$ \AA). Afterwards, high intensities of H$^+$ ions at 75 keV transmitted through thick (3000-4000 \AA) single-crystal gold films in the $<110>$ directions were detected \cite{nel63}. Other examples for keV range ions have been shown in ref. \cite{Ras05} where 3 keV P$^+$ ions moving into layers of 500 \AA \, of various crystals were studied. The channeling effect is also exploited in high energy Physics e.g. to extract high energy ions from a beam by means of bent crystals or to study diffractive Physics by analysing scattered ions along the beam direction (see e.g. ref. \cite{Baur00}). Recently \cite{droby2} it has been pointed out the possible role which this effect can play in the evaluation of the detected energy of recoiling nuclei in crystals, such as the NaI(Tl)\footnote{For completeness, it is worth to note that luminescent response for channeling in NaI(Tl) was already studied in ref. \cite{lunt} for MeV-range ions.}. \begin{figure} [!t] \centering \includegraphics[width=9.0cm] {fig1.eps} \caption{Simplified schema of the channeling effect in the NaI(Tl) lattice. The axial channeling occurs when the angle of the motion direction of an ion with the respect to the crystallographic axis is less than a characteristic angle, $\Psi_c$, depicted there (see for details Sec. 2). Two examples for channeled and unchanneled ions are also shown (dashed lines).} \label{fg:schema_chan} \end{figure} In fact, the channeling effect can occur in crystalline materials due to correlated collisions of ions with target atoms. In particular, the ions through the open channels have ranges much larger than the maximum range they would have if their motion would be either in other directions or in amorphous materials. Moreover, when a low-energy ion goes into a channel, its energy losses are mainly due to the electronic contributions. This implies that a channeled ion transfers its energy mainly to electrons rather than to the nuclei in the lattice and, thus, its quenching factor (namely the ratio between the detected energy in keV electron equivalent [keVee] and the kinetic energy of the recoiling nucleus in keV) approaches the unity. It is worth to note that this fact can have a role in corollary analyses in the Dark Matter particle direct detection experiments, when WIMP (or WIMP-like) candidates are considered. In fact, since the routine calibrations of the detectors are usually performed by using $\gamma$ sources (in order to avoid induced radioactivity in the materials), the quenching factor is a key quantity to derive the energy of the recoiling nucleus after an elastic scattering. Generally, for scintillation and ionization detectors this factor has been inferred so far by inducing tagged recoil nuclei through neutron elastic scatterings \cite{Mis_neut}; however, as it will be discussed in Sec. 3, the usual analysis carried out on similar measurements does not allow to account for the channeled events. A list of similar values for various nuclei in different detectors can be found e.g. in ref. \cite{RNC}. In particular, commonly in the interpretation of the dark matter direct detection results in terms of WIMP (or WIMP-like) candidates the quenching factors are assumed to be constant values without considering e.g. their energy dependence, the properties of each specific used detector and the experimental uncertainties. An exception was in the DAMA/NaI corollary model dependent analyses for WIMP (or WIMP-like) candidates \cite{RNC,ijmd,epj06,ijma2} where at least some of the existing uncertainties on the $q_{Na}$ and $q_I $ values, measured with neutrons, were included. In this paper the possible impact of the channeling effect in NaI(Tl) crystals is discussed in a phenomenological framework and comparisons on some of the corollary analyses carried out in terms of WIMP (or WIMP-like) candidates \cite{RNC,ijmd,epj06,ijma2}, on the basis of the 6.3 $\sigma$ C.L. DAMA/NaI model independent evidence for particle Dark Matter in the galactic halo\footnote{We remind that various possibilities for some of the many possible astrophysical, nuclear and particle Physics scenarios have have been analysed by DAMA itself both for some WIMP/WIMP-like candidates and for light bosons \cite{RNC,ijmd,epj06,ijma2,ijma}, while other corollary analyses are also available in literature, such as e.g. refs. \cite{Bo03,Bo04,Botdm,khlopov,Wei01,foot,Saib,droby1,droby2}. Many other scenarios can be considered as well.}, are given. | In this paper the channeling effect of recoiling nuclei induced by WIMP and WIMP-like elastic scatterings in NaI(Tl) crystals has been discussed. Its possible effect in a reasonably cautious modeling has been presented as applied to some given simplified scenarios in corollary quests for the candidate particle for the DAMA/NaI model independent evidence. This further shows the role of the existing uncertainties and of the correct description and modeling of all the involved processes as well as their possible impact in the investigation of the candidate particle. Some of them have already been addressed at some extent, such as the halo modeling \cite{halo,RNC,ijmd}, the possible presence of non-thermalized components in the halo (e.g. caustics \cite{siki} or SagDEG \cite{epj06} contributions), the accounting for the electromagnetic contribution to the WIMP (or WIMP-like) expected energy distribution \cite{ijma2}, candidates other than WIMPs (e.g. \cite{ijma} and in literature), etc.. Obviously, many other arguments can be addressed as well both on DM candidate particles and on astrophysical, nuclear and particle physics aspects; for more see \cite{RNC,ijmd,ijma,epj06,ijma2} and in literature. In particular, we remind that different astrophysical, nuclear and particle Physics scenarios as well as the experimental and theoretical associated uncertainties leave very large space also e.g. for significantly lower cross sections and larger masses. | 7 | 10 | 0710.0288 |
0710 | 0710.0846_arXiv.txt | s{ The Alpha Magnetic Spectrometer (AMS), to be installed on the International Space Station (ISS) in 2008, is a cosmic ray detector with several subsystems, one of which is a proximity focusing Ring Imaging \CK\ (RICH) detector. This detector will be equipped with a dual radiator (aerogel+NaF), a lateral conical mirror and a detection plane made of 680 photomultipliers and light guides, enabling precise measurements of particle electric charge and velocity. Combining velocity measurements with data on particle rigidity from the AMS Tracker it is possible to obtain a measurement for particle mass, allowing the separation of isotopes. \\ A Monte Carlo simulation of the RICH detector, based on realistic properties measured at ion beam tests, was performed to evaluate isotope separation capabilities. Results for three elements --- H (Z=1), He (Z=2) and Be (Z=4) --- are presented. } \vspace{-0.9cm} | AMS will provide a major improvement on existing data for isotopic abundances in cosmic rays. Simulation results indicate that the separation of light isotopes using the combination of RICH data and tracker rigidity measurements is feasible. The dual radiator configuration of NaF and aerogel makes isotope separation of light elements possible for energies in the range from 0.5 to 10 GeV/nucleon, approximately. Best mass resolutions are $\sim$~2\% at 3 GeV/nucleon for aerogel, and $\sim$~3\% at 1 GeV/nucleon for NaF. Techniques presented here may also be applied in the separation of antimatter isotopes which is of great importance in dark matter studies. \vspace{-0.2cm} | 7 | 10 | 0710.0846 |
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0710 | 0710.2558_arXiv.txt | We highlight the potential importance of gaseous TiO and VO opacity on the highly irradiated close-in giant planets. The atmospheres of these planets naturally fall into two classes that are somewhat analogous to the M- and L-type dwarfs. Those that are warm enough to have appreciable opacity due to TiO and VO gases we term the ``pM Class'' planets, and those that are cooler, such that Ti and V are predominantly in solid condensates, we term ``pL Class'' planets. The optical spectra of pL Class planets are dominated by neutral atomic Na and K absorption. We calculate model atmospheres for these planets, including pressure-temperature profiles, spectra, and characteristic radiative time constants. Planets that have temperature inversions (hot stratospheres) of $\sim$2000 K and appear ``anomalously'' bright in the mid infrared at secondary eclipse, as was recently found for planets \hh\ and \hd, we term the pM Class. Molecular bands of TiO, VO, H$_2$O, and CO will be seen in emission, rather than absorption. This class of planets absorbs incident flux and emits thermal flux from high in their atmospheres. Consequently, they will have large day/night temperature contrasts and negligible phase shifts between orbital phase and thermal emission light curves, because radiative timescales are much shorter than possible dynamical timescales. The pL Class planets absorb incident flux deeper in the atmosphere where atmospheric dynamics will more readily redistribute absorbed energy. This leads to cooler day sides, warmer night sides, and larger phase shifts in thermal emission light curves. We briefly examine the transit radii for both classes of planets. The boundary between these classes is particularly dependent on the incident flux from the parent star, and less so on the temperature of the planet's internal adiabat (which depends on mass and age), and surface gravity. Around a Sun-like primary, for solar composition, this boundary likely occurs at $\sim$0.04-0.05 AU, but uncertainties remain. We apply these results to pM Class transiting planets that are observable with the \emph{Spitzer Space Telescope}, including \hd, WASP-1b, TrES-3b, TrES-4b, \hh, and others. The eccentric transiting planets HD 147506b and HD 17156b alternate between the classes during their orbits. Thermal emission in the optical from pM Class planets is significant red-ward of 400 nm, making these planets attractive targets for optical detection via Kepler, COROT, and from the ground. The difference in the observed day/night contrast between $\upsilon$ Andromeda b (pM Class) and \he\ (pL Class) is naturally explained in this scenario. | 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. Given this incredible complexity, it is worthwhile to first look for physical processes that may be common to groups of planets. In addition to a mass and radius, one can further characterize a planet by studying its atmosphere. The visible atmosphere is a window into the composition of a planet and contains clues to its formation history \cp[e.g.,][]{Marley07b}. Of premier importance in this class of highly irradiated planets is how stellar insolation affects the atmosphere, as this irradiation directly affects the atmospheric structure, temperatures, and chemistry, the planet's cooling and contraction history, and even its stability against evaporation. Since irradiation is perhaps the most important factor in determining the atmospheric properties of these planets, we examine the insolation levels of the 23 known transiting planets. We restrict ourselves to those planets more massive than Saturn, and hence for now exclude treatment of the ``hot Neptune'' GJ 436b, which is by far the coolest known transiting planet. \mbox{Figure~\ref{flux}} illustrates the stellar flux incident upon the planets as a function of both planet mass (\mbox{Figure~\ref{flux}}\emph{a}) and planet surface gravity (\mbox{Figure~\ref{flux}}\emph{b}). In these plots diamonds indicate transiting planets and triangles indicate other interesting hot Jupiters, for which \emph{Spitzer Space Telescope} data exist, but which do not transit. The first known transiting planet, \hd, is seen to be fairly representative of these planets in terms of incident flux. Planets OGLE-TR-56b and OGLE-TR-132b are somewhat separate from the rest of the group because they receive the highest stellar irradiation. Both orbit their parent stars in less than 2 days and are prototypes of what has been called the class of ``very hot Jupiters'' \cp{Konacki03,Bouchy04} with orbital periods less than 3 days. However, orbital period is a poor discriminator between ``very hot'' and merely ``hot,'' as \he\ clearly shows. Labeled a ``very hot Jupiter'' upon its discovery, due to its short 2.2 day period \citep{Bouchy05}, \he\ actually receives a comparatively modest amount of irradiation due to its relatively cool parent star. Therefore, perhaps a classification based on incident flux, equilibrium temperature, or other attributes would be more appropriate. In this paper we argue that based on the examination of few physical processes that two classes of hot Jupiter atmospheres emerge with dramatically different spectra and day/night contrasts. Equilibrium chemistry, the depth to which incident flux will penetrate into a planet's atmosphere, and the radiative time constant as a function of pressure and temperature in the atmosphere all naturally define two classes these irradiated planets. Our work naturally builds on the previous work of \ct{Hubeny03} who first investigated the effects of TiO and VO opacity on close-in giant planet atmospheres as a function of stellar irradiation. These authors computed optical and near infrared spectra of models with and without TiO/VO opacity. In general they found that models with TiO/VO opacity feature temperature inversions and molecular bands are seen in emission, rather than absorption. Two key questions from the initial \ct{Hubeny03} investigation were addressed but could not be definitely answered were: 1) if a relatively cold planetary interior would lead to Ti/V condensing out deep in the atmosphere regardless of incident flux, thereby removing gaseous TiO and VO, and, 2) if this condensation did not occur, at what irradiation level would TiO/VO indeed be lost at the lower atmospheric temperatures found at smaller incident fluxes. Later \ct{Fortney06} investigated model atmospheres of planet \hh\ including TiO/VO opacity at various metallicities. Particular attention was paid to the temperature of the deep atmosphere \emph{P-T} profiles (as derived from an evolution model) in relation to the Ti/V condensation boundary. Similar to \ct{Hubeny03}, they found a temperature inversion due to absorption by TiO/VO and computed near and mid-infrared spectra that featured emission bands. Using the \emph{Spitzer} InfraRed Array Camera (IRAC) \ct{Harrington07} observed \hh\ in secondary eclipse with \emph{Spitzer} at 8 $\mu$m and derived a planet-to-star flux ratio consistent with a \ct{Fortney06} model with a temperature inversion due to TiO/VO opacity. At that point, looking at the work of \ct{Fortney06} and especially \ct{Hubeny03}, \ct{Harrington07} could have postulated that all objects more irradiated than \hh\ may possess inversions due to TiO/VO opacity, but given the single-band detection of \hh, caution was in order. More recently, based on the four-band detection of flux from \hd\ by \ct{Knutson08}, \ct{Burrows07c} find that a temperature inversion, potentialy due to TiO/VO opacity, is necessary to explain this planet's mid-infrared photometric data. Based on their new \hd\ model and the previous modeling investigations these authors posit that planets warmer than \hd\ may features inversions, while less irradiated objects such as \he\ do not, and discuss that photochemical products and gaseous TiO/VO are potential absorbers which may lead to this dichotomy. We find, as has been previously shown, that those planets that are warmer than required for condensation of titanium (Ti) and vanadium (V)-bearing compounds will possess a temperature inversion at low pressure due to absorption of incident flux by TiO and VO, and will appear ``anomalously'' bright in secondary eclipse at mid-infrared wavelengths. Thermal emission in the optical will be significant \cp{Hubeny03,Lopez07}. Furthermore, here we propose that these planets will have large day/night effective temperature contrasts. We will term these very hot Jupiters the ``pM Class,'' meaning gaseous TiO and VO are the prominent absorbers of optical flux. The predictions of equilibrium chemistry for these atmospheres are similar to dM stars, where absorption by TiO, VO, H$_2$O, and CO is prominent \cp{Lodders02b}. Planets with temperatures below the condensation curve of Ti and V bearing compounds will have a gradually smaller mixing ratio of TiO and VO, leaving Na and K as the major optical opacity sources \cp*{BMS}, along with H$_2$O, and CO. We will term these planets the pL class, similar to the dL class of ultracool dwarfs. These planets will have relatively smaller secondary eclipse depths in the mid infrared and significantly smaller day/night effective temperature contrasts. As discussed below, published \emph{Spitzer} data are consistent with this picture. The boundary between these classes, at irradiation levels (and atmospheric temperatures) where Ti and V may be partially condensed is not yet well defined. In this paper we begin by discussing the observations to date. We then give an overview of our modeling methods and the predicted chemistry of Ti and V. We calculate pressure-temperature (\emph{P-T}) profiles and spectra for models planets. For these model atmospheres we then analyze in detail the deposition of incident stellar flux and the emission of thermal flux, and go on to calculate characteristic radiative time constants for these atmospheres. We briefly examine transmission spectra before we apply our models to known highly irradiated giant planets. Before our discussion and conclusions we address issues of planetary classification. | Though 1D radiative-convective equilibrium model atmospheres we have addressed the class of atmospheres, the pM Class, for which TiO and VO are extremely strong visible absorbers \cp{Hubeny03}. This absorbed incident flux drives these planets to have hot ($\sim$2000+ K) stratospheres. Therefore, these planets will appear very bright in the mid-infrared, with brightness temperatures larger than their equilibrium temperatures. This is the case for \hh\ \cp{Harrington07}, and \hd\ \ct{Knutson08}. In addition, these planets will have large day/night temperature contrasts because radiative time constants at photospheric pressures are much shorter than reasonable advective timescales. The hottest point of the planet should be the substellar point, which absorbs the most flux, leading to perhaps negligible phase shift between the times of maximum measured thermal emission and when the day side is fully visible. This appears to be the situation for planet $\upsilon$ And b, observed by \ct{Harrington06}. Given that its irradiation level is intermediate between \hd\ and \hh\ we find that this planet is pM Class. Due to the fast radiative times, the day side of these planets may have \emph{P-T} profiles that do not deviate much from radiative equilibrium models. Although atmospheric dynamics will surely be vigorous, winds will be unable to advect gas before it cools to space. For these planets, high irradiation, the presence of gaseous TiO and VO, a hot stratosphere, the location of the hottest atmosphere at the substellar point, and a large day/night temperature contrast all go hand-in-hand. On the other hand, we have shown that in the pL Class, dominated by absorption by H$_2$O, Na, and K, photospheric pressures and temperatures prevail such that advective timescales and radiative timescales are similar \cp[see also][]{Seager05}. Since atmospheric dynamics will be important for the redistribution of energy, the consequences for the structure and thermal emission of these atmospheres will be quite complex. The efficiency of energy redistribution will vary with planetary irradiation level, surface gravity, and rotation rate. pL Class planets will have smaller day/night temperature contrasts and measurable phase shifts in thermal emission light curves that will be wavelength-dependent. In addition, these planets may show variability in secondary eclipse depth \cp[e.g.~][]{Rauscher07}. However, without a better understanding of the dynamics it is difficult to make detailed predictions at this time. Secondary eclipse depths should range somewhere between values expected for a ``full redistribution'' model and inefficient redistribution. The published secondary eclipse data for pL Class planets \T\ and \he\ are all consistent with this prediction \cp{Fortney05,Fortney07b}. In addition, the 8 $\mu$m light curves for 51 Peg b and \hd\ \cp{Cowan07} and \he\ \cp{Knutson07b} are consistent with this prediction as well. Examination of \mbox{Figure~\ref{flux}} shows that HD 179949b is a pM Class planet, and indeed \ct{Cowan07} found the largest phase variation it their small sample for this planet, but the unknown orbital inclination makes definitive conclusions difficult. Additional observational results will soon help to test the models presented here. We find that transiting planets WASP-1b, TrES-4b, TrES-3b, OGLE-Tr-10b, and TrES-2b will be in the pM Class, along with non-transiters $\upsilon$ And b and HD 179949. The low-irradiation boundary of this class is not yet clear, and planet \hd\ shows that temperature inversions persist to irradiation levels where TiO/VO are expected to begin being lost to condensation \cp{Burrows07c}. At still lower irradiation levels, the limited data for \he\ lead us to conclude it is pL Class \cp{Fortney07b}. Secondary eclipse data for XO-2b, HAT-P-1b, and WASP-2b will be important is determining how temperature inverstions (and the TiO/VO abundances) wane with irradiation level. Just as in dM and dL stars, the condensation of Ti and V is expected to be gradual process, so we fully expect transition objects between the distinct pM and pL class members. We will soon have additional information that will help shed light on the atmosphere of \hd\. The 8 $\mu$m light curve for \hd\ obtained by \ct{Cowan07} shows little phase variation. However, if our theory connecting TiO/VO opacity and temperature inversions to large day/night contrasts is correct, we expect to see a large variation. Soon H.~Knutson and collaborators will obtain half-orbit light curves for \hd\ at 8 and 24 $\mu$m. The quality should be comparable to that obtained by \citet{Knutson07b} for \he, and will put our theory to the test. HD 147506, with an eccentricity of 0.517 \cp{Bakos07b}, and HD 17156, with an eccentricity of 0.67 \cp{Fischer07,Barbieri07}, will be extremely interesting cases as the flux they receive varies by factors of 9 and 26, respectively, between periapse and apoapse. They should each spend part of their orbits as pL class and part as pM class. This makes predictions difficult, but large day/night temperatures differences at apoapse are likely. There has recently been considerable discussion on the relative merits of multi-dimensional dynamical models and 1D radiative-convective model atmospheres for these highly irradiated planets. Both kinds of studies provide interesting predictions. While it could be claimed that 1D radiative-convective models are unrealistic because they ``lack dynamics,'' they do include very detailed chemistry, vast opacity databases, and advanced non-gray radiative transfer, which all dynamics models for these planets lack. Given the large diversity in predictions among the various dynamical models, which use a host of simplifications, the next step will be combining dynamics and radiative transfer, an idea which has been mentioned or advanced by a number of authors \cp{Seager05,Barman05,Fortney06b,Burrows06,Dobbs07}. Eventually we will be able to give up 1D $f$-type parameters to treat the incident flux in a more realistic fashion for these exotic atmospheres. We recently started working toward this goal in \ct[][see also \citealp{Dobbs07}]{Fortney06b}; work continues, and we believe it will have a promising future. We think that the predictions we have made here with a 1D model will provide a framework for understanding the observations to come. Additional observations for both pL and pM class planets, along with additional theoretical and modeling efforts, should further clarify our understanding. | 7 | 10 | 0710.2558 |
0710 | 0710.3148_arXiv.txt | In this paper, warm inflationary models on a brane scenario are studied. Here we consider slow-roll inflation and high-dissipation regime in a high-energy scenario. General conditions required for these models to be realizable are derived. We describe scalar and tensor perturbations for these scenarios. Specifically we study power-law potentials considering a dissipation parameter to be a constant on the one hand and $\phi$ dependent on the other hand. We use recent astronomical observations to restrict the parameters appearing in our model. | It is well known that many long-standing problems of the Big Bang model, namely the horizon problem, flatness, homogeneity and the numerical density of monopoles, may find a natural explanation in the frame of the inflationary universe model \cite{Guth1981,Inflation,Inflationary}. Perhaps the most relevant feature of the inflationary universe model is that it provides a causal interpretation for the origin of the observed anisotropy in the cosmic microwave background (CMB) radiation, and also the distribution of large-scale structures \cite{WMAP1,WMAP3}. But the inflationary universe model has problems too. One of the problems in it is how to attach the observed universe to the end of the inflationary epoch; there are three possible solutions to this problem: reheating \cite{KolbandTurner}, preheating \cite{TBKLS} and warm inflation \cite{Berera1995,deOliveira1998,Berera1999,Bellini1998}. In this work we focus on the latter. In standard inflationary universe models, the acceleration of the universe is driven by a scalar field (named inflaton) with a specific scalar potential. These kinds of model are divided into two regimes, the slow-roll and reheating ephocs. In the slow-roll period the universe inflates and all interactions between the inflaton scalar field and any other field are typically neglected. Subsequently, a reheating period is invoked to end the period of inflation. After reheating, the universe is filled with radiation \cite{afterReheating,Inflation}, and then the universe gets connected with the Big Bang model. Warm inflation is an alternative mechanism to have successful inflation and avoid the reheating period \cite{Berera1995}. In this kind of model, dissipative effects are important during inflation, so that radiation production occurs concurrently with the inflationary expansion. The inflaton interacts with a thermal bath via a friction term, where phenomenologically the decay of the scalar field is described by means of an interaction Lagrangian. For instance, the authors of Ref.\cite{new} take the interaction terms of the form $\frac{1}{2}\lambda^2\phi^2\chi^2$ and $g\chi\bar{\psi}\psi$, where the inflationary period presents a two-stage decay chain $\phi\rightarrow\chi\rightarrow\psi$. In this case, they reported that the damping term $\Gamma$ becomes $\frac{\lambda^3g^2\phi}{256\pi^2}$. Note that if the scalar field changes a little bit, then the friction coefficient $\Gamma$ remains almost constant. From the point of view of statistical mechanics, the interaction between quantum fields and a thermal bath could be illustrated by a general fluctuation-dissipation relation \cite{new2}. Warm inflation was criticized on the basis that the inflaton cannot decay during the slow-roll phase \cite{new3}. However, in recent years, it has been shown that the inflaton can indeed decay during the slow-roll phase (see \cite{new4} and references therein) whereby it now rests on solid theoretical grounds. On the other hand, the inclusion of the damping term into the model needs a very special scheme: for instance, in Ref.\cite{Profe} it was found that when $\Gamma$ is constant, in the slow-roll approximation a wrong value for the number of e-fold was obtained: meanwhile, in the power-law approach this problem is absent. Warm inflation ends when the universe heats up to become radiation dominated. At this epoch the universe stops inflating and `smoothly' enters into a radiation dominated Big Bang phase \cite{Berera1995}. The matter components of the universe are created by the decay of either the remaining inflationary field or the dominant radiation field \cite{TaylorandBerera2000}. In standard inflationary universe models the quantum fluctuations associated to the inflaton scalar field generate the density perturbations seeding the structure formation at a late time in the evolution of the universe. Instead, in warm inflation models, the density fluctuations arise from thermal rather than quantum fluctuations \cite{Berera2004,Berera1995}. These fluctuations have their origin in the hot radiation and influence the inflaton through a friction term in the equation of motion of the inflaton scalar field \cite{Berera1996}. Several aspects of warm inflationary universe models have been studied in the past few years. The goal of the present work is to investigate warm inflationary models on brane scenarios, where the total energy density $\rho=\rho_{\phi}+\rho_{\gamma}$ is found on the brane \cite{Profe2}. The universe is filled with a self-interacting scalar field of energy density $\rho_{\phi}$ and a radiation field with energy density $\rho_{\gamma}$. The motivation for introducing brane scenarios is the increasing interest in higher dimensional cosmological models, motivated by superstring theory, where the matter fields (related to open string modes) are confined to a lower dimensional brane, while gravity (closed string modes) can propagate in the bulk \cite{Strings}. Shiromizu et al. \cite{Shiromizu} have found the four-dimensional Einstein's equations projected onto the brane. These projections introduce some differences in the fundamental field equations, such as the Friedmann equation and, therefore, in the equations that describes the linear perturbations theory \cite{Brane}. Our aim is quantify the modifications of the warm inflation model in the brane scenario for arbitrary inflaton potentials on the brane. In order to do this we study the linear theory of cosmological perturbations for a warm inflationary model on a brane. The perturbations are expressed in term of different parameters appearing in our model: these parameters will be constrained from the WMAP three-year data \cite{WMAP3}. In section II we describe the dynamics of our model and we establish some approximations that we use in this work. In section III we investigate the linear theory of perturbations. Here we calculate the scalar perturbations in the longitudinal gauge and also the tensor perturbations. In section IV we take a power-law potential and we investigate the high-energy and high-dissipation regime considering a power-law dissipation coefficient $\Gamma$. Finally, in section V, we give some conclusions. We have used units in which $c=\hbar=1$. | In this paper we have considered a warm inflationary scenario on a brane. We have restricted ourselves to a high-dissipation and high-energy regime. In the slow-roll approximation we have found a general relationship between radiation and scalar field densities; see Eq.(\ref{eq21}). In relation to the perturbations we have considered that the there does no exist any interaction between the brane and the bulk, i.e., we have neglected back-reaction due to metric perturbations in the fifth dimension. We note that a full investigation is required to discover when back-reaction will have a significant effect in the perturbations. With the above-mentioned restriction we have obtained explicitly the contributions of the adiabatic and entropy modes. We have shown that the dissipation parameter plays a crucial role in producing the entropy mode (see Eq.(\ref{eq62b})). A general relation for the density perturbations is given in Eq.(\ref{eq41}). The tensor pertubations are generated via stimulated emission into the existing thermal background (see Eq.(\ref{eq52})) and the tensor-scalar ratio is modified by a temperature-dependent factor. We have studied a power-law potential for different dependence of the dissipation coefficient $\Gamma$. From the normalization of the WMAP three-year data, the potential becomes of the order of $V(\phi_0)\sim10^{-20}M_4^4$ when it leaves the horizon at the scale of $k_0=0.002\text{Mpc}^{-1}$. As in the situation studied in Ref.\cite{Chaotic}, the value of the potential depends on $M_5$. Here we have considered $M_5\sim10^{-5}M_4$. To fulfill the approximations in our model we have restricted the range of the parameters in which warm inflation on a brane can occur (see FIGS. \ref{fig:HT}, \ref{fig:VM5}). In order to show some explicit results we have chosen the following set of values for the parameters, $T_r\simeq T=3.5\times10^{-6}M_4$, $r=1000$ and $s=100$. First, we have considered a chaotic potential and a constant dissipation coefficient; for this case we have found that the spectrum is driven toward an scale-invariant spectrum and the running of the spectral index is extremely small. The situation was similar when we considered a variable dissipation coefficient. In the case of a power-law potential with $n=4$ and a constant dissipation coefficient we have found that the value to $n_s$ is smaller than the previous case but out of the range given by WMAP three-year data. On the other hand, the running of the spectral index is closer to the value given by WMAP three-year data but one order of magnitud smaller. Finally the best situation is found when we consider a constant dissipation coefficient and a $n=6$ potential. In this case both parameters, $n_s$ and $\alpha_s$, are in the ranges specified by WMAP. It is necessary to note that with our approximations we have found that for the case in which $n=6$ and $m=0$ is favored in the light of the recent results reported by WMAP three-year data. On the other hand, we have considered some approximations; as a result, we were not able to find any real solution in the allowed range of parameters. For example, in the slow-roll approximation, in the case $n=4$ and, when we try to consider a variable coefficient of dissipation $\Gamma$, we did not find any real solution. However, we think that we could find a solution if in place of using this approximation we assume a power-law for the scale factor. We intend to return to this point (and others) in the near future. | 7 | 10 | 0710.3148 |
0710 | 0710.1651_arXiv.txt | We report on an investigation of the environments of the SLACS sample of gravitational lenses. The local and global environments of the lenses are characterized using SDSS photometry and, when available, spectroscopy. We find that the lens systems that are best modelled with steeper than isothermal density profiles are more likely to have close companions than lenses with shallower than isothermal profiles. This suggests that the profile steepening may be caused by interactions with a companion galaxy as indicated by N-body simulations of group galaxies. The global environments of the SLACS lenses are typical of non-lensing SDSS galaxies with comparable properties to the lenses, and the richnesses of the lens groups are not as strongly correlated with the lens density profiles as the local environments. Furthermore, we investigate the possibility of line-of-sight contamination affecting the lens models but do not find a significant over-density of sources compared to lines of sight without lenses. | The Sloan Lens ACS Survey \citep[SLACS;][]{bolton} is a large sample of strong gravitational lenses derived from the Sloan Digital Sky Survey (SDSS). The lens sample has proved to be particularly useful due to the high quality of the data; all of the SLACS lenses have known lens and source redshifts, stellar velocity dispersions for the lensing galaxies have been measured from the SDSS spectra, all of the systems have {\em Hubble Space Telescope} ({\em HST}) ACS imaging in two bands for accurate non-parametric lens modelling, and all of the sources are extended and therefore provide additional constraints on the mass profile of the lensing galaxy. Furthermore, the density of background sources in the {\em HST} imaging allows a weak lensing analysis of the ensemble sample of lenses \citep{gavazzi}; the major drawback of the lens sample is that it is unlikely to find any variable sources that would provide time delays. \citet{koopmansSLACS} have used the SLACS lenses to show that early-type galaxies have isothermal total inner density profiles with very little intrinsic scatter (approximately 6 per cent) assuming a uniformity in the environments of the lenses that has not been rigorously tested. While the SLACS lenses seem to lie on the Fundamental Plane \citep{bolton,treu,bolton07} and do not differ noticeably from other SDSS galaxies with similar luminosities and stellar velocity dispersions, the local environments of the lenses might affect the mass profiles of the lensing galaxies \citep[e.g.,][]{rusin,dobke,augerb}. If some of the lenses are being perturbed by neighbouring galaxies the intrinsic scatter of the density slope for {\em isolated} early-type galaxies might be even smaller than 6 per cent. Furthermore, it has been suggested that line-of-sight (LOS) contamination significantly affects the SLACS lenses \citep{guimaraes} and the density slope might be expected to be shallower than originally reported. We report on a spectroscopic and photometric evaluation of the environments and lines of sight of the 15 SLACS lenses investigated by \citet{koopmansSLACS}. A weighting scheme is used to determine the effective number of potential perturbing companions to each lens galaxy. We also characterize the `richness' of the global environment of each lens field and quantify the number of galaxies along the LOS to the lens. Throughout this paper the term `global environment' is used to describe the group, cluster, or field in which the lens resides while the `local environment' describes the environment within $\approx 100$ \hinv kpc of the lensing galaxy. A $\Lambda$CDM cosmology with $\Omega_M = 0.27$ and $\Omega_\Lambda = 0.73$ is used to determine all physical distances, which are measured in \hinv units. | We find that the global environments of the SLACS lenses are typical of other massive early-type galaxies found in the SDSS. Two of the steeper than isothermal systems lie in very over-dense regions but the remaining lenses all have richness values that fall near the peak of the richness distribution (Figure \ref{figure_comp_rich}). Furthermore, the suggestion that the SLACS lenses are affected by LOS contamination does not seem to be merited by the data. While there are other galaxies along the lines of sight to the lens systems, the LOS densities do not significantly deviate from the densities along comparable lines of sight. We therefore expect that the parameter estimates should not be affected as proposed by \citet{guimaraes}. The SDSS photometric data indicate that the SLACS lenses are only slightly more likely to be associated with companion galaxies than comparable lenses selected from the SDSS, though the uncertainty of the photometric identifications makes the difference negligible. We therefore conclude that SLACS lenses do lie in typical environments both globally and locally. However, we also find that lens systems with steeper than isothermal density slopes are preferentially associated with companion galaxies compared to lenses with shallower density slopes. N-body simulations suggest that interactions with neighbouring galaxies can induce a steepening in the density slope \citep{dobke} and there are other lens systems with companion galaxies that are found to be best modelled with steeper than isothermal profiles \citep[e.g.,][]{rusin,augerb}. The interaction-induced steepening is a transient effect and the density profile of the galaxy will return to isothermal approximately 0.5-2 Gyr after the encounter with the neighbour \citep{dobke}. This may account for the large range of $N_w$ for lenses with nearly isothermal profiles; isothermal lenses with a companion may be in the relaxed state before or after an encounter with the neighbour galaxy. This stripping mechanism may also account for local observations of dark matter deficient galaxies \citep[e.g.,][]{romanowsky,proctor}. We note that one steeper-profile system, SDSSJ1250+0523, is not photometrically associated with any neighbouring galaxies; this perhaps illustrates the limits of using photometry to find perturbing companions or demonstrates that other factors also influence the slope of the density profile. Additionally, we have not found an environmental bias to account for the shallower lenses. However, if a lens galaxy is embedded in a cluster, the joint profile of the cluster and galaxy would tend to be modelled with a shallower than isothermal profile if a single-component power law is used (ignoring interactions between the two halos). This effect is dependent on the location of the lens with respect to the centre of the cluster, which our photometric analysis is unable to address. More complete field spectroscopy would better characterize the global environments of these lens systems and allow correlations between the mass slopes and cluster centre offsets to be investigated. Furthermore, a spectroscopic investigation of the local environments of the complete sample of SLACS early-type lenses would confirm the correlation indicated by our photometric analysis and provide strong evidence for truncation caused by galaxy interactions. | 7 | 10 | 0710.1651 |
0710 | 0710.0711_arXiv.txt | { We present results of $H\alpha$ imaging for 42 galaxies in the nearby low-density cloud Canes Venatici I populated mainly by late-type objects. Estimates of the $H\alpha$ flux and integrated star formation rate ($SFR$) are now available for all 78 known members of this scattered system, spanning a large range in luminosity, surface brightness, $HI$ content and $SFR$. Distributions of the CVnI galaxies versus their $SFR$, blue absolute magnitude and total hydrogen mass are given in comparison with those for a population of the nearby virialized group around M81. We found no essential correlation between star formation activity in a galaxy and its density environment. A bulk of CVnI galaxies had enough time to generate their baryon mass with the observed $SFR$. Most of them possess also a supply of gas sufficient to maintain their observed $SFR$s during the next Hubble time. } | The distribution over the sky of 500 galaxies of the Local volume with distances within 10 Mpc shows considerable inhomogeneities due to the presence of groups and voids. Apart from several virialized groups, like the group around the galaxy M81, an amorphous association of nearby galaxies in the Canes Venatici constellation was noted by many authors (Karachentsev 1966, de Vaucouleurs 1975, Vennik 1984, Tully 1988). The boundaries of it are rather uncertain. Roughly, in a circle of radius $\sim20\degr$ around the galaxy NGC~4736 there are about 80 known galaxies with $D<10$ Mpc, which corresponds to a density contrast on the sky $\Delta N/N\sim4$. Inside of this complex the distribution of galaxies is also inhomogeneous, showing some clumps differing in their location on the sky and distances in depth. About 70\% of the population of the cloud accounts for irregular dwarf galaxies, whose masses are obviously insufficient to keep such a system in the state of virial equilibrium. As data on distances of galaxies accumulated, a possibility appeared of studying the kinematics of the cloud in details. As was shown by Karachentsev et al. (2003), the CVnI cloud is in a state close to the free Hubble expansion, having a characteristic crossing time of about 15 Gyr. Being a scattered system with rare interactions between galaxies, the nearby CVnI cloud is a unique laboratory for studying star formation processes in galaxies running independently, without a noticeable external influence. Kennicutt et al. (1989), Hoopes et al. (1999), van Zee (2000), Gil de Paz et al. (2003), James et al. (2004) and Hunter \& Elmegreen (2004) conducted observations in the $H\alpha$ line of three dozen galaxies of this complex, which made it possible to determine the star formation rate ($SFR$) in them. However, more than half of other members of the cloud proved to be out of vision of these authors. Our task consisted in the completion of $H\alpha$-survey of the population of the CVnI cloud. The results of our observations and their primary analysis are presented in this paper. | A systematic survey of $H{\alpha}$-emission in the nearest scattered cloud CVnI shows that in most of its galaxies the process of active star formation is on despite the low density contrast of this cloud and rather rare interaction between its members. By making full use of our $H{\alpha}$-survey, we can estimate the mean density of $SFR$ in the cloud. A conical volume in the distance range from 2 to 10 Mpc, resting in a sky region of $\sim1500$ square degrees, makes 153 Mpc$^3$. In this volume we have a summary value $\Sigma(SFR)=18.6M_{\sun}$yr$^{-1}$, which yields an average density of star formation rate $\dot{\rho}_{SFR}=0.12M_{\sun}$yr$^{-1}$Mpc$^{-3}$. The obtained value turns out to be a little less than $\dot{\rho}_{SFR}=0.165M_{\sun}$yr$^{-1}$Mpc$^{-3}$ for a ``sell of homogeneity'' embracing the group M81 (Karachentsev \& Kaisin, 2007). According to Nakamura et al. (2004), Martin et al. (2005) and Hanish et al. (2006), the average star formation rate per 1 Mpc$^3$ at the present epoch (z=0) is (0.02--0.03)$M_{\sun}$yr$^{-1}$Mpc$^{-3}$. Consequently, the CVnI cloud has excess of $SFR$ density 4--6 times higher in comparison with the mean global quantity, which roughly corresponds to the cloud density contrast on the sky, $\Delta N/N\sim4$. Herefrom we conclude that the CVnI cloud is characterised by a usual norm of star formation in its galaxies. | 7 | 10 | 0710.0711 |
0710 | 0710.3120.txt | We present a new measurement of the volumetric rate of Type Ia supernova up to a redshift of 1.7, using the Hubble Space Telescope (HST) GOODS data combined with an additional HST dataset covering the North GOODS field collected in 2004. We employ a novel technique that does not require spectroscopic data for identifying Type Ia supernovae (although spectroscopic measurements of redshifts are used for over half the sample); instead we employ a Bayesian approach using only photometric data to calculate the probability that an object is a Type Ia supernova. This Bayesian technique can easily be modified to incorporate improved priors on supernova properties, and it is well-suited for future high-statistics supernovae searches in which spectroscopic follow up of all candidates will be impractical. Here, the method is validated on both ground- and space-based supernova data having some spectroscopic follow up. We combine our volumetric rate measurements with low redshift supernova data, and fit to a number of possible models for the evolution of the Type Ia supernova rate as a function of redshift. The data do not distinguish between a flat rate at redshift $>$ 0.5 and a previously proposed model, in which the Type Ia rate peaks at redshift $\sim$ 1 due to a significant delay from star-formation to the supernova explosion. Except for the highest redshifts, where the signal to noise ratio is generally too low to apply this technique, this approach yields smaller or comparable uncertainties than previous work. | %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%5 The empirical evidence for the existence of dark energy came from observations of Type Ia supernovae~\citep{bib:riess, bib:P99, bib:P03}, which are believed to arise from the thermonuclear explosion of a progenitor white dwarf after it approaches the Chandrasekhar mass limit~\citep{bib:chan}. However, the physics of Type Ia supernova production is not well understood. The two most plausible scenarios for the white dwarf to accrete the necessary mass are the single degenerate case, where the white dwarf is located in a binary system; and the double degenerate case, where two white dwarfs merge. The Type Ia supernova rate is correlated with the star formation history (SFH), and thus a measurement of the rate as a function of redshift helps constrain the possible type Ia progenitor models. In addition to its importance for understanding Type Ia supernovae as astronomical objects, a good grasp of the Type Ia supernova rate to high redshifts is important for the next generation of proposed space-based supernova cosmology experiments, such as SNAP~\citep{bib:snap}. It is therefore of great practical interest to determine the rate of Type Ia supernovae at redshifts $>$ 1. The subject of Type Ia supernova rates has been addressed by many authors in the past. Existing rate measurements have been mostly limited to redshift ranges $<$ 1: the results of~\cite{bib:cappellaro},~\cite{bib:hardin},~\cite{bib:madgwick}, and~\cite{bib:blanc} measure the rates at redshifts $\leq$ $\sim$0.1;~\cite{bib:neill},~\cite{bib:tonry}, and~\cite{bib:pain}, at intermediate redshifts of 0.47, 0.50, and 0.55, respectively; and~\cite{bib:barris}, up to a redshift of 0.75. The only published measurement of the rates at redshifts $>$ 1 is that of~\cite{bib:dahlen}, who analyzed the GOODS dataset. There are several important differences that distinguish our work from that of~\cite{bib:dahlen}. First, we augment the GOODS sample with the HST data collected during the Spring-Summer 2004 high redshift supernova searches. Second, our methods of calculating the control time (the time during which a supernova search is potentially capable of finding supernova candidates) and the efficiency to identify a supernova are based on a detailed Monte Carlo simulation technique using a library of supernova templates. Third, we adopt a novel approach to typing supernovae, using photometric data and a Bayesian probability method described in~\cite{bib:ourpaper}. The Bayesian technique is able to perform classification using only photometric data, and therefore does not require spectroscopic follow up. Optionally, photometric or spectroscopic redshifts can be used to improve the classification accuracy. Our initial requirements on potential supernova candidates are more stringent in terms of the number of points on the light curve and the signal to noise of those points than those of~\cite{bib:dahlen}; thus some of the candidates they identified will fail our cuts. However, we are able to reliably separate Type Ia supernova from other supernovae types based on their Bayesian probability, with an efficiency that is readily quantifiable, thus allowing us to use larger data samples. Our approach therefore avoids the problems that arise in estimating the efficiency for the decision to schedule spectroscopic follow up based on a potentially low signal-to-noise initial detection. The Bayesian classification technique uses photometric data, and does not require any spectroscopic followup. This is an advantage for future large-area surveys (such as the Dark Energy Survey, Pan-STARRS, and LSST) that will discover thousands of supernova candidates, but are unlikely to be able to obtain spectroscopic data for all of them, to distinguish Type Ia supernovae from core collapse supernovae and other variable objects. The technique described here can be considered a prototype of the kind of analysis that could be performed on these future large data sets to identify Type Ia supernovae for cosmological studies. There is a clear trade-off involved in using photometric measurements alone: if the quality of the photometric data is poor, then the efficiency of this technique to identify Type Ia supernovae is reduced; on the other hand, this technique enables larger samples of Type Ia from imaging surveys to be identified for cosmological studies, without the need for time-consuming spectroscopic follow up. Note that although the method is able to perform the supernova typing with photometric data alone (\emph{i.e.}, it does not require spectroscopic data, either redshifts or types), it is certainly able to use the extra information that is available, and in fact 70\% of the supernova candidates discussed in the present work have redshifts which were obtained spectroscopically. It is also worth noting that while in this paper we only analyze the Type Ia supernova rates, the Bayesian classification technique can be used to classify other types as well, making it possible to measure the rates of non-Type Ia supernovae in a similar fashion. These analyses will be presented in future publications. The paper is organized as follows. In section~\ref{sec:data} we describe the data samples used in the analysis. In section~\ref{sec:candselection} we describe the supernova candidate selection and typing process. In section 4 we calculate the control time, survey area, and search efficiency, and determine the volumetric Type Ia supernova rate from our data sample. A comparison of the rates with those reported in the literature is given in section~\ref{sec:comp}, and fits of the rates to different models relating the Type Ia supernova rates to the SFH are given in section~\ref{sec:sf}. A summary is given in section~\ref{sec:concl}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | \label{sec:concl} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% We have analyzed the rates of Type Ia supernovae up to a redshift of 1.7 using two samples collected with the HST: the GOODS data, and the 2004 ACS sample collected in the Spring-Summer 2004 covering the GOODS North field. Using only the data from two broadband filters, F775W and F850LP, we applied a novel technique for identifying Type Ia supernovae based on a Bayesian probability approach. This method allows us to automatically type supernova candidates in large samples, properly taking into account all known sources of systematic error. We also make use of the best currently available full spectral templates for five different supernova types for the candidate typing, as well as for calculating the efficiency of our supernova search, and the control time. These templates will undoubtedly be improved over the next several years as more supernova data becomes available. Current and upcoming supernova surveys will not only provide a better understanding of individual supernova types, but may also uncover new types of supernovae, which can then be added to the Bayesian classification framework. Likewise, a better understanding of the many parameters that affect supernova observations will improve the classification scheme, which will result in better constraints on the measured rates. The calculations of the supernova finding efficiency, the control time, and the survey area are all done taking into account the specific observing configurations pertinent for the surveys, such as exposure times, cadences, and the orientations of the GOODS tiles. We carried out a comparison of the predicted and observed numbers of supernovae in redshift bins of $\Delta \bar{z}$ = 0.1, for two different models of the Type Ia supernova rates: a redshift-independent rate and a power-law redshift-dependent rate. We find that the available data fit both models equally well. For comparison with previous work, particularly that of~\cite{bib:dahlen}, who also analyzed a large subset of the data used here, we calculated the volumetric Type Ia supernova rates in four redshift bins, 0.2 $\leq$ $\bar z$ $<$ 0.6, 0.6 $\leq$ $\bar z$ $<$ 1.0, 1.0 $\leq$ $\bar z$ $<$ 1.4, and 1.4 $\leq$ $\bar z$ $<$ 1.7. We find that our results are generally consistent with those of~\cite{bib:dahlen}. Due to the larger of number supernova candidates which this Bayesian classification technique makes available, we obtain smaller or equal uncertainties in all the bins up to $z$ = 1.7. In the highest redshift bin we obtain a larger uncertainty because the signal to noise ratio is generally too low to apply this technique. We fitted the resulting rates to two leading models used in recent literature: the two-component model and a Gaussian time delay model. The former model implies an increase in the Type Ia supernova rates at highest redshifts; while the latter, a decrease. We find that the statistics of the present sample does not definitively discriminate between the two scenarios -- only one supernova in this work and two supernovae in~\cite{bib:dahlen} contribute to the important highest-redshift bin. Significantly larger surveillance time would be required to arrive at a conclusive statement on the trends for the Type Ia rates at high redshifts. In the future, several ambitious new surveys are planned that will collect photometric data for thousands of supernovae in order to improve the constraints on dark energy. Individual spectroscopic follow up for every supernova candidate is likely to be impractical in these surveys. The Bayesian classification method described here has the ability to classify supernovae using photometric measurements alone, and is a promising technique for these future surveys. %%%%%%%% | 7 | 10 | 0710.3120 |
0710 | 0710.4129.txt | {} {To compute the chemical evolution of spiral bulges hosting Seyfert nuclei, based on updated chemical and spectro-photometrical evolution models for the bulge of our Galaxy, to make predictions about other quantities measured in Seyferts, and to model the photometric features of local bulges. The chemical evolution model contains updated and detailed calculations of the Galactic potential and of the feedback from the central supermassive black hole, and the spectro-photometric model covers a wide range of stellar ages and metallicities.} {We computed the evolution of bulges in the mass range $2\times 10^{9}-10^{11}M_{\odot}$ by scaling the efficiency of star formation and the bulge scalelength as in the inverse-wind scenario for elliptical galaxies, and considering an Eddington limited accretion onto the central supermassive black hole.} {We successfully reproduced the observed relation between the mass of the black hole and that of the host bulge. The observed nuclear bolometric luminosity emitted by the supermassive black hole is reproduced only at high redshift or for the most massive bulges; in the other cases, at $z \simeq 0$ a rejuvenation mechanism is necessary. The energy provided by the black hole is in most cases not significant in the triggering of the galactic wind. The observed high star formation rates and metal overabundances are easily achieved, as well as the constancy of chemical abundances with the redshift and the bulge present-day colours. Those results are not affected if we vary the index of the stellar IMF from $x=0.95$ to $x=1.35$; a steeper IMF is instead required in order to reproduce the colour-magnitude relation and the present $K$-band luminosity of the bulge.} {We show that the chemical evolution of the host bulge, with a short formation timescale of $\sim 0.1$ Gyr, a rather high efficiency of star formation ranging from $11$ to $50$ Gyr$^{-1}$ according to the bulge mass and an IMF flatter with respect to the solar neighbourhood, combined with the accretion onto the black hole is sufficient to explain the main observed features of Seyfert galaxies.} | The outstanding question of the co-evolution of Active Galactic Nuclei (AGNs) and their host galaxies has received considerable attention in the past decades, since various pieces of evidence pointed to a link between the formation of supermassive black holes (BHs) and the formation and evolution of their host spheroids: for example, the usual presence of massive dark objects at the centre of nearby spheroids (Ford et al. 1997; Ho 1999; Wandel 1999); the correlation between the BH mass and the stellar velocity dispersion of the host (for quiescent galaxies, Ferrarese \& Merritt 2000; Gebhardt et al. 2000a; Tremaine et al. 2002; for active galaxies, Gebhardt et al. 2000b; Ferrarese et al. 2001; Shields et al., 2003; Onken et al. 2004; Nelson et al. 2004) or its mass (Kormendy \& Richstone 1995; Magorrian et al. 1998; Marconi \& Hunt 2003; Dunlop et al. 2003); the similarity between light evolution of quasar (QSO) population and the star formation history of galaxies (Cavaliere \& Vittorini 1998; Haiman et al. 2004); the establishment of a good match among the optical QSO luminosity function, the luminosity function of star-forming galaxies and the mass function of dark matter halos (DMHs) at $z\sim 3$ (Haenhelt et al., 1998). The most widely accepted explanation for the luminosity emitted by an AGN, is radiatively efficient gas accretion onto a central supermassive BH. %Some theoretical model link quasar activity with major merger events %in a hierarchical galaxy formation process (Kauffmann \& Haenhelt %2000; Kauffmann et al., 2003). The outflows from AGNs can profoundly affect the evolution of the host galaxy, e.g. by quenching or inducing the star formation (e.g., see Ciotti \& Ostriker 2007, and references therein). The mutual feedback between galaxies and QSOs was used as a key to solve the shortcomings of the semianalytic models in galaxy evolution, e.g. the failure to account for the surface density of high-redshift massive galaxies (Blain et al., 2002; Cimatti et al., 2002) and for the $\alpha$-enhancement as a function of mass (Thomas et al., 2002), since it could provide a way to invert the hierarchical scenario for the assembly of galaxies and star formation (see e.g. Monaco et al., 2000; Granato et al., 2004; Scannapieco et al., 2005). The study of the chemical abundances of the QSOs was first undertaken by Hamann \& Ferland (1993), who combined chemical evolution and spectral synthesis models to interpret the N~{\scshape v}/C~{\scshape iv} and N~{\scshape v}/He~{\scshape ii} broad emission line ratios, and found out that the high metallicities and the abundance ratios of the broad-line region are consistent with the outcomes of the models for giant elliptical galaxies (Arimoto \& Yoshii, 1987; Matteucci \& Tornamb\`e, 1987; Angeletti \& Giannone, 1990), where the timescales of star formation and enrichment are very short and the initial mass function (IMF) is top-heavy. In the same year, Padovani \& Matteucci (1993) and Matteucci \& Padovani (1993) employed the chemical evolution model of Matteucci (1992) to model the evolution of radio-loud QSOs, which are hosted by massive ellipticals, following in detail the evolution of several chemical species in the gas. They supposed that the mass loss from dying stars after the galactic wind provides the fuel for the central BH and modeled the bolometric luminosity as $L_{bol}=\eta\dot{M}c^2$, with a typical value for the efficiency of $\eta = 0.1$ and were successful in obtaining the estimated QSO luminosities and the observed ratio of AGN to host galaxy luminosity. Then, they studied the evolution of the chemical composition of the gas lost by stars in elliptical galaxies and spiral bulges for various elements (C, N, O, Ne, Mg, Si and Fe), and found out that due to the high star-formation rate (SFR) of spheroids at early times the standard QSO emission lines were naturally explained. The relatively weak observed time dependence of the QSO abundances for $t \gtrsim 1$ Gyr was also predicted. The model of Matteucci \& Padovani (1993) still followed the classic wind scenario, where the efficiency of star formation decreases with increasing galactic mass and which was found to be inconsistent with the correlation between spheroid mass and $\alpha$-enhancement (Matteucci 1994). Moreover, Padovani \& Matteucci (1993) pointed out that if all mass lost by stars in the host galaxy after the wind were accreted by the central BH, the final BH mass would be up to two orders of magnitude larger than observed. Other works (Fria\c ca \& Terlevich 1998; Romano et al. 2002; Granato et al., 2004), which had a more refined treatment of gas dynamics, limited their analysis of chemical abundances to the metallicity $Z$ and the [Mg/Fe] ratio and their correlation with the galactic mass. All these studies were mainly devoted to studying the co-evolution of radio-loud QSOs and their host spheroids, which are elliptical galaxies. Now we want to extend the approach of Padovani \& Matteucci (1993) to AGNs hosted by spiral bulges, with a more recent chemical evolution model for the bulge with the introduction of the treatment of feedback from the central BH and a more sophisticated dealing of the accretion rate. Since Seyfert nuclei are preferentially hosted by disk-dominated galaxies (Adams, 1977; Yee, 1983; MacKenty, 1990; Ho et al. 1997) our study can be applied to this class of objects. The paper is organized as follows: in \S 2 we illustrate the chemical and photometrical evolution model, in \S 3 we show our calculations of the potential energy and of the feedback from supernovae (SNe) and from the AGN, in \S 4 we discuss our results concerning the black hole masses and luminosities, the chemical abundances and the photometry, and in \S 5 we draw some conclusions. | 7 | 10 | 0710.4129 |
|
0710 | 0710.5628_arXiv.txt | Dust has been detected in the recurrent nova RS\,Ophiuchi on several occassions. I model the historical mid-infrared photometry and a recent Spitzer Space Telescope spectrum taken only half a year after the 2006 eruption. The dust envelope is little affected by the eruptions. I show evidence that the eruptions and possibly the red giant wind of RS\,Oph may sculpt the interstellar medium, and show similar evidence for the recurrent dwarf nova T\,Pyxidis. | \subsection{Asymptotic Giant Branch stars and red supergiants} Stars with an initial mass between $\sim1$ and $\sim40$ M$_\odot$ become hydrogen/helium-shell-burning Asymptotic Giant Branch (AGB) stars or core-helium-burning red supergiants (RSG). These phases are characterised by cool, molecular atmospheres giving rise to M spectral types (S or C for some chemically peculiar AGB stars), strong radial pulsations of this atmosphere on timescales of a year or more, and a high luminosity ($L>2,000$ L$_\odot$). The pulsation may lift the atmosphere high enough such that dust can condense \citep{BowenWillson1991}. Radiation pressure on the grains then drives a dust wind which, via collisions with molecular hydrogen, drags the gas along with it. At a typical wind speed $v_\infty\sim5$ to 30 km s$^{-1}$ these stars lose mass at rates from $\dot{M}\sim10^{-7}$ to over $10^{-4}$ M$_\odot$ yr$^{-1}$ \citep{vanLoonEtal1999}. In AGB stars this leads to the removal of the mantle and the premature death of the star, leaving the truncated core behind as a cooling white dwarf. The more massive RSGs either explode or evolve back to higher surface temperatures, and it is not yet clear how much mass they will have shed during the preceding RSG stage. Mass-loss rates of cool, luminous stars are most easily estimated from the reprocessed radiation emitted by the dust grains at infrared (IR) wavelengths, which is particularly conspicuous in the $\lambda\sim10$ to 40 $\mu$m region. The main problem with this method is that the dust constitutes only a minor fraction of the total mass in the wind, typically $\sim1$:$200$ \citep{Knapp1985} but poorly known for all but the dustiest stars. Carbon monoxide has strong rotational transitions at mm wavelengths which can be detected in relatively nearby stars, but despite being the most abundant molecule after H$_2$ this too is a trace species of which the mass fraction is uncertain --- and modelling the CO line requires knowledge of the temperature profile throughout the envelope which depends, amongst other things, on the dust content. \subsection{First ascent Red Giant Branch stars} Low-mass stars ($M_{\rm initial}<2$ M$_\odot$) evolve along a first ascent Red Giant Branch (RGB) as hydrogen-shell burning stars with an inert helium core. They reach a maximum luminosity of only $L_{\rm tip}\sim2,000$ L$_\odot$, and it becomes problematic for them to drive a wind. The situation is worsened by the fact that many RGB stars do not pulsate strongly and slowly enough to sufficiently increase the scaleheight, and the dust condensation is therefore unlikely to reach completeness. With a low, uncertain dust:gas mass ratio and possibly higher dilution of the dust envelope as it kinematically decouples from the bulk gas, measured mass-loss rates will tend to be too low. The threshold below which this happens is not well known, partly because of uncertainties in the opacity of the grains as they nucleate and grow: \citet{GailSedlmayr1987} estimate $\dot{M}\gg 10^{-6}$ M$_\odot$ yr$^{-1}$ but \citet{NetzerElitzur1993} place it at $\dot{M}>10^{-7}$ M$_\odot$ yr$^{-1}$. \citet{JudgeStencel1991} show that RGB mass loss must be driven by another mechanism, but that this appears to be similarly efficient as a dust-driven wind. The warmer stars further down the RGB have more prominent chromospheres, which may be accompanied by the generation of Alfv\'en or acoustic waves that provide a possible alternative mechanism for driving a wind. Mass-loss rates estimated from the IR emission for the brightest, dustiest RGB stars are typically $\dot{M}\sim10^{-7}$ to $10^{-6}$ M$_\odot$ \citep{vanLoonEtal2006,OrigliaEtal2007}. Mass-loss rates from $\dot{M}\sim10^{-9}$ to $10^{-6}$ M$_\odot$ yr$^{-1}$ have been determined from the blue-displaced cores of strong optical absorption lines and emission in some of these lines, but these are very uncertain too because of the sensitivity to the precise excitation and ionization conditions in the wind. For the most luminous RGB stars the IR and optical estimates tend to agree within an order of magnitude \citep{McDonaldvanLoon2007}. \begin{figure}[!t] \plotone{Jacco_vanLoon_fig1.eps} \caption{RS\,Oph compared to isochrones from \citet{GirardiEtal2000}.} \end{figure} | Dust in RS\,Oph has been detected in between several eruptions, and at only half a year since the 2006 eruption. The silicate dust species and mass-loss rate of $\dot{M}\sim2$ to $3\times10^{-8}$ M$_\odot$ yr$^{-1}$ are not at odds with the expectations for well-developed dust-driven winds of single stars. But because RS\,Oph has a relatively warm photosphere, weak pulsation and not very high luminosity, this might in fact suggest that the mass-loss rate {\it is} enhanced by the effects of the companion white dwarf. The eruptions do {\it not} seem to have a dramatic effect on the dust envelope. An accretion rate onto the white dwarf surface of $\dot{M}>10^{-7}$ M$_\odot$ yr$^{-1}$ would probably require an enhanced flow through the Lagrangian point L$_1$. Considering the pattern in the time intervals between historic eruptions, including the 1907 eruption \citep{Schaefer2004}, one might heuristically expect the next eruption to be due around 2015 (Fig.\ 10). \begin{figure}[!b] \plotone{Jacco_vanLoon_fig10.eps} \caption{Recent outbursts of RS\,Oph and a heuristic prediction for the next.} \end{figure} | 7 | 10 | 0710.5628 |
0710 | 0710.2839_arXiv.txt | {% Neutron stars contain persistent, ordered magnetic fields that are the strongest known in the Universe. However, their magnetic fluxes are similar to those in magnetic A and B stars and white dwarfs, suggesting that flux conservation during gravitational collapse may play an important role in establishing the field, although it might also be modified substantially by early convection, differential rotation, and magnetic instabilities. The equilibrium field configuration, established within hours (at most) of the formation of the star, is likely to be roughly axisymmetric, involving both poloidal and toroidal components. The stable stratification of the neutron star matter (due to its radial composition gradient) probably plays a crucial role in holding this magnetic structure inside the star. The field can evolve on long time scales by processes that overcome the stable stratification, such as weak interactions changing the relative abundances and ambipolar diffusion of charged particles with respect to neutrons. These processes become more effective for stronger magnetic fields, thus naturally explaining the magnetic energy dissipation expected in magnetars, at the same time as the longer-lived, weaker fields in classical and millisecond pulsars.} | The purpose of this talk is to present and discuss some of the physical processes that are likely to be relevant in determining the structure and evolution of magnetic fields in neutron stars, almost regardless of their (so far largely unknown) internal composition and state of matter. For this reason, I do not discuss the fascinating, exotic issues of quark matter, superfluidity, superconductivity, and the like, but emphasize the much more pedestrian concepts of stable stratification and non-ideal magnetohydrodynamics (MHD) processes such as ambipolar diffusion. A review with a very different focus has recently been given by Geppert (2006). | The extremely strong magnetic fields found in neutron stars are nevertheless weak enough to be balanced by small perturbations in the density, pressure, and chemical composition of the stably stratified, degenerate, multi-species fluid found in their interior. The field likely originates from the fairly strong magnetic flux inherited from the progenitor star, possibly modified by a combination of convection, differential rotation, and magnetic instabilities acting during the short, protoneutron star stage following collapse. The equilibrium field set up during this stage likely involves linked toroidal and poloidal field components. It can decay through dissipative processes such as weak interactions and ambipolar diffusion, which change the chemical composition of the matter, allowing it to move. These processes become particularly effective at high field strengths, possibly accounting for the energy release inside magnetars, which also keeps the stellar interior hot enough for the magnetic field to rearrange substantially. This forces changes as well in the crust of the star, which can be broken by strong fields, whereas weaker ones can evolve by a combination of Hall drift and resistive dissipation. The internal dissipation processes are ineffective for pulsar-strength fields, which can easily survive for a pulsar lifetime. None of these processes depends essentially on the exotic properties of a neutron star, such as Cooper pairing or quark matter, which can however modify the time scales for these processes to occur. | 7 | 10 | 0710.2839 |
0710 | 0710.5244_arXiv.txt | In this paper we present a deep and homogeneous i-band selected multi-waveband catalogue in the COSMOS field covering an area of about $0.7\sq\degr$. Our catalogue with a formal 50\% completeness limit for point sources of $i\sim 26.7$ comprises about 290~000 galaxies with information in 8 passbands. We combine publicly available u, B, V, r, i, z, and K data with proprietary imaging in H band. We discuss in detail the observations, the data reduction, and the photometric properties of the H-band data. We estimate photometric redshifts for all the galaxies in the catalogue. A comparison with 162 spectroscopic redshifts in the redshift range $ 0 \lsim z \lsim 3$ shows that the achieved accuracy of the photometric redshifts is \mbox{$\Delta z / (z_{spec}+1) \lsim 0.035$} with only $\sim 2$\% outliers. We derive absolute UV magnitudes and investigate the evolution of the luminosity function evaluated in the restframe UV (1500~\AA). { There is a good agreement between the LFs derived here and the LFs derived in the FORS Deep Field. We see a similar brightening of M$^\ast$ and a decrease of $\phi^\ast$ with redshift.} The catalogue including the photometric redshift information is made publicly available. | \label{sec:intro} In the last decade our knowledge about the evolution of global galaxy properties over a large redshift range has improved considerably. The 2dF Galaxy Redshift Survey (2dFGRS; \citealt{colless:1}), the Sloan Digital Sky Survey (SDSS; \citealt{stoughton:1}), and the 2MASS survey \citep{2MASS} have provided very large local galaxy samples with spectroscopic and/or photometric information in various passbands. Thanks to these data sets we are now able to assess very accurate local ($z\sim 0.1$) reference points for many galaxy evolution measurements like the luminosity function, the star formation activity, the spatial clustering of galaxies, the stellar population, the morphology, etc. In the redshift range between $0.2 \lsim z \lsim 1 $ pioneering work has been done in the context of the Canada France Redshift Survey \citep{lilly:3}, the Autofib survey \citep{ellis:1} and in the Canadian Network for Observational Cosmology survey \citep{yee:1}. They provide accurate distances and absolute luminosities by spectroscopic followup of optically selected galaxies, thus being able to probe basic properties of galaxy evolution. Moreover the K20-survey \citep{cimatti:2} as well as the MUNICS survey \citep{drory:2,feulner:1} extend the analysis into the near infrared regime (for $0.2 \lsim z \lsim 1.5 $). An important step towards probing the galaxy properties also in the high redshift regime around $z \sim 3$ and $z\sim 4$ was the work of \citet{steidel_lbg:1} and \citet{steidel_lbg:2}. They used colour selection to discriminate between low and high redshift galaxies \citep[see ][for a review]{giavalisco:3}. The so-called Lyman-break galaxies (LBGs, mainly starburst galaxies at high redshift) are selected by means of important features in the UV spectrum of star-forming galaxies. The next milestones in pushing the limiting magnitude for detectable galaxies to fainter and fainter limits were the space based Hubble Deep Field North (HDFN; \citealt{HDF96}) and Hubble Deep Field South \citep[HDFS; ][]{HDFS00,HDFS00a} \citep[see ][for a review]{ferguson:1}. Although of a limited field of view of about $5\sq\arcmin$ only, the depth of the HDFs allowed the detection of galaxies up to a redshift of 5 and even beyond. In the past years the space based HDFs were supplemented by many more multi-band photometric surveys like the NTT SUSI deep Field (NDF; \citealt{arnouts_ntt}), the Chandra Deep Field South (CDFS; \citealt{arnouts_cdfs}), the William Herschel Deep Field (WHDF; \citealt{mccracken:1,metcalf:1}), the Subaru Deep Field/Survey (SDF; \citealt{maihara:1,ouchi:2}), the COMBO-17 survey \citep{combo17:1}, FIRES \citep{labbe:1}, the FORS Deep Field (FDF; \citealt{fdf_data}), the Great Observatories Origins Deep Survey (GOODS; \citealt{giavalisco:1}), the Ultra Deep Field (UDF and UDF-Parallel ACS fields; \citealt{giavalisco:2,bunker:1, bouwens:2004}), the VIRMOS deep survey \citep{lefevre:3}, GEMS \citep{rix:1}, the Keck Deep Fields \citep{sawicki:2}, and the Multiwavelength Survey by Yale-Chile (MUSYC; \citealt{gawiser:1, quadri:1}). With the advent of all these deep multi-band photometric surveys the photometric redshift technique (essentially a generalisation of the drop-out technique) can be used to identify high-redshift galaxies. Photometric redshifts are often determined by means of template matching algorithm that applies Bayesian statistics and uses semi-empirical template spectra matched to broad-band photometry (see also \citealt{baum:1, koo:1, brunner:1, Soto:1, benitez:1, bender:1, borgne:1, firth:1}). Redshifts of galaxies that are several magnitudes fainter than typical spectroscopic limits can be determined reliably with an accuracy of \mbox{$\Delta z / (z_{spec}+1)$} of 0.02 to 0.1. In this context the COSMOS survey (\citet{scoville:1}; see also http://www.astro.caltech.edu/cosmos/ for an overview) combines deep to very-deep multi-waveband information in order to extend the analysis of deep pencil beam surveys to a much bigger volume, thus being able to drastically increase the statistics and detect also very rare objects. For this, the survey covers an area of about $2\sq\degr$ with imaging by space-based telescopes (Hubble, Spitzer, GALEX, XMM, Chandra) as well as large ground based telescopes (Subaru, VLA, ESO-VLT, UKIRT, NOAO, CFHT, and others). In this paper we combine publicly available u, B, V, r, i, z, and K COSMOS data with proprietary imaging in the H band to derive a homogeneous multi-waveband catalogue suitable for deriving accurate photometric redshifts. In Section~\ref{sec:imaging_data} we give an overview of the near-infrared (NIR) data acquisition and we describe our 2-pass data reduction pipeline used to derive optimally (in terms of signal-to-noise for faint sources) stacked images in Section~\ref{sec:imaging_reduction}. We also present NIR galaxy number counts and compare them with the literature.\\ In Section~\ref{sec:isel_catalogue} we present the deep multi-waveband i-band selected catalogue and discuss its properties, whereas the data reduction of the spectroscopic redshifts is described in Section~\ref{sec:spec}. In Section~\ref{sec:photoz} we present the photometric redshift catalogue, discuss the accuracy of the latter and show the redshift distribution of the galaxies. In Section~\ref{sec:uvlf} we derive the redshift evolution of the restframe UV luminosity function and luminosity density at 1500~\AA\ from our i-selected catalogue before we summarise our findings in Section~\ref{sec:summary_conclusion}. We use AB magnitudes and adopt a $\Lambda$ cosmology throughout the paper with \mbox{$\Omega_M=0.3$}, \mbox{$\Omega_\Lambda=0.7$}, and \mbox{$H_0=70 \, \mathrm{km} \, \mathrm{s}^{-1} \, \mathrm{Mpc}^{-1}$}. | \label{sec:summary_conclusion} In this paper we present the data acquisition and reduction of NIR Js, H, and K' bands in the COSMOS field. We describe a 2-pass reduction pipeline to reduce NIR data. The 2-pass pipeline is optimised to avoid flat-field errors introduced if the latter are constructed from science exposures. Moreover we present and implement a method to stack images of different quality resulting in an optimal S/N ratio for faint sky dominated {point sources}. The Js and K' band cover an area of about \mbox{$200\sq\arcmin$} (1 patch) whereas the H band covers about $0.85\sq\degr$ (15 patches) in total. The 50\% completeness limits are 22.67, $\sim 21.9$, and 21.76 in the Js, H, and K' band, respectively. The number counts of all NIR bands nicely agree with the number counts taken from literature. Furthermore we present a deep and homogeneous i-band selected multi-waveband catalogue in the COSMOS field by combining publicly available u, B, V, r, i, z, and K bands with the H band. The clean catalogue with a formal 50\% completeness limit for point sources of $i\sim 26.7$ comprises about 290~000 galaxies with information in 8 passbands and covers an area of about $0.7\sq\degr$ (12 patches). We exclude all objects with corrupted magnitudes in only one of the filters from the catalogue in order to have a catalogue as homogeneous as possible. Photometric redshifts for all objects are derived and a comparison with 162 spectroscopic redshifts in the redshift range $ 0 \lsim z \lsim 3$ shows that the achieved accuracy of the photometric redshifts is \mbox{$\Delta z / (z_{spec}+1) \lsim 0.035$} with only $\sim 2$\% outliers. Please note that in order to break the degeneracy between high redshift and low redshift solutions we included also the GALEX FUV and NUV filters in the photometric redshift estimation which considerably reduced the number of outliers. The multi-waveband catalogue including the photometric redshift information is made publicly available. The data can be downloaded from {\verb http://www.mpe.mpg.de/~gabasch/COSMOS/ } We derive absolute UV magnitudes and a comparison in a magnitude-redshift diagram with the FDF shows good agreement. Moreover we investigate the evolution of the luminosity function evaluated in the restframe UV (1500~\AA). We find a substantial brightening of M$^\ast$ and a decrease of $\phi^\ast$ with redshift: from \mbox{$\langle z \rangle\sim 0.5$} to \mbox{$\langle z \rangle\sim 4.5$} the characteristic magnitude increases by about 3 magnitudes, whereas the characteristic density decreases by about 80 -- 90\%. We compare the redshift evolution of the UV luminosity density in the COSMOS field and the FDF up to a redshift of $z\sim 5$. Below a redshift of $z\sim 2.5$ the mean UV luminosity density in COSMOS is systematically higher by about $1\sigma$ if compared to the FDF. At $2.5 \lsim z \lsim 5$ both UV luminosity densities agree very well. It is worth noting the remarkably good agreement between the UV LF as well as the UV LD despite the fact, that the FDF is about 60 times smaller than the COSMOS field analysed here. | 7 | 10 | 0710.5244 |
0710 | 0710.2464_arXiv.txt | {} {We studied the temporal and spectral evolution of the synchrotron emission from the high energy peaked BL~Lac object \1e.} {Two recent observations have been performed by the \xmm\,\,and \sw\,satellites; we carried out X-ray spectral analysis for both of them, and photometry in optical-ultraviolet filters for the \sw\,\,one. Combining the results thus obtained with archival data we built the long-term X-ray light curve, spanning a time interval of 26~years, and the Spectral Energy Distribution (SED) of this source.} {The light curve shows a large flux increasing, about a factor of six, in a time interval of a few years. After reaching its maximum in coincidence with the \xmm\,\,pointing in December~2000 the flux decreased in later years, as revealed by \sw\,. The very good statistics available in the 0.5-10 keV \xmm\,\,X-ray spectrum points out a highly significant deviation from a single power law. A log-parabolic model with a best fit curvature parameter of 0.25 and a peak energy at $\sim$~1~keV describes well the spectral shape of the synchrotron emission. The simultaneous fit of \sw\,\,UVOT and XRT data provides a milder curvature ($b\sim0.1$) and a peak at higher energies ($\sim15$~keV), suggesting a different state of source activity. In both cases UVOT data support the scenario of a single synchrotron emission component extending from the optical/UV to the X-ray band.} {New X-ray observations are important to monitor the temporal and spectral evolution of the source; new generation $\gamma$-ray telescopes like AGILE and GLAST could for the first time detect its inverse Compton emission.} | BL~Lac objects are thought to be radio-loud Active Galactic Nuclei (AGNs) observed in a direction very close to the axis of a relativistic jet outflowing from the inner nuclear region (\cite{Urry95}). This interpretation could explain most of the characteristics of these sources like compact and flat-spectrum radio emission, superluminal motion revealed by VLBI imaging, high and variable radio and optical polarization, non-thermal continuum emission extending from radio to $\gamma$-ray frequencies, an almost featureless optical spectrum and the fast variability at all frequencies. BL~Lac objects are generally characterized by a double bump structure in the broad band Spectral Energy Distribution (SED). The low frequency bump is attributed to synchrotron radiation emitted by relativistic electrons in the jet; inverse Compton scattering by the same electron population on the synchrotron radiation is thought to be at the origin of the high frequency bump. The peak of the first bump may vary in a rather wide range of frequencies: from the IR/optical band for the low energy peaked BL~Lacs (LBLs) to the UV/X-ray band for the high energy peaked BL~Lacs (HBLs) (\cite{Giommi94}, \cite{Padovani}). \1e, also named BZB J1210+3929 in the recent Multifrequency Catalogue of Blazars (\cite{BZcat}) is one of the X-ray selected BL~Lacertae objects of the \ein\,\,Medium Sensitivity Survey (\cite{EMSS}). It was discovered as a serendipitous source located about five arc-minutes north of one of the most intensively studied AGNs, the bright Seyfert galaxy NGC~4151. For this reason \1e has been observed on many occasions by all the imaging X-ray instruments that have operated since the \ein\,\,observatory. Despite its relatively high redshift (z=0.615) HST was able to detect the bright (M$_{\rm R}$\,=\,$-$24.4) host galaxy which is of elliptical type (\cite{Scarpa00}). In this paper we report the long-term X-ray light curve which spans over 26 years. We also report and compare the spectral analysis of the most recent observations carried out by the \xmm\,\,(\cite{jansen}) and \sw\,\,(\cite{gehrels}) satellites; we discuss the possibility to model the synchrotron emission of this source with a single log-parabolic model. The peak of the synchrotron component lies in the X-ray band, and for this reason the source can safely be classified as an HBL. We finally report the Spectral Energy Distribution (SED) of \1e compiled from non-simultaneous multi-frequency archival data. | \label{Discussion} Adding to multi-frequency archival data the results of our analysis we built the radio to X-ray SED of \1e (Fig.\ref{SED}) which follows the characteristic trend of extreme HBL sources (\cite{Padovani}) with the synchrotron emission covering the entire frequency interval and peaking well into the X-ray band. The inverse Compton component is not detected and therefore is expected at higher energies. The \sw\,\,satellite is sensitive both in the X-ray and in the optical-ultraviolet band. This provides the opportunity to test if the Spectral Energy Distribution of a source like \1e can be explained by a single synchrotron component, or if multiple emission components are present. For this purpose we first fitted a log-parabolic model to simultaneous \sw\,\,XRT and UVOT fluxes (solid line in Fig.~\ref{combo}) and obtained a rather low value for the curvature parameter ($b\sim0.1$) in contrast with the one obtained by fitting XRT data only, as reported in Table~\ref{tab1}; the energy peak lies at about 15~keV or more. Comparing these results with those from \xmm, \sw\,\,evidently caught \1e in a different state of activity, characterized by a milder curvature and a peak of the synchrotron component shifted to higher energies. Anyway firm conclusions cannot be drawn due to a relatively poor statistics and to the fact that the found $E_p$ value likely lies at energies higher than 10~keV, outside the range of \sw\,\,XRT. Aware of the limits of using non simultaneous observations for a variable source, yet we veri\-fied if \sw\,\,UVOT and \xmm\,\,MOS1 data are compatible with a log-parabolic model. We took points representing XSPEC model of observation~I and fit them on a wider frequency interval (dot-dashed line in Fig.\,\ref{combo}): UVOT points are sistematically above the extrapolation of the log-parabolic model. We took care to compare this result with the one obtained by fitting directly X-ray rebinned data and obtained an almost coincident result. Then we fitted both data sets with a log-parabola (dotted line in Fig.\,\ref{combo}) and estimated the curvature parameter: we obtained $b=0.18\pm0.01$, a value at about 2$\sigma$ of the one derived fitting only \xmm\,\,data (see Table \ref{tab1}). At this point we tested the log-parabolic model with $b$ fixed at 0.18 and found $\chi^2_r/d.o.f. = 1.02/376$, a value practically coincident with the one obtained leaving $b$ free to vary, with only a small difference in the $a$ parameter. This result encouraged us in concluding that a single synchrotron component is in a good agreement with the emission observed. \1e is certainly worth monitoring in the next years. New observations in the X-ray band would add essential information to the light curve behaviour: if any kind of regularity in flux increasing and fading should emerge, it would be possible to put physical constraints to establish the nature of the mechanism responsible for these variations. Moreover, observations with instruments which extend to higher energies than XRT, like those on board \textit{Suzaku}, would extend the known Spectral Energy Distribution and would allow us to obtain a better parametrization of the spectral curvature. Finally, \1e may be an interesting target for AGILE and the forthcoming GLAST $\gamma$-ray mission which could detect for the first time the inverse Compton component of this HBL source. | 7 | 10 | 0710.2464 |
0710 | 0710.0182_arXiv.txt | Using high-resolution N-body simulations, we examine whether a major dry merger mitigates the difference in the radial density distributions between red and blue globular clusters (GCs). To this end, we study the relation between the density slope of the GCs in merger progenitors and that in a merger remnant, when the density distribution is described by $n_{\rm GC}\propto r^{-\alpha}$. We also study how our results depend on the merger orbit and the size of the core radius of the initial GC density distribution. We find that a major dry merger makes the GC profile flatter, and the steeper initial GC profile leads to more significant flattening, especially if the initial slope is steeper than $\alpha\sim3.5$. Our result suggests that if there is a major dry merger of elliptical galaxies whose red GCs have a steeper radial profile than the blue GCs, as currently observed, and their slopes are steeper than $\alpha\sim3.5$, the difference in the slopes between two populations becomes smaller after dry mergers. Therefore, the observed slopes of red and blue GCs can be a diagnostic of the importance of dry merger. The current observational data show that the red and blue GCs have more comparable and shallower slopes in some luminous galaxies, which may indicate that they have experienced dry mergers. | Globular clusters (GCs) in elliptical galaxies have been intensively studied in consideration of explaining the formation of both their host elliptical galaxies and GCs themselves \citep[see][for a review]{araa06}. GCs are attractive as tracers of the star formation history of their host galaxies \citep[e.g.][]{yi04,strader06}, because some properties of GC systems are correlated with the properties of their host galaxies \citep[e.g.,][]{brodie91,djorgovski92}. It is thought that the formation of GCs is triggered by starburst accompanying gas-rich galaxy merging \citep[e.g.][]{schweizer87,ashman92} or starburst that might happen with multiple dissipational collapses \citep{forbes97}. Forming young star clusters are found and they are expected to become star clusters like current old GCs in local galaxy mergers \citep{schweizer06}. An important aspect of GC systems in elliptical galaxies is a color bimodality \citep[e.g.,][]{zepf93,geisler96,gebhardt99,larsen01,peng06}. It is also found that red GCs are more centrally concentrated than blue GCs \citep[e.g.][]{forte05,bassino06b, tamura06}, which could put additional constraints on their formation scenario \citep{bekki02}. The radial profile of each GC subpopulation is well described by a power-law distribution, especially at the outer radii, and the red GCs have a steeper slope than the blue GCs. Although the origin of this color bimodality is still uncertain, it is probably closely related to the formation history of their host elliptical galaxies \citep[e.g.,][]{yoon06,strader07,kundu07}. Classically, three scenarios have been proposed to explain the color bimodality; major gas-rich mergers, in situ formation of multiple dissipational collapses, and dissipationless accretion. An explanation suggested by \citet{ashman92} is that red GCs are metal-rich clusters which might be formed by gas-rich disk-disk mergers. Therefore, the red GCs might be younger than blue GCs. Meanwhile, \citet{forbes97} explain that blue GCs might be formed in the first stage of dissipational collapse and red GCs might be formed after the truncation of the blue GC formation. Another explanation given by \citet{cote98} includes accretion of blue GCs from small galaxies into the already formed red GCs. More recently, \citet{beasley02} demonstrate that the bimodality can be explained in elliptical galaxy formation based on a hierarchical clustering scenario. On the other hand, recent research suggests that in the late stage of evolution, early-type galaxies might have experienced dry merging where merger progenitors do not have much gas, nor accompany star formation. The number density evolution of red galaxies has been discussed in observations of COMBO-17 \citep{bell04} and DEEP2 surveys \citep{faber05}, and such studies suggest that the density change can be understood by the dominance of dry merging after z $\simlt$ 1 \citep[but see also][]{yamada05,cimatti06,bundy07,scarlata07}. The evolution of galaxy clustering also implies late effects on the evolution of massive red galaxies from dry merging \citep{white07}. Moreover, the observations show that dry merging does occur \citep{vandokkum05,tran05,rines07}, while the observed features of galaxies are well explained in cosmological simulations \citep[e.g.][]{kawata06}. Recent theoretical studies of dry merger simulations of ellipticals show that merger remnants maintain their properties on the fundamental plane and other scaling relations \citep[e.g.][]{nipoti03,boylan05,robertson06,ciotti07}. The dry merging of binary ellipticals also can explain the formation of boxy-type ellipticals \citep{naab06}. \citet{bekki06a} demonstrate that the observed correlation between a spatial distribution of GCs and the total luminosity of ellipticals can be explained by sequential dissipationless major mergers, because the radial density profile of GCs progressively flatten after each major dry merger. This study raises the important question of whether or not the slopes of the density profiles of red and blue GCs persist after major dry merging. For example, we now consider that the density profiles of red and blue GCs in progenitor elliptical galaxies are described by $n_{\rm GC, red}\propto r^{-\alpha_{\rm p,red}}$ and $n_{\rm GC, blue}\propto r^{-\alpha_{\rm p,blue}}$, and these profiles in a major dry merger remnant become $n_{\rm GC, red}\propto r^{-\alpha_{\rm r,red}}$ and $n_{\rm GC, blue}\propto r^{-\alpha_{\rm r,blue}}$. The current observations suggest that $\alpha_{\rm r,red}>\alpha_{\rm r,blue}$. However, if a dry merger flattens the red GC density profile more than the blue GC density profile, $\alpha_{\rm p,red}-\alpha_{\rm r,red}>\alpha_{\rm p,blue}-\alpha_{\rm r,blue}$, in the remnant galaxy the difference between $\alpha_{\rm r,red}$ and $\alpha_{\rm r, blue}$ becomes smaller. In this case, the observed difference in the slopes of the density profiles of red and blue GCs in nearby ellipticals can be a valuable diagnostic for the importance of the dry merger in the evolution of ellipticals. To clarify this issue, the question becomes how the flattening of GC profiles during dry merging depends on the initial distributions of the GCs. We use numerical simulations of major dry mergers to study the dependence of GC distributions in merger remnants on the initial distributions in merger progenitors. Then, we can compare $\alpha_{\rm p,red}-\alpha_{\rm r,red}$ and $\alpha_{\rm p,blue}-\alpha_{\rm r,blue}$, for different sets of $\alpha_{\rm p,red}$ and $\alpha_{\rm p,blue}$. Since major dry mergers between two equal-mass merger progenitors must leave the most significant effects on merger remnants, compared with minor mergers, we study only equal-mass mergers in this paper. In \S2, we explain details of our dry merger simulations and initial GC distributions. The changes in GC spatial distributions due to major mergers are shown in \S3. We discuss the implication of our results in \S4 that is followed by conclusion. | Our main result is that steeper initial GC profiles experience stronger flattening as shown in Figure \ref{fig:alpha}. $\alpha_{\rm p} \approx 3.5$ is in boundary between strong and weak flattening. The results imply that if the initial slopes of both red and blue GCs are steeper than $\alpha_{\rm p} \approx 3.5$, the difference in the slopes between two populations of GCs will become much smaller, independent of merger orbits. In particular, even only one dry merger can make the slope flatter dramatically. Moreover, it also makes both slopes be around $\alpha_{\rm r} \approx 3.5 - 4.5$. For example, if $\alpha_{\rm p,red}=5$ and $\alpha_{\rm p, blue}=4$, Figure \ref{fig:alpha} suggests that $\alpha_{\rm r,red} \approx 4.3$ and $\alpha_{\rm p, blue} \approx 3.6$. The difference in the slope between red and blue GCs becomes significantly small. Therefore, the difference in the slopes of red and blue GCs can be a constraint of the number of major dry mergers. On the other hand, if the initial GC distribution has a slope shallower than $\alpha_{\rm p} \approx 3.5$, the slope changes very little, and almost no change in the case of $\alpha_{\rm p} \approx 2$. Therefore, if the initial distributions of red and blue GCs in merger progenitors are steeper than $\alpha_{\rm p} = 2$, and the galaxies experienced a number of major dry mergers, it leads the distributions of both red and blue GCs to become close to $\alpha_{\rm r} \sim 2$ for both red and blue GCs, and the difference in the slope of spatial distributions becomes difficult to be measured. It is also worth noting that if the initial core size is larger for blue GCs than for red GCs, but their initial slopes are the same, a major dry merger can make the slope for the blue GCs shallower than for the red GCs, as shown in Figure \ref{fig:r_c_dependence1}. Therefore, the final slope is not a simple function of the initial slope. Our results suggest that dry mergers make the slopes of the density profiles for red and blue GCs shallower and similar. Several studies \citep[e.g.][]{kissler97,vandenbergh98,forbes05,lauer07, emsellem07} claim that the various observed properties for ellipticals show a transition around ${\rm M_{V} \sim -21}$, i.e. ${\rm \sim 10^{11}\ M_{\sun}}$. Some of these bimodalities have been interpreted in the picture of dry merger hypothesis for the growth of massive ellipticals \citep[e.g.][]{capetti06}. For example, \citet{emsellem07} show that slowly rotating ellipticals might be mainly affected by dry mergers. The slowly rotating ellipticals are more luminous than ${\rm M_{B} \approx -20.5}$, while less luminous ellipticals are fast rotators. It may indicate that less luminous galaxies have not experienced any major dry mergers. Therefore, it is interesting to examine the observed distributions of red and blue GCs as a function of the stellar mass of ellipticals. If more luminous galaxies have experienced dry mergers, our results predict that the slopes for red and blue GCs in bright galaxies are shallower and similar. Unfortunately, the current observational samples are not enough to test it statistically. Below we provide some discussion based on the current limited measurements. NGC 1399 is one of ellipticals whose spatial distributions of red and blue GCs are well-observed \citep{dirsch03,bassino06b}. When assuming ${\rm M/L_{V} = 5}$ and $B-V=0.9$ for an elliptical galaxy, as used in \citet{forbes05}, the stellar mass of NGC 1399 is about ${\rm 5 \times\ 10^{11}\ M_{\odot}}$ (${\rm M_{B} = -21.8}$) \citep{araa06}. The derived power-law indices of projected radial density profiles are $1.9\pm0.06$ and $1.6\pm0.10$ for red and blue GCs, respectively \citep{bassino06b}. If the distributions are spherically symmetric, the three-dimensional radial distributions have power-law slopes of $\alpha \approx 2.9$ and $\approx 2.6$ for red and blue GCs, respectively. Note that the real slope might be slightly steeper, because the distribution of GCs is not infinite. The slope difference between red and blue GC distributions is $\Delta \alpha \approx 0.3$ while both GC populations show $\alpha \approx 3$. In addition to NGC 1399, NGC 1407 is a brightest group galaxy with ${\rm M_V=-21.86}$, and has a bimodal color distribution of GCs. \citet{forbes06} report that the projected slopes of the red and blue GC density profiles are respectively $1.50\pm0.05$ and $1.65\pm0.29$, i.e., expected three dimensional slopes of $\alpha_{\rm red}=2.50$ $\alpha_{\rm blue}=2.65$, which are statistically the same as each other. On the other hand, NGC 1374 and NGC 1379 are less luminous than NGC 1399, having ${\rm M_{V} = -20.4}$ and $-20.6$, respectively. The red and blue GCs of the two galaxies are studied in \citet{bassino06a}. The derived projected density distributions of red and blue GCs in NGC 1374 have slopes of 3.2 and 2.3. Hence, the expected slopes in their three-dimensional density profiles are $\alpha_{\rm red}\sim 4.2$ and $\alpha_{\rm blue}\sim3.3$. The $\Delta \alpha$ in NGC 1379 is $\sim 0.6$, and the red and blue GCs have projected power-law slopes of $\sim 2.9$ and $\sim 2.3$, respectively. The low-luminosity galaxies have a larger difference in the density slopes of the red and blue GCs, and their slopes are steeper than the more luminous galaxies discussed above. Combining with our results, these four sample suggests that luminous galaxies, like NGC 1399 and NGC 1407, may have experienced some dry mergers, while less luminous galaxies, such as NGC 1374 and NGC 1379, may have not been formed through dry mergers, which is consistent with a scenario that the significance of dry mergers depends on the stellar mass. However, we also find some giant galaxy that does not follow the above trend. Recently, \citet{tamura06} measured the density profiles of red and blue GCs for giant ellipticals: M87 and NGC 4552. We have fitted the projected density profile in Figure 5 of their paper by a power-law profile, and the derived power-law index is 2.4 and 1.4 for red and blue GCs in M87 and 1.8 and 1.2 for red and blue GCs in NGC 4552, respectively. Although these giant ellipticals have relatively shallower slopes, the difference in slopes of red and blue GCs is significant. According to our result, these giant ellipticals are unlikely to have experienced significant dry mergers. Some giants might form without dry mergers, or some other factor, such as a number of minor mergers, might affect the slopes of GCs. Again, so far there are not many studies which derive the density distributions of both red and blue GCs. This problem limits the application of our results to the current observational data. However, we expect that future observations will improve the statistics for nearby ellipticals, and our results will add an important framework to interpret the galaxies' evolutionary history. It will be also valuable to compare spatial distributions of blue and red GCs with other expected properties from dry mergers such as surface brightness profile. | 7 | 10 | 0710.0182 |
0710 | 0710.5857_arXiv.txt | The classical picture of GUT baryogenesis has been strongly modified by theoretical progress concerning two nonperturbative features of the standard model: the phase diagram of the electroweak theory, and baryon and lepton number changing sphaleron processes in the high-temperature symmetric phase of the standard model. We briefly review three viable models, electroweak baryogenesis, the Affleck-Dine mechanism and leptogenesis and discuss the prospects to falsify them. All models are closely tied to the nature of dark matter, especially in supersymmetric theories. In the near future results from LHC and gamma-ray astronomy will shed new light on the origin of the matter-antimatter asymmetry of the universe. | The cosmological matter-antimatter asymmetry can be dynamically generated if the particle interactions and the cosmological evolution satisfy Sakharov's conditions \cite{sak67}, \begin{itemize} \item baryon number violation, \item $C$ and $C\!P$ violation, \item deviation from thermal equilibrium. \end{itemize} Although the baryon asymmetry is just a single number, it provides an important connection between particle physics and cosmology. In his seminal paper, 40 years ago, Sakharov not only stated the necessary conditions for baryogenesis, he also proposed a specific model. The origin of the baryon asymmetry were $C\!P$ violating decays of superheavy `maximons' with mass $\mathcal{O}(M_\mathrm{P})$ at an initial temperature $T_i \sim M_\mathrm{P}$. The $C\!P$ violation in maximon decays was related to the $C\!P$ violation observed in $K^0$-decays, and the violation of baryon number led to a proton lifetime $\tau_p > 10^{50}\ \mathrm{years}$, much larger than current estimates in grand unified theories. \begin{figure}[t] \includegraphics[height=6cm]{fig3v}\hspace{1cm} \includegraphics[height=6cm,width=8cm]{pot} \caption{{\it Left:} Critical temperature $T_c$ of the electroweak transition as function of $R_{HW}=m_H/m_W$; from \cite{cfh98}. {\it Right:} Effective potential of the Higgs field $\varphi$ at temperature $T>T_c$. \label{fig:ew}} \end{figure} At present there exist a number of viable scenarios for baryogenesis. They can be classified according to the different ways in which Sakharov's conditions are realized. In grand unified theories baryon number ($B$) and lepton number ($L$) are broken by the interactions of gauge bosons and leptoquarks. This is the basis of classical GUT baryogenesis (cf.~\cite{kt90}). In a similar way, lepton number violating decays of heavy Majorana neutrinos lead to leptogenesis \cite{fy86}. In the simplest version of leptogenesis the initial abundance of the heavy neutrinos is generated by thermal processes. Alternatively, heavy neutrinos may be produced in inflaton decays or in the reheating process after inflation. Because in the standard model baryon number, $C$ and $C\!P$ are not conserved, in principle the cosmological baryon asymmetry can also be generated at the electroweak phase transition \cite{krs85}. A further mechanism of baryogenesis can work in supersymmetric theories where the scalar potential has approximately flat directions. Coherent oscillations of scalar fields can then generate large asymmetries \cite{ad85}. The theory of baryogenesis crucially depends on nonperturbative properties of the standard model, first of all the nature of the electroweak transition. A first-order phase transition yields a departure from thermal equilibrium. Fig.~1 shows the phase diagram of the electroweak theory, i.e. the critical temperature in units of the Higgs mass, $T_c/m_H$, as function of the Higgs mass in units of the W-boson mass, $R_{HW}=m_H/m_W$ \cite{cfh98,lr98}. For small Higgs masses the phase transition is first-order; above a critical Higgs mass, $m_H > m_H^c \simeq 72$~GeV, it turns into a smooth crossover \cite{bp94,klx96}. This upper bound for a first-order transition has to be compared with the lower bound from LEP, $m_H > 114$~GeV. Hence, there is no departure from thermal equilibrium at the electroweak transition in the standard model. The second crucial nonperturbative aspect of baryogenesis is the connection between baryon number and lepton number in the high-temperature, symmetric phase of the standard model. Due to the chiral nature of the weak interactions $B$ and $L$ are not conserved \cite{tho76}. At zero temperature this has no observable effect due to the smallness of the weak coupling. However, as the temperature reaches the critical temperature $T_c$ of the electroweak phase transition, $B$ and $L$ violating processes come into thermal equilibrium \cite{krs85}. The rate of these processes is related to the free energy of sphaleron-type field configurations which carry topological charge. In the standard model they lead to an effective interaction of all left-handed fermions \cite{tho76} (cf.~Fig.~2), \begin{equation} O_{B+L} = \prod_i \left(q_{Li} q_{Li} q_{Li} l_{Li}\right)\; , \end{equation} which violates baryon and lepton number by three units, \begin{equation} \Delta B = \Delta L = 3\;. \label{sphal1} \end{equation} The sphaleron transition rate in the symmetric high-temperature phase has been evaluated by combining an analytical resummation with numerical lattice techniques \cite{bmr00}. The result is, in accord with previous estimates, that $B$ and $L$ violating processes are in thermal equilibrium for temperatures in the range \begin{equation} T_{EW} \sim 100\ \mbox{GeV} < T < T_{SPH} \sim 10^{12}\ \mbox{GeV}\;. \end{equation} \begin{figure}[t] \includegraphics[height=6cm]{lg1} \caption{One of the 12-fermion processes which are in thermal equilibrium in the high-temperature phase of the standard model. \label{fig:sphal} } \end{figure} Sphaleron processes have a profound effect on the generation of the cosmological baryon asymmetry. An analysis of the chemical potentials of all particle species in the high-temperature phase yields the following relation between the baryon asymmetry and the corresponding $L$ and $B-L$ asymmetries, \begin{equation}\label{basic} \langle B\rangle_T = c_S \langle B-L\rangle_T = {c_S\over c_S-1} \langle L\rangle_T\;. \end{equation} Here $c_S$ is a number ${\cal O}(1)$. In the standard model with three generations and one Higgs doublet one has $c_s= 28/79$. We conclude that lepton number violation is necessary in order to generate a cosmological baryon asymmetry\footnote{In the case of Dirac neutrinos, which have extremely small Yukawa couplings, one can construct leptogenesis models where an asymmetry of lepton doublets is accompanied by an asymmetry of right-handed neutrinos such that the total lepton number is conserved and $\langle B-L \rangle_T = 0$ \cite{dlx00}.}. However, it can only be weak, because otherwise any baryon asymmetry would be washed out. The interplay of these conflicting conditions leads to important contraints on neutrino properties and on possible extensions of the standard model in general. \begin{figure}[t] \includegraphics[height=6cm]{ewbg} \caption{Sketch of nonlocal electroweak baryogenesis. From \cite{ber02}. } \end{figure} | 40 years after Sakharov's work on the cosmological matter-antimatter asymmetry we have several viable models of baryogenesis, the most predictive ones being electroweak baryogenesis and leptogenesis. In fact, based on our theoretical understanding of the electroweak phase diagram, electroweak baryogenesis in the standard model has already been excluded by the LEP bound on the Higgs mass. Supersymmetric electroweak baryogenesis will soon be tested at the LHC. Detailed studies of the nonequilibrium leptogenesis process have led to the preferred neutrino mass window $10^{-3}\ {\rm eV} < m_i < 0.1\ {\rm eV}$ in the simplest scenario with hierarchical heavy neutrinos. The consistency with the experimental evidence for neutrino masses has dramatically increased the popularity of the leptogenesis mechanism. It is exciting that new experiments and cosmological observations will probe the absolute neutrino mass scale in the coming years. However, more work is needed on the full quantum mechanical treatment of leptogenesis, in particular the flavour dependence. All baryogenesis mechanisms are closely related to the nature of dark matter. A discovery of the standard supergravity scenario at LHC could be consistent with electroweak baryogenesis but would rule out the simplest version of thermal leptogenesis. On the other hand, evidence for gravitino dark matter can be consistent with leptogenesis. Finally, the discovery of macroscopic dark matter like Q-balls would point towards nonperturbative dynamics of scalar fields in the early universe and therefore favour Affleck-Dine baryogenesis. | 7 | 10 | 0710.5857 |
0710 | 0710.2193_arXiv.txt | Rapid mass transfer in a binary system can drive the accreting star out of thermal equilibrium, causing it to expand. This can lead to a contact system, strong mass loss from the system and possibly merging of the two stars. In low metallicity stars the timescale for heat transport is shorter due to the lower opacity. The accreting star can therefore restore thermal equilibrium more quickly and possibly avoid contact. We investigate the effect of accretion onto main sequence stars with radiative envelopes with different metallicities. We find that a low metallicity ($Z<10^{-3}$) $4\Msun$ star can endure a 10 to 30 times higher accretion rate before it reaches a certain radius than a star at solar metallicity. This could imply that up to two times fewer systems come into contact during rapid mass transfer when we compare low metallicity. This factor is uncertain due to the unknown distribution of binary parameters and the dependence of the mass transfer timescale on metallicity. In a forthcoming paper we will present analytic fits to models of accreting stars at various metallicities intended for the use in population synthesis models. | The majority of stars are found in binaries, many of which can interact, for example by exchanging mass, resulting in an evolution very distinct from isolated stars. Although the fraction of stars in binaries might be different for earlier generations of stars, formed in metal-poor environments, they are certainly worth a systematic study. In this work we discuss the effect of accretion onto main sequence stars as function of metallicity. In a second contribution to these proceedings we discuss how the ranges for different cases of mass transfer depend on metallicity \citep{caseABC07}. Mass transfer takes place when one of the stars exceeds a certain critical radius, the Roche lobe radius. The first phase of mass transfer is usually so fast that the accreting star is driven out of thermal equilibrium and expands \citep{Benson70,Yungelson73}, potentially so much that the two stars come into contact. Contact binaries are not well understood, but they probably involve strong mass loss from the system and merging of the two stars. Instead of using full binary evolution models, we choose to reduce the large parameter space of binaries by studying the behavior of models of isolated stars under controlled conditions. We follow the approach of \citet{Kippenhahn+Meyer77} and \citet{Neo+ea77}, who studied the evolution of the radius of accreting main sequence stars. We extend their work, to study the effect of metallicity, using up-to-date input physics. Here we present the results of an exploratory study and its possible implications for binaries at low metallicity. In a forth coming paper we will present analytic fits to a finer grid of models, intended for the use in population synthesis models. Expansion due to thermal timescale mass transfer is commonly neglected or taken into account using a simple approximate criterion. Our fits will provide an improvement, which is easy to implement. \begin{figure} \includegraphics[ angle=-90, width=\columnwidth]{deminkea_poster1_fig1.ps} \caption{Radius versus mass of an initially 4 \Msun star at solar metallicity accreting with different accretion rates. \label{R_vs_m} } \end{figure} | We find that a $4\Msun$ low metallicity star ($Z<10^{-3}$) can endure a 10 to 30 times higher accretion rate before it expands to a certain radius compared to solar metallicity stars. This suggests that at low metallicity fewer binaries come into contact during rapid mass transfer. A rough estimate based on very simplified assumptions indicate that the effect could be up to a factor 2, i.e. only half as many may binaries evolve into contact at low metallicity(Z < 10{-3}) compared to solar metallicity. This factor is uncertain and probably only an upper limit as it depends on the initial distribution of binary parameters, on how the typical mass transfer rate depends on metallicity \citep[it is probably higher at low metallicity, e.g.][]{Langer_ea00} and the binary parameters, on the specific entropy of the accreted material and on the efficiency of mass transfer \citep[see for example][]{DeMink+ea07}. In a forthcoming paper we will present analytic fits to our models intended for the use in population synthesis models. \begin{theacknowledgments} We would like to thank Peter Eggleton for providing his stellar evolution code, John Eldridge for the opacity tables and Rob Izzard for interesting discussions and suggestions. \end{theacknowledgments} | 7 | 10 | 0710.2193 |
0710 | 0710.3706_arXiv.txt | The Magellanic Clouds (MCs) are interacting SBm galaxies similar to many that exist in Universe. They are the largest neighbouring satellites of the Milky Way, reflecting a typical environment of a large galaxy surrounded by satellites. They contain stars which are as old as the Universe as well as newly forming and this extended range of star formation is a highly valuable source to understand the process of formation and evolution of galaxies in general. The MCs are overall more metal poor than the Galaxy and therefore may hold information about the Universe at its early stages. They are located at a fairly well known distance, which makes it easier to measure details of their stellar component and structure. They are also fortunately located in a region of sky only lightly affected by Galactic reddening, which translates into the capability of detecting their faint stellar populations. The MCs belong to a complex system, the Magellanic System, which has in total four distinct components: the Large Magellanic Cloud (LMC), the Small Magellanic Cloud (SMC), the Bridge connecting the two Clouds and the Stream attached to the SMC. The latter two are predominantly formed of gas and are of tidal origin. | The Magellanic System has yet many challenging aspects that new surveys, with the increased quality of the coming data and new theoretical models and their ability to explain detail observations, aim to resolve in the next decade. Prior to new facilities like GAIA, JWST and ALMA we need to exploit data from VISTA and similarly powerful telescopes at other wavelengths. Surveys like VMC will provide unique and high quality data for science and training of young astronomers. | 7 | 10 | 0710.3706 |
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0710 | 0710.3530_arXiv.txt | We present the first measurements of the angular correlation function of galaxies selected in the far (\fuvcenter) and near (\nuvcenter) Ultraviolet from the \galex survey fields overlapping SDSS DR5 in low galactic extinction regions. The area used covers $120$ sqdeg (\galex - MIS) down to magnitude AB $=22$, yielding a total of 100,000 galaxies. The mean correlation length is $\sim3.7 \pm 0.6$~Mpc and no significant trend is seen for this value as a function of the limiting apparent magnitude or between the \galex bands. This estimate is close to that found from samples of blue galaxies in the local universe selected in the visible, and similar to that derived at $z\simeq3$ for LBGs with similar rest frame selection criteria. This result supports models that predict anti-biasing of star forming galaxies at low redshift, and brings an additional clue to the downsizing of star formation at $z<1$. | In the current paradigm of structure formation, the bulk of the most massive systems form in a cold dark matter-dominated universe by the merging of less massive units formed earlier. In parallel to this hierarchical evolution, recent observations point to the so-called ``downsizing'', namely the fact that in galaxies having high baryonic masses the bulk of stars formed at high redshift ($z \gtrsim 1$), while in galaxies having low baryonic masses the bulk of stars formed at lower redshift \citep[][and also \citet{DeLucia_2006} and \citet{Neistein_2006} for results from simulations]{Cowie_1996, Heavens_2004, Bundy_2006, Jimenez_2005}. The star formation efficiency shows a strong decline at $0<z<1$, as measured by the evolution of the star formation rate density \citep{Hopkins_2006, Lilly_1996, Schiminovich_2005, Sullivan_2000, Wilson_2002}. These epochs also see the bulk of the build-up of the bimodality in galaxy properties of the local universe, which is apparent in their color distribution \cite{Baldry_2004}, morphologies \citep{Kauffmann_2004}, spectral class \citep{Madgwick_2002} and spatial distribution \citep{Budavari_2003}. Understanding the full picture is complex as this evolution is the result of the interplay of several physical processes \citep{Faber_2005} and combine the effects of initial galaxy formation conditions (``nature'') with galaxy evolution events (``nurture'') \citep{Kauffmann_2004}. In this context, tracers that measure over cosmic time galaxy populations selected with homogeneous physical criteria are of primary interest. They help compare observations to simulation predictions over a large range of redshifts with reduced uncertainties, and allow a study of the redshift evolution of galaxy properties derived from different surveys. The ultraviolet (UV) range of the spectrum meets these conditions: UV luminosities provide a good measure of recent star formation within galaxies \citep{Kennicutt_1998}, modulo attenuation by dust, and has been widely used at high redshifts to study the properties of the Lyman Break Galaxies (LBGs) \citep{Giavalisco_2001, Shapley_2003,Steidel_1995}. As large amounts of data are now becoming available at lower redshifts as part of the GALEX surveys \citep{Martin_2005}, the restframe UV spectral domain is presently well sampled over the full $0<z<6$ redshift range. Furthermore, comparison of results from high and low $z$ UV-selected samples is eased by the fact that the UV luminosity density fractions\footnote{We define the UV luminosity density fraction of a given sample as the ratio of the UV luminosity density encompassed by the sample over the total UV luminosity density at the same redshift.} probed at high and low $z$ are similar \citep[][hereafter Paper II]{Heinis_2007}, due to the strong luminosity evolution of the UV luminosity function \citep{Arnouts_2005}. Noticeably, during the epochs probed by GALEX the properties of active star forming galaxies show a very fast evolution. The wealth of UV-selected data now available at low redshifts enables statistical studies in the context of the downsizing of star formation, and in particular searches for links between the star formation properties and galaxy environment in terms of galaxy or dark matter density. Here we focus on the evolution with redshift of the link of star formation with Dark Matter and particularly the evolution of the class of Dark Matter halos hosting actively star forming galaxies since $z\sim1$. This can be achieved by the study of the clustering of galaxies: at high redshift, LBGs studies show that UV selected galaxies inhabit high galaxy density regions \citep{Steidel_1998}, and are strongly biased with respect to the underlying Dark Matter, with more actively star forming galaxies being more biased \citep[][see \citet{Giavalisco_2002} for a review on the properties of LBGs]{Adelberger_2005, Giavalisco_2001, Foucaud_2003}. We propose to extend such studies to low redshifts using similar selection criteria. This paper is the first in a series and presents the methods and first results of angular clustering measurements from GALEX data. The following section presents the datasets and the derivation of the redshift distributions. Section \ref{sec_acf_methods} presents two methods to derive the angular correlation function from a set of fields, and a discussion about the behavior of these methods with respect to photometry inhomogeneity. In section \ref{sec_acf} we present our results on the angular correlation functions and correlation lengths. To provide the crucial link to Dark Matter halos, we use the analytical \citet{Mo_2002} formalism that we present in sec. \ref{sec_mo_white}. We end with a short discussion in sec. \ref{sec_discussion}. Throughout the paper a ${\Lambda}CDM$ cosmology is assumed with matter density $\Omega_m = 0.3$, vacuum energy density $\Omega_\Lambda=0.7$, and a Hubble parameter $h=0.7$ where $H_0 = 70 $km s$^{-1}$ Mpc$^{-1}$. All correlation length values taken from the literature have been converted accordingly using equation (4) in \citet{Magliocchetti_2000}. | 7 | 10 | 0710.3530 |
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0710 | 0710.0207.txt | % Aims -- Methods --Results -- Conclusions The first steps of planet formation are marked by the growth and crystallization of sub--micrometer--sized dust grains accompanied by dust settling toward the disk midplane. In this paper we explore whether the first steps of planet formation are affected by the presence of medium--separation stellar companions. We selected two large samples of disks around single and binary T~Tauri stars in Taurus that are thought to have only a modest age spread of a few Myr. The companions of our binary sample are at projected separations between $\sim$10 and 450\,AU with masses down to about 0.1\,M$_{\sun}$. We used the strength and shape of the 10\,\micron{} silicate emission feature as a proxy for grain growth and for crystallization respectively. The degree of dust settling was evaluated from the ratio of fluxes at two different mid--infrared wavelengths. We find no statistically significant difference between the distribution of 10\,\micron{} silicate emission features from single and binary systems. In addition, the distribution of disk flaring is indistinguishable between the single and binary system samples. These results show that the first steps of planet formation are not affected by the presence of a companion at tens of AU. | Two--third of the G stars in the solar neighborhood are members of multiple--star systems (e.g. \citealt{dm91}). These binaries and multiple systems are often found to harbor giant planets (e.g. \citealt{bondes07}). Similarly, young low--mass pre--main sequence stars are very frequently members of multiple systems, mostly binaries \citep{mathieu00,duchene07}. This suggests that planet formation around single stars such as our Sun may be atypical and urges us to understand the effects of stellar companions on planet formation. We tackle this question from an observational point of view. Because numerical simulations of grain agglomeration suggest short timescales for the formation of planetesimals (only a few 10$^4$\,yr, e.g. \citealt{beckwith00}), it is crucial to know how stellar companion(s) affect the dust processing in the first few million years. Grain growth and the settling of dust grains towards the disk midplane are thought to represent the first steps in the planet--formation process (e.g. \citealt{ls93} for a review). The study of disks around intermediate--mass stars also indicates a link between grain growth and crystallinity. High crystallinity was found only when grains larger than the dominant sub--micron interstellar grains were present \citep{van05}. In the context of these findings and dust evolution models (e.g. \citealt{dd04}), older disks are expected to have more processed dust (larger grains and crystals) than younger disks. In addition, their disk structure should be flatter because of the gradual settling of large dust grains towards the disk midplane. However, recent observations show that the degree of dust processing can be very different even for coeval disks around stars of similar spectral type in the same star--forming region (e.g. \citealt{prz03,apai05}). This demonstrates that dust evolution is not uniquely controlled by stellar age and luminosity but at least one additional parameter is present. There are two studies suggesting that stellar multiplicity could play a major role in the initial dust processing. \citet{meeus03} found that among three coeval T~Tauri disks in the Chamaeleon~I star--forming region the closest binary system (projected separation of $\sim$120\,AU) sports the strongest contribution from large ($\sim$\,2\,\micron) grains and has the highest crystalline mass fraction. Similarly, \citet{sterzik04} pointed out that the disk of a young brown dwarf with a companion at $\ga$\,30\,AU shows more processed dust than the disk around a single brown dwarf. Although the small samples inhibit any firm conclusions, these results suggest that companions might trigger rapid dust evolution. Intuitively this may happen in different ways. A companion could speed up dust evolution by dynamically stirring the circumstellar dust grains and leading to an enhanced collision and grain growth rate (e.g. \citealt{dubrulle95}). In addition, the dynamical stirring may lead to an increased mixing that could also expose larger amounts of dust to temperatures high enough ($\ge$\,800\,K) to be crystallized. In this paper we compare two carefully constructed samples of disks around single and binary\footnote{a few of the systems in our study are triple or quadruple systems. For simplicity, we refer to the whole sample as binaries)} stars with a narrow age spread to test the hypothesis that binary systems have disks with more processed dust and flatter structures. In Sect.~\ref{S:sample} we describe our samples. The data reduction of the {\it Spitzer} spectra and of the 24\,\micron{} MIPS photometry is presented in Sect.~\ref{S:redu}. We summarize our results in Sect.~\ref{S:results} and discuss in Sect.~\ref{S:discussion} their implications on planet formation in single and binary systems. %{\bf COMMENT: note that stellar multiplicity does not preclude planet formation} %{\bf COMMENT: note no significant difference in NIR excess between binaries and singles PPVI, nor in the frequencies of IRAs detections. But differences in mm fluxes (smaller masses)--similar temperatures surface densities, similar planet formation} %{\bf COMMENT: from PPV Duchene et al. among filed-stars the frequency of high-order multiple systems is 4 times lower than that of binaries.Shall we shift to binaries rather than multiple systems?} %{\bf COMMENT: there is a statistically significant difference in the mass distribution of planets around binariee and single stars: massive planets in short orbits (0.05AU) are found mainly around tight binaries.} | \label{S:discussion} Our study shows a large diversity in the 10\,\micron{} silicate emission features and SED slopes of T~Tauri disks. We found that neither the dust processing nor the disk flaring correlates with the multiplicity of the sources. These results are particularly interesting for two aspects that will be discussed in the following. \subsection{Medium--separation binaries and planet formation} A stellar companion induces tidal forces in a disk that become particularly strong at resonance points. Resonant interactions result in the excitation of density waves that can truncate a disk and act to modify the binary eccentricity (see e.g. \citealt{lubow00} for a review). Theoretical calculations of binary--disk interactions predict that circumstellar disks will be truncated at 0.2--0.5 times the binary semimajor axis $a$, with the exact values depending on eccentricity, mass ratio, and disk viscosity \citep{art94}. These theoretical expectations are supported by millimeter observations of binaries tracing the optically thin dust emission and thus the total disk mass in the system. There is evidence for a diminished millimeter flux (hence disk mass) among the 1--100\,AU binaries in comparison to wider binaries or single stars (e.g. \citealt{mathieu00} for a review). This result is qualitatively consistent with the circumstellar disks of medium--separation binaries being tidally truncated at 0.2--0.5$a$. Two--thirds of our sources have stellar companions between 0.1\arcsec--1\arcsec{}, with a mean projected separation of 0.4\arcsec{} or 56\,AU at the distance of Taurus. Therefore, the typical truncation radius for disks in our sample is $>$11--28\,AU, well outside the location of Jupiter and Saturn in our Solar System. Even in other systems these outer radii are found to be devoid of giant planets (e.g. \citealt{kasper07}). This fact suggests that the formation of terrestrial and giant planets may proceed undisturbed in disks around medium--separation binaries even if these disks are constrained in size. %Although disks in binaries are constrained in size, the formation of terrestrial and giant planets could then proceed undisturbed if the circumstellar material is similar in surface density and temperature to that in disks around single stars. These parameters are difficult to constrain from observations, but there are indications of a similar pathway for planet formation in disks around single and medium--separation binaries. Early investigations of young TTSs found no significant difference in the frequency of near-- and mid--infrared excess emission between single and binary star systems (e.g. \citealt{simonprato95,jensen96}). With the 60\micron{} IRAS flux probing dust $\la$10\,AU from the central star, these measurements demonstrate that binary systems as often have disks as single stars do. Recently \citet{monin07} analyzed the separation distributions of binaries with and without disks and found no statistical difference. Since most of their binaries have projected separations $>$20\,AU, their result shows that medium--and wide-- separation binaries do not have a significant effect on the circumstellar disk lifetime. Our work indicates that these disks also evolve in a similar way. The extent of dust processing in the disk surface layer and the degree of dust settling in binary disk systems do not statistically differ from those in disks around single stars. This suggests that the first few Myr of disk evolution in the terrestrial (and maybe out to the giant) planet--forming region are not affected by medium--separation stellar companions. Whether the disk evolution proceeds undisturbed for tens of millions of years until planets are fully formed cannot yet be assessed observationally. \citet{bouwman06} estimate a mean disk dispersion timescale of $\sim$5\,Myr for close ($\le$4AU) binaries in contrast to a timescale of $\approx$9\,Myr for single star systems. They argue that the time available to form planets in close binary systems is considerably shorter than that in disks around single stars, which may inhibit planet formation. The only two medium--separation binaries in their sample hint for a disk dispersal timescale comparable to that of single stars suggesting a similar disk evolution for single and medium--separation binary systems over the first $\sim$\,10\,Myr. Exoplanet surveys offer us a glimpse into the frequency and properties of giant planets in multiple star systems. Recently \citet{eu07} reported 42 planets orbiting binary and multiple stars (see, their Table 1). \citet{bondes07} analyze a subsample of radial velocity planet host stars with uniform planet detectability and demonstrate that the overall frequency of giant planets in binaries is not statistically different from that of planets in single stars. However, they find indications for a lower frequency of radial velocity planets in the subgroup of close-- and medium--separation binaries ($<\,50-100$\,AU). In a complementary study, \citet{desbar07} find that the mass distribution of planets in binaries with separations $<\,300-500$\,AU is statistically different from that around wider binaries and single stars: Massive planets in short--period orbits are found predominantly around close-- and medium--separation binaries. Taken together, the results from the frequency and properties of exoplanets suggest that a stellar companion with separation less than a few hundred AU affects giant planet formation and/or the subsequent migration. Numerical simulations seem to support this notion. \citet{kley00} shows that a fairly eccentric ($e_{\rm bin}=0.5$) stellar companion at 50--100\,AU enhances the growth rate of a Jupiter mass planet embedded in a circumstellar disk and makes its inward migration more rapid. Recently, \citet{kn07} confirm these trends by following the evolution of a 30\,M$_\earth$ protoplanet in a disk truncated by a stellar companion at 18.5\,AU and $e_{\rm bin}=0.36$, like the $\gamma$~Cep binary system. Our study shows that the early evolution of protoplanetary disks surrounding binary stars is similar to that in single stars indicating that that the differences in the exoplanet properties arise in the later stages of their formation and/or migration. Whether terrestrial planet formation is also affected by medium--separation binaries cannot be yet addressed observationally. Our study shows that the initial dust processing is not impacted by the presence of a stellar companion. Based on the fact that the build--up of planetesimals as large as the $\sim$500--km Vesta has occurred in the first 3.8$\pm$1.3~Myr of the Solar nebula \citep{kleine02}, it is reasonable to speculate on the basis of our study that the formation of planetesimals in binary and single systems proceed along, if not on identical avenues. Another indication supporting this suggestion comes from the finding of a similar incidence of debris disks in Gyr--old single and binary stars \citep{trilling06}. If the debris dust is produced by colliding asteroids, then the similar rate of debris dust in binaries implies that planetesimal formation is not inhibited by the presence of stellar companions. Recent simulations of the later stages of terrestrial planet formation show that rocky planets can form in a wide variety of binary systems \citep{quinta07}. The binary periastron is the most important parameter in limiting the number of forming planets and their range of orbits. \citet{quinta07} show that binaries with periastron $\ga$10\,AU, comprising most of the medium--separation binaries investigated in this paper, can form terrestrial planets over the entire range of orbits allowed for single stars. As a result more than 50\% of the binary systems in the Milky Way \citep{dm91} are wide enough to allow the formation of Earth--like planets. \subsection{The diversity in silicate features and SEDs} Although small sample statistics suggested a correlation between stellar multiplicity and initial dust processing \citep{meeus03,sterzik04,sic07}, our study demonstrates that medium--separation stellar companions do not appreciably affect the growth and crystallization of dust in circumstellar disks. Given the criteria applied to select our samples, we can also exclude that age, spectral type, and stellar environment can account for the large variety of observed silicate emission features and SED slopes in our study. There may be several other factors contributing to this diversity that will be fully explored in an upcoming contribution. In the following we briefly mention two of them: Turbulence in circumstellar disks not only drives the accretion of gas onto the central star but also replenishes the disk atmosphere with more grains that can be larger in size. If the grains inferred from the 10\,\micron{} silicate emission feature reflect the level of disk turbulence, the strength of the features should depend on the stellar accretion rates. \citet{sic07} note that stars with strong features tend to have large accretion rates in their sample of several Myr old intermediate-- and low--mass stars. This trend may be the result of turbulence determining the grain population in the disk atmosphere. Alternatively, the trend could be due to the more massive stars (that have typically larger accretion rates) in their sample heating larger disk area and thus producing stronger silicate emission features (see, e.g. \citealt{kessler07}). The tentative correlation seen in the sample of \citet{sic07} needs to be confirmed using a larger and more homogeneous sample of stars with well--determined accretion rates. %Stellar high--energy photons (UV and X--rays) may also affect the distribution of grain sizes and the composition of solids in circumstellar disks. Low--mass TTs are strong emitters of relative hard KeV X-rays (Wolk et al. 2005) that could evaporate the grain mantle or the complete grain. Voit (1992) calculated that silicate ?? grains with a radius $<$10\AA evpaorates completed when subject to an X-ray energu of zz, which could be reached in the disk atmospheres for grains closer than yy from the central star. This could produce a grain size distribution biased toward larger grains. I DON'T UDERSTAND THE DRAINE A EQUANTION IT GIVES UNREASONABLE RESULTS!!! It says that 50um gains can survive (not evaporate as close as 2 stellar radii.) %Low-mass TTSs also show for about 25\% of their time intense X--ray flares that can carry energies as high as a few MeV (Wolk et al. 2005). Such flares could melt and crystallize grains and may be responsible for the flash melted chondrules in meteorites (e.g. Shu et al. 2001). To asses the effect of stellar high--energy photons on small grains it would be interesting to correlate stellar UV/X-ray variability with changes in the strength and shape of the 10\,\micron{} silicate emission features. Different initial conditions for the collapsing cores may also leave their imprints on the formation and evolution of circumstellar disks. This possibility has been explored by \citet{dullemond06} to explain crystallization of dust grains in the early stages of disk evolution. In their model the level of crystallinity depends crucially on the rotation rate of the collapsing cloud core because this determines the radius at which the infalling matter reaches the disk: rapidly rotating clouds would evolve into disks with low crystallinity, while slowly rotating clouds into disks with high crystallinity. In this paper we explored the effect of a stellar companion on the initial growth and settling of dust grains in circumstellar disks. We constructed two large samples of disks around single and binary TTSs with a narrow age spread and a spectral type distribution for the single stars identical to that of the primary stars in the binary sample. We used the strength of the 10\,\micron{} silicate emission feature derived from IRS/{\it Spitzer} spectra as a proxy for grain growth and the SED slope of circumstellar disks as a proxy for dust settling. Our results can be summarized as follows: \\ \smallskip -- there is no statistically significant difference between the distribution of 10\,\micron{} silicate emission features from single and binary systems. \\ -- the distribution of disk flaring is indistinguishable between the single and binary system samples. \\ \smallskip These results show that stellar companions at projected separations of $\ga$\,10\,AU do not appreciably affect the degree of crystallinity nor the degree of grain growth. Based on the combination of these and other results we argue that the formation of planetesimals and possibly terrestrial planets is not inhibited in a circumstellar disk perturbed by a medium--separation stellar companion. %%%%%%%%%% Single star spectra \begin{figure*} \includegraphics[angle=90,scale=0.7]{f8.ps} \caption{Infrared spectra for the sample of single stars.\label{s1}} \end{figure*} \begin{figure*} \includegraphics[angle=90,scale=0.7]{f9.ps} \caption{Infrared spectra for the sample of single stars.\label{s2}} \end{figure*} \begin{figure*} \includegraphics[angle=90,scale=0.7]{f10.ps} \caption{Infrared spectra for the sample of single stars.\label{s3}} \end{figure*} \begin{figure*} \includegraphics[angle=90,scale=0.7]{f11.ps} \caption{Infrared spectra for the sample of single stars.\label{s4}} \end{figure*} %%%%%%%%%% Binary star spectra \begin{figure*} \includegraphics[angle=90,scale=0.7]{f12.ps} \caption{Infrared spectra for the sample of binary stars.\label{b1}} \end{figure*} \begin{figure*} \includegraphics[angle=90,scale=0.7]{f13.ps} \caption{Infrared spectra for the sample of binary stars.\label{b2}} \end{figure*} \begin{figure*} \includegraphics[angle=90,scale=0.7]{f14.ps} \caption{Infrared spectra for the sample of binary stars.\label{b3}} \end{figure*} \begin{figure*} \includegraphics[angle=90,scale=0.7]{f15.ps} \caption{Infrared spectra for the sample of binary stars.\label{b4}} \end{figure*} | 7 | 10 | 0710.0207 |
0710 | 0710.1473_arXiv.txt | In this contribution we study integrated properties of dynamically segregated star clusters. The observed core radii of segregated clusters can be 50\% smaller than the ``true'' core radius. In addition, the measured radius in the red filters is smaller than those measured in blue filters. However, these difference are small ($\lesssim10\%$), making it observationally challenging to detect mass segregation in extra-galactic clusters based on such a comparison. Our results follow naturally from the fact that in nearly all filters most of the light comes from the most massive stars. Therefore, the observed surface brightness profile is dominated by stars of similar mass, which are centrally concentrated and have a similar spatial distribution. | 7 | 10 | 0710.1473 |
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0710 | 0710.0063_arXiv.txt | An investigation of the dynamo instability close to the threshold produced by an ABC forced flow is presented. We focus on the on-off intermittency behavior of the dynamo and the counter-effect of the Lorentz force in the non-linear stage of the dynamo. The Lorentz force drastically alters the statistics of the turbulent fluctuations of the flow and reduces their amplitude. As a result much longer burst (on-phases) are observed than what is expected based on the amplitude of the fluctuations in the kinematic regime of the dynamo. For large Reynolds numbers, the duration time of the ``On'' phase follows a power law distribution, while for smaller Reynolds numbers the Lorentz force completely kills the noise and the system transits from a chaotic state into a ``laminar'' time periodic flow. The behavior of the On-Off intermittency as the Reynolds number is increased is also examined. The connections with dynamo experiments and theoretical modeling are discussed. \pacs{47.65.-d,47.20.Ky,47.27.Sd,52.65.Kj} | Dynamo action, the self amplification of magnetic field due to the stretching of magnetic field lines by a flow, is considered to be the main mechanism for the generation of magnetic fields in the universe \cite{mhdbooks}. To that respect many experimental groups have successfully attempted to reproduce dynamos in liquid sodium laboratory experiments \cite{gailitis2000,gailitis2001,gailitis2004,muller2000,stieglitz2001,monchaux2007,berhanu2007}. The induction experiments \cite{odier1998,peffley2000a,peffley2000b,frick2002,bourgoin2002,nornberg2006a,nornberg2006b,stepanov2006,volk2006,bourgoin2006} studying the response of an applied magnetic field inside a turbulent metal liquid represent also a challenging science. With or without dynamo instability the flow of a conducting fluid forms complex system, with a large degree of freedoms and a wide branch of non linear behaviors. In this work we focus on one special behavior: the On-Off intermittency or blowout bifurcation \cite{pomeau1980,platt1993}. On-off intermittency is present in chaotic dynamical systems for which there is an unstable invariant manifold in the phase space such that the unstable solutions have a growth rate that varies strongly in time taking both positive and negative values. If the averaged growth rate is sufficiently smaller than the fluctuations of the instantaneous growth rate, then the solution can exhibit on-off intermittency where bursts of the amplitude of the distance from the invariant manifold are observed (when the growth rate is positive) followed by a decrease of the amplitude (when the growth rate is negative). (See \cite{sweet2001a,sweet2001b} for a more precise definition). On-Off intermittency has been observed in different physical experiments including electronic devices, electrohydrodynamic convection in nematics, gas discharge plasmas, and spin-wave instabilities \cite{phys-onoff}. In the MHD context, near the dynamo instability onset, the On-Off intermittency has been investigated by modeling of the Bullard dynamo \cite{leprovost2006}. Using direct numerical simulation \cite{sweet2001a,sweet2001b} were able to observe On-Off intermittency solving the full MHD equations for the ABC dynamo, (here we present an extended work of this particular case). On-Off intermittency has also been found recently for a Taylor-Green flow \cite{ponty2007b}. Finally, recent liquid metal experimental results (VKS) \cite{pinton2007} show some intermittent behavior, with features reminiscent of on-off self-generation that motivated our study. For the MHD system we are investigating the evolution of the magnetic energy $E_b=\frac{1}{2}\int {\bf b}^2 dx^3$ is given by $\partial_t { E_b} = \int {\bf b( \cdot b \nabla) u - \eta (\nabla b)^2} dx^3$. If the velocity field has a chaotic behavior in time the right hand side of the equation above can take positive or negative values and can be modeled as multiplicative noise. A simple and proved very useful way to model the behavior of the magnetic field during the on-off intermittency is using a stochastic differential equation (SDE-model) \cite{pomeau1980,platt1993,fujisaka,Yu1990,ott1004,Platt1994,Heagy1994,Venka1995,Venka1996,aumaitre2006,aumaitre2005}: \begin{equation} \partial_t E_b = (a+\xi) E_b - NL(E_b) \label{SDE} \end{equation} where $ E_b$ is the magnetic energy, $a$ is the long time averaged growth rate, $\xi$ models the noise term typically assumed to be white (see however \cite{aumaitre2006,aumaitre2005}) and of amplitude $D$ such that $\langle \xi(t)\xi(t') \rangle = 2D\delta(t-t')$. $NL$ is a non-linear term that guaranties the saturation of the magnetic energy to finite values typically taken to be $NL(X)=X^3$ for investigations of supercritical bifurcations or $NL(X)=X^5-X^3$ for investigations of subcritical bifurcations. Alternative, an upper no-flux boundary is imposed at $E_b=1$. In all these cases (independent of the non-linear saturation mechanism) the above SDE leads to the stationery distribution function that for $0<a<D$ has a singular behavior at $E_b=0$: $P(E_b)\sim E_b^{a/D-1}$ indicating that the systems spends a lot of time in the neighborhood of $E_b=0$. This is singularity is the signature of On-Off intermittency. Among other predictions of the SDE model here we note that the distribution of the duration time of the ``off" phases follows a power law behavior $PDF(\Delta T_{off})\sim \Delta T_{off}^{-1.5}$, all moments of the magnetic energy follow a linear scaling with $a$, $\langle E_b^m \rangle \sim a$, and for $a=0$ the set of the burst has a fractal dimension $d=1/2$ \cite{Platt1994,Heagy1994,Venka1995,Venka1996}. In this dynamical system eq.(\ref{SDE}) however the noise amplitude or the noise proprieties do not depend on the amplitude of the magnetic energy. However, in the MHD system, when the non-linear regime is reached, the Lorentz force has a clear effect on the the flow such as the decrease of the small scale fluctuation, and the decrease of the local Lyapunov exponent \cite{cattaneo1996,zienicke1998}. Some cases, the flow is altered so strongly that the MHD dynamo system jumps into an other attractor, that cannot not sustain any more the dynamo instability \cite{brummell}. Although the exact mechanism of the saturation of the MHD dynamo is still an open question that might not have a universal answer, it is clear that both the large scales and the turbulent fluctuations are altered in the non-linear regime and need to be taken into account in a model. \begin{figure} \includegraphics[width=8cm]{fig1.eps} \caption{A typical example of a burst. The top panel shows the evolution of the kinetic energy (top line) and magnetic energy (bottom line). The bottom panel shows the evolution of the magnetic energy in a log-linear plot. During the on phase of the dynamo the amplitude of the noise of the kinetic energy fluctuations is significantly reduced. The runs were for the parameters $Gr=39.06$ and $G_M=50.40$. \label{fig1}} \end{figure} Figure \ref{fig1} demonstrates this point, by showing the evolution of the kinetic and magnetic energy as the dynamo goes through On- and Off- phases. During the On phases although the magnetic field energy is an order of magnitude smaller than the kinetic energy both the mean value and the amplitude of the observed fluctuations of the kinetic energy are significantly reduced. As a result the On-phases last a lot longer than what the SDE-model would predict. With our numerical simulations, we aim to describe which of the On-Off intermittency proprieties are affected through the Lorentz force feed-back. This paper is structured as follows. In the next section \ref{Nmethod} we discuss the numerical method used. In section \ref{Onset} we present the table of our numerical runs and discuss the dynamo onset. Results for small Reynolds numbers investigating the transition from a laminar dynamo to on-off intermittency are presented in \ref{Route}, and the results on fully developed on-off intermittency behavior are given in section \ref{Onoff}. Conclusions, and implications on modeling and on the laboratory experiments are given in the last section \ref{Cons}. | \label{Cons} In this work we have examined how the on-off intermittency behavior of a near criticality dynamo is changed as the kinematic Reynolds is varied, and what is the effect of the Lorentz force in the non-linear stage of the dynamo. The predictions of \cite{Platt1994,Heagy1994,Venka1995,Venka1996}, linear scaling of the averaged magnetic energy with the deviation of the control parameter from its critical value, fractal dimensions of the bursts, distribution of the ``off" time intervals, and singular behavior of the pdf of the magnetic energy that were tested numerically in \cite{sweet2001a,sweet2001b} were verified for a larger range of Kinematic Grashof numbers when On-Off intermittency was present. Note however that all these predictions are based on the statistics of the flow in the kinematic stage of the dynamo. However it was found that the Lorentz force can drastically alter the On-Off behavior of the dynamo in the non-linear stage by quenching the noise. For small Grashof numbers the Lorentz force can trap the original chaotic system in the linear regime in to a time periodic state resulting to no On-Off intermittency. At larger Grashof numbers $Gr>20$ On-Off intermittency was observed but with long durations of the ``on" phases that have a power law distribution. These long ``on" phases result in a pdf that peaks at finite values of $E_b$. This peak can be attributed to the presence of a subcritical instability or to the quenching of the hydrodynamic ``noise" at the nonlinear stage or possibly a combination of the two. In principle the SDE model (eq.\ref{SDE}) can be modified to include these two effects: a non-linear term that allows for a subcritical bifurcation and a $E_b$ dependent amplitude of the noise. There many possibilities to model the quenching of the noise, however the nonlinear behavior might not have a universal behavior and we do not attempt to suggest a specific model. The relative range of the On-Off intermittency was found to decrease as the Reynolds number was increased possibly reaching an asymptotic regime. However the limited number of Reynolds numbers examined did not allow us to have a definite prediction for this asymptotic regime. This question is of particular interest to the dynamo experiments \cite{gailitis2000,gailitis2001,gailitis2004,muller2000,stieglitz2001,monchaux2007,berhanu2007} that until very recently \cite{pinton2007} have not detected On-Off intermittency . There are many reasons that could explain the absence of detectable On-Off intermittency in the experimental setups, like the strong constrains imposed on the flow \cite{gailitis2004,muller2000} that do not allow the development of large scale fluctuations or the Earths magnetic field that imposes a lower threshold for the amplitude of the magnetic energy. Numerical investigations at higher resolution and a larger variety of flows or forcing would be useful at this point to obtain a better understanding. | 7 | 10 | 0710.0063 |
0710 | 0710.0580_arXiv.txt | {Using lattice techniques we investigate the generation of long range cosmological magnetic fields during a cold electroweak transition. We will show how magnetic fields arise, during bubble collisions, in the form of magnetic strings. We conjecture that these magnetic strings originate from the alignment of magnetic dipoles associated with EW sphaleron-like configurations. We also discuss the early thermalisation of photons and the turbulent behaviour of the scalar fields after tachyonic preheating.} \FullConference{The XXV International Symposium on Lattice Field Theory\\ July 30 - August 4 2007\\ Regensburg, Germany} \begin{document} | There have been many theoretical attempts to explain the origin of large scale cosmological magnetic fields (LSMF)~\cite{giova}. The main difficulty resides in understanding their correlation scale which ranges from the size of galaxies to clusters and super-clusters with an amplitude of the order of micro-gauss, pointing to a primordial origin. Following the work initiated in~\cite{lat05}, we address this issue in the context of a cold electroweak transition taking place after a period of hybrid inflation. The EW transition has been in the heart of many proposals to address magnetogenesis, linking it in many cases with the generation of the baryon asymmetry. The results presented here resemble the mechanism proposed by Vachaspati connecting the appearance of magnetic fields to that of sphalerons and Z-strings~\cite{vachas}. The model we have considered is a hybrid inflation model with the bosonic field content of the Standard Model coupled, via the Higgs field, to a singlet inflaton: \begin{eqnarray} {\cal L} = - \frac{1}{4}G^a_{\mu\nu}G^{\mu\nu}_a - \frac{1}{4}F^{Y}_{\mu\nu}F_{Y}^{\mu\nu}+ {\rm Tr}\Big[(D_\mu\Phi)^\dag D^\mu\Phi\Big ]+ \frac{1}{2}(\partial_\mu\chi)^2 - V(\Phi,\chi)\\ {\rm V} (\Phi,\chi) = {\rm V}_0 + \frac{1}{2}(g^2\chi^2-m^2)\,|\Phi|^2 + \frac{\lambda}{4} |\Phi|^4 + \frac{1}{2} \mu^2 \chi^2 \, \end{eqnarray} The couplings are fixed to the standard model values for several ratios of the Higgs to W masses. For concreteness, we have fixed the inflaton to Higgs coupling by the relation: $g^2= 2\lambda $. To solve the time evolution of the system, starting at the end of inflation, we have performed a numerical evolution based on a suitable classical approximation (more details can be found in Refs.~\cite{marga2,lat05,articulo}). In this work we discuss the results for $\mh = 4.65\ \mw$. Results for more realistic values will be presented in~\cite{articulo}. | We have analysed numerically the proposal that long range magnetic fields could be generated during a cold electroweak transition after a period of low scale hybrid inflation. The generation mechanism is mainly based on two facts: \begin{itemize} \item At the SSB stage bubble-like structures, associated to local maxima in the Higgs-field norm, appear. Points outside the bubble front, remaining close to the false vacuum, form string like structures. \item Bubble collisions give rise to sphaleron-like configurations attached to the location of zeroes of the Higgs field. For $\theta_W \ne 0$ these sphalerons behave as magnetic dipoles. \end{itemize} These two ingredients together lead to an allignment of the sphalerons' dipoles forming magnetic string networks as we have observed in our simulations. Some results concerning the coherence and intensity of the generated magnetic fields have been presented. A further analysis of the time evolution of the magnetic field, as well as the study of the dependence of the $\mh$ to $\mw$ ratio will be presented in~\cite{articulo}. We have also discussed some features of the late time behaviour of the system, among them, thermalisation of photon radiation and turbulence in the scalar fields. | 7 | 10 | 0710.0580 |
0710 | 0710.2250_arXiv.txt | Networks of reactions on dust grain surfaces play a crucial role in the chemistry of interstellar clouds, leading to the formation of molecular hydrogen in diffuse clouds as well as various organic molecules in dense molecular clouds. Due to the sub-micron size of the grains and the low flux, the population of reactive species per grain may be very small and strongly fluctuating. Under these conditions rate equations fail and the simulation of surface-reaction networks requires stochastic methods such as the master equation. However, the master equation becomes infeasible for complex networks because the number of equations proliferates exponentially. Here we introduce a method based on moment equations for the simulation of reaction networks on small grains. The number of equations is reduced to just one equation per reactive specie and one equation per reaction. Nevertheless, the method provides accurate results, which are in excellent agreement with the master equation. The method is demonstrated for the methanol network which has been recently shown to be of crucial importance. | Chemical networks in interstellar clouds consist of gas-phase and grain-surface reactions \citep{Hartquist1995,Tielens2005}. Reactions that take place on dust grains include the formation of molecular hydrogen \citep{Gould1963,Hollenbach1971b} as well as reaction networks producing ice mantles and various organic molecules. Unlike gas phase reactions in cold clouds that mainly produce unsaturated molecules, surface processes are dominated by hydrogen-addition reactions that result in saturated, hydrogen-rich molecules, such as H$_2$CO, CH$_3$OH, NH$_3$ and CH$_4$. In particular, recent experiments show that methanol cannot be efficiently produced by gas phase reactions \citep{Geppert2006}. On the other hand, there are indications that it can be efficiently produced on ice-coated grains \citep{Watanabe2005}. Therefore, the ability to perform simulations of the production of methanol and other complex molecules on grains is of great importance \citep{Garrod2006}. Unlike gas-phase reactions, simulated using rate equation models \citep{Pickles1977,Hasegawa1992}, grain-surface reactions require stochastic methods such as the master equation \citep{Biham2001,Green2001}, or Monte Carlo (MC) simulations \citep{Charnley2001}. This is due to the fact that under interstellar conditions, of extremely low gas density and sub-micron grain sizes, surface reaction rates are dominated by fluctuations which cannot be accounted for by rate equations \citep{Tielens1982,Charnley1997,Caselli1998,Shalabiea1998}. A significant advantage of the master equation over MC simulations is that it consists of differential equations, which can be easily coupled to the rate equations of gas-phase chemistry. Furthermore, unlike MC simulations that require the accumulation of statistical information over long times, the master equation provides the probability distribution from which the reaction rates can be obtained directly. However, the number of equations increases exponentially with the number of reactive species, making the simulation of complex networks infeasible \citep{Stantcheva2002,Stantcheva2003}. The recently proposed multi-plane method dramatically reduces the number of equations, by breaking the network into a set of fully connected sub-networks \citep{Lipshtat2004}, enabling the simulation of more complex networks. However, the construction of the multi-plane equations for large networks turns out to be difficult. In this Letter we introduce a method based on moment equations which exhibits crucial advantages over the multi-plane method. The number of equations is further reduced to the smallest possible set of stochastic equations, including one equation for the population size of each reactive specie (represented by a first moment) and one equation for each reaction rate (represented by a second moment). Thus, for typical sparse networks the complexity of the stochastic simulation becomes comparable to that of the rate equations. Unlike the master equation (and the multi-plane method) there is no need to adjust the cutoffs - the same set of equations applies under all physical conditions. Unlike the multi-plane equations, the moment equations are linear and for steady state conditions can be easily solved using algebraic methods. Moreover, for any given network the moment equations can be easily constructed using a diagrammatic approach, which can be automated \cite{Barzel2007}. | In summary, we have introduced a method, based on moment equations, for the simulation of chemical networks taking place on dust-grain surfaces in interstellar clouds. The method provides highly efficient simulations of complex reaction networks under the extreme conditions of low gas density and sub-micron grain sizes, in which the reaction rates are dominated by fluctuations and stochastic simulations are required. The number of equations is reduced to one equation for each reactive specie and one equation for each reaction, which is the lowest possible number for such networks. This method enables us to efficiently simulate networks of any required complexity without compromising the accuracy. It thus becomes possible to incorporate the complete network of surface reactions into gas-grain models of interstellar chemistry. To fully utilize the potential of this method, further laboratory experiments are needed, that will provide the activation energy barriers for diffusion, desorption and reaction processes not only for hydrogen but for all the molecules involved in these networks. We thank A. Lipshtat for helpful discussions. This work was supported by the Israel Science Foundation and the Adler Foundation for Space Research. | 7 | 10 | 0710.2250 |
0710 | 0710.1409_arXiv.txt | {% The origin of low-luminosity Type IIP supernovae is unclear: they have been proposed to originate either from massive ($\sim 25~M_{\sun}$) or low-mass ($\sim 9~M_{\sun}$) stars. }{% We wish to determine parameters of the low-luminosity Type IIP supernova 2003Z, to estimate a mass-loss rate of the presupernova, and to recover a progenitor mass. }{% We compute the hydrodynamic models of the supernova to describe the light curves and the observed expansion velocities. The wind density of the presupernova is estimated using a thin shell model for the interaction with circumstellar matter. }{% We estimate an ejecta mass of $14.0\pm1.2~M_{\sun}$, an explosion energy of $(2.45\pm0.18)\times10^{50}$ erg, a presupernova radius of $229\pm39~R_{\sun}$, and a radioactive $^{56}$Ni amount of $0.0063\pm0.0006~M_{\sun}$. The upper limit of the wind density parameter in the presupernova vicinity is $10^{13}$ g\,cm$^{-1}$, and the mass lost at the red/yellow supergiant stage is $\leq 0.6~M_{\sun}$ assuming the constant mass-loss rate. The estimated progenitor mass is in the range of $14.4-17.4~M_{\sun}$. The presupernova of SN~2003Z was probably a yellow supergiant at the time of the explosion. }{% The progenitor mass of SN~2003Z is lower than those of SN~1987A and SN~1999em, normal Type IIP supernovae, but higher than the lower limit of stars undergoing a core collapse. We propose an observational test based on the circumstellar interaction to discriminate between the massive ($\sim 25~M_{\sun}$) and moderate-mass ($\sim 16~M_{\sun}$) scenarios. } | \label{sec:intro} Type II plateau supernovae (SNe~IIP) with the plateau of $\sim 100$ days in the light curve are believed to be an outcome of a core collapse of the $9-25~M_{\sun}$ stars (e.g., Heger et al. \cite{HFWLH_03}). This paradigm assuming the Salpeter mass spectrum suggests that about 66\% of all SNe~IIP should be produced by progenitors, i.e., stars on the main sequence, in the $9-15~M_{\sun}$ range. At present this general picture remains unconfirmed. It could be verified via the determination of ejecta masses from the hydrodynamic modeling for a sufficiently large sample of SNe~IIP. Unfortunately, only few SNe~IIP have the well-observed light curves and spectra needed to reliably reconstruct the basic SN parameters. It is not therefore surprising that up to now ejecta mass has been determined using the detailed hydrodynamic simulations for only SN~1987A and SN~1999em. It is noteworthy that, from the point of view of explosion mechanism, SN~1987A is a normal SN~IIP; it has the ejecta mass, the explosion energy, and the amount of ejected $^{56}$Ni comparable to those of SN~1999em. Recently an interesting subclass of SNe~IIP, so called low-luminosity SNe~IIP, was selected observationally (Pastorello et al. \cite{PZT_04}). This family is characterized by luminosities, expansion velocities, and radioactive $^{56}$Ni masses which are significantly lower than those of normal SNe~IIP. Two views on the origin of low-luminosity SNe~IIP have been proposed. Turatto et al. (\cite{TMY_98}) have suggested the origin from massive stars, $\geq~25~M_{\sun}$, which presumably form black holes and eject a low amount of $^{56}$Ni; alternatively, these SNe might originate from low-mass stars, $\sim 9~M_{\sun}$, which are expected to eject a low amount of $^{56}$Ni (Chugai \& Utrobin \cite{CU_00}; Kitaura et al. \cite{KJH_06}). The investigation of the origin of these SNe~IIP began with some confusion. The point is that SN~1997D, the first low-luminosity SN~IIP, was detected long after the explosion and, therefore, was erroneously claimed to possess a short ($\sim 50$ days) plateau (Turatto et al. \cite{TMY_98}). This, in turn, provoked a conclusion that SN~1997D originated from a low-mass ($\sim 9~M_{\sun}$) main-sequence star (Chugai \& Utrobin \cite{CU_00}). The subsequent discovery of several low-luminosity SNe~IIP with a long plateau of $\sim 100$ days (Pastorello et al. \cite{PZT_04}) falsified this conclusion. Until now there were no attempts to model hydrodynamically low-luminosity SNe~IIP on the basis of new observational data. The estimate of the ejecta masses of low-luminosity SN~1999br and SN~2003Z was made only using a semi-analytical model (Zampieri et al. \cite{ZPT_03}; Zampieri \cite{Zam_05}). The semi-analytical model, being a sensible tool for the first order estimates, cannot, however, substitute the hydrodynamic simulations. In this paper we study the well-observed low-luminosity Type IIP SN~2003Z (Pastorello \cite{Pas_03}; Knop et al. \cite{KHBD_07}). Our approach is based on the hydrodynamic modeling of the light curves and expansion kinematics. This paper is organized as follows. The observational data are presented in Sect.~\ref{sec:obsdat}. Section~\ref{sec:model} describes the modeling procedure. Results, specifically, the SN~2003Z parameters and progenitor mass are presented in Sect.~\ref{sec:results}. The implications of the results are discussed in Sect.~\ref{sec:discon}. A distance to SN~2003Z of 21.68 Mpc is adopted using the Hubble constant $H_0=70$ km\,s$^{-1}$\,Mpc$^{-1}$ and a recession velocity of the host galaxy NGC 2742 $v_{\mathrm{cor}}=1518$ km\,s$^{-1}$, corrected for the Local Group infall to the Virgo cluster and taken from the Lyon Extragalactic Data base. There are no observational signatures of the interstellar absorption in the host galaxy (Pastorello \cite{Pas_03}) so a total extinction is taken to be equal to the Galactic value $A_{B}=0.167$ (Schlegel et al. \cite{SFD_98}). | \label{sec:discon} The goal of this study was to recover the basic parameters of SN 2003Z and to get an idea about progenitor masses of low-luminosity SNe~IIP. We estimated the ejecta mass to be $14.0\pm1.2~M_{\sun}$, the explosion energy $(2.45\pm0.18)\times10^{50}$ erg, the pre-SN radius $229\pm39~R_{\sun}$, and the $^{56}$Ni mass $0.0063\pm0.0006~M_{\sun}$. Using the ejecta/wind interaction model, we found the upper limit of the density parameter of the pre-SN wind and, assuming the constant mass-loss rate at the RSG/YSG stage, constrained the progenitor mass by the range of $14.4-17.4~M_{\sun}$. The estimate of the wind density thus allows us to avoid the uncertainty in the progenitor mass stemming from the weak sensitivity of the hydrodynamic model to the helium core mass. The only normal SN~IIP, studied in a similar way to SN~2003Z, is SN~1999em. Its pre-SN radius is $\approx 500~R_{\sun}$, the ejecta mass is $\approx 19~M_{\sun}$, the explosion energy is $\approx 1.3\times10^{51}$ erg, and the $^{56}$Ni mass is $\approx 0.036~M_{\sun}$ (Utrobin \cite{Utr_07}), while the progenitor mass is $\approx 22~M_{\sun}$ (Chugai et al. \cite{CCU_07}). A comparison of SN~2003Z with SN~1999em taken together with the fact that the low-luminosity SNe~IIP are very similar indicates that this variety of SNe~IIP originates from less massive progenitors than normal SNe~IIP. Parameters of three SNe~IIP --- SN~1987A (Utrobin \cite{Utr_05}), SN~1999em (Utrobin \cite{Utr_07}; Chugai et al. \cite{CCU_07}), and SN~2003Z (Table \ref{tab:sumtab}) --- determined on the basis of the similar hydrodynamic modeling, suggest a picture in which ordinary SNe~IIP originate from massive progenitors around $20~M_{\sun}$, while low-luminosity SNe~IIP originate from less massive progenitors around $\sim 16~M_{\sun}$, close to the average mass of the $9-25~M_{\sun}$ range traditionally associated with SNe~IIP. If our result for SN~2003Z is correct, then the explosion energy and the amount of ejected $^{56}$Ni should significantly increase for the progenitors between $\sim 16~M_{\sun}$ and $\sim 20~M_{\sun}$ (Fig.~\ref{fig:nienms}). Interestingly, the empirical correlation between the explosion energy and the $^{56}$Ni mass for normal SNe~IIP has been demonstrated by Nadyozhin (\cite{Nad_03}). It should be emphasized that these correlations, valid for SNe~IIP, may not be applicable to other SNe~II. At least SN~1994W (Type IIn event) was found to have a low amount of ejected $^{56}$Ni, $\sim 0.015~M_{\sun}$, but a normal explosion energy, $\approx 1.3\times10^{51}$ erg (Chugai et al. \cite{CBC_04}). \begin{table}[t] \caption[]{Hydrodynamic models for SN 1987A, SN 1999em, and SN~2003Z.} \label{tab:sumtab} \centering \begin{tabular}{@{ } c @{ } c @{ } c @{ } c @{ } c @{ } c @{ } c @{ } c @{ }} \hline\hline \noalign{\smallskip} SN & $R_0$ & $M_{env}$ & $E$ & $M_{\mathrm{Ni}}$ & $v_{\mathrm{Ni}}^{max}$ & $v_{\mathrm{H}}^{min}$ & $M_\mathrm{ms}$ \\ & $(R_{\sun})$ & $(M_{\sun})$ & ($10^{51}$ erg) & $(10^{-2} M_{\sun})$ & (km\,s$^{-1}$) & (km\,s$^{-1}$) & $(M_{\sun})$ \\ \noalign{\smallskip} \hline \noalign{\smallskip} 87A & 35 & 18 & 1.5 & 7.65 & 3000 & 600 & 21.5 \\ 99em & 500 & 19 & 1.3 & 3.60 & 660 & 700 & 22.2 \\ 03Z & 229 & 14 & 0.245 & 0.63 & 535 & 360 & 15.9 \\ \noalign{\smallskip} \hline \end{tabular} \end{table} Two alternative conjectures about the origin of low-luminosity SNe~IIP have been proposed: (1) progenitors are massive stars, $\geq 25~M_{\sun}$ (Turatto et al. \cite{TMY_98}); (2) these SNe originate from low-mass stars, $\sim 9~M_{\sun}$ (Chugai \& Utrobin \cite{CU_00}; Kitaura et al. \cite{KJH_06}). The present estimate of the progenitor mass of SN~2003Z, $14.4-17.4~M_{\sun}$, is in the disparity with both the high-mass and low-mass scenarios. Our estimate of the progenitor mass is essentially based on the assumption that the mass-loss rate at the RSG/YSG stage was constant. In fact, the derived wind density refers to the close vicinity of SN~2003Z $r<R_{\mathrm w}=vt\sim 10^{16}$ cm (where $v\approx7\times10^8$ cm s$^{-1}$ is the SN expansion velocity in the outer layers and $t\sim10^7$ sec is the characteristic age of the SN). This linear scale with the wind velocity $u>10$ km\,s$^{-1}$ corresponds to the wind history during the latest $R_{\mathrm w}/u <300$ yr before the explosion. Most of the RSG/YSG stage, i.e. $4\times10^5-10^6$ yr, is thus hidden from our sight. One might suggest that the mass-loss rate was substantially more vigorous in the past than immediately before the SN explosion. If this questionable possibility were the case, the $25~M_{\sun}$ progenitor would lose $\sim 10~M_{\sun}$, and the high-mass scenario for SN~2003Z would be saved. Of note, the essentially higher wind density of pre-SN in this case suggests that the dense wind at the large radii could be revealed via radio and X-ray observations at the large age, $t\geq10$ yr. A remarkable property of all three SNe is strong mixing between the helium core and the hydrogen envelope indicated by a low minimal velocity of hydrogen matter (Table \ref{tab:sumtab}). Generally, the model with the unmixed helium core shows a bump in the light curve at the end of the plateau (Utrobin \cite{Utr_07}); the absence of the bump in the available light curves of SNe~IIP, in addition to the narrow-topped H$\alpha$ emission at the nebular phase, indicates that substantial mixing between the helium core and the hydrogen envelope is a universal phenomenon for SNe~IIP. This mixing could be induced by the convection during the growth of helium core (Barkat \& Wheeler \cite{BW_88}) or/and during the SN explosion (Kifonidis et al. \cite{KPSJM_06}). \begin{figure}[t] \resizebox{\hsize}{!}{\includegraphics{fig9.eps}} \caption{% Explosion energy (\textbf{a}) and $^{56}$Ni mass (\textbf{b}) versus progenitor mass for three core-collapse SNe. } \label{fig:nienms} \end{figure} It is noteworthy that the radii of all three pre-SNe~IIP (Table \ref{tab:sumtab}) are notably lower than a radius of typical massive RSG, e.g., the radius of $\approx 800~R_{\sun}$ for $\alpha$ Ori (Harper et al. \cite{HBL_01}). We have already shown that the SN~2003Z pre-SN was the YSG rather than the RSG. For SN~1999em the $22~M_{\sun}$ pre-SN with the luminosity of $\approx 10^5~L_{\sun}$ is characterized by the effective temperature of $\sim 4700$~K, which is intermediate between the RSG and YSG values. Interestingly, the pre-SN of SN~2004et, a normal SN~IIP, was probably a YSG as well (Li et al. \cite{LVFC_05}). These cases suggest that not only the pre-SN~1987A, but a possibly significant fraction of pre-SNe~IIP, experience an excursion on the Hertzsprung-Russell (HR) diagram from the RSG towards the YSG before the explosion. Note, some pre-SNe~IIP, when becoming the YSG, could get into the instability strip and manifest themselves before the explosion as long period ($P\sim50-100$ days) Cepheids. If this is the case, the recovery of the Cepheid period of a pre-SN on the basis of a set of pre-explosion observations could provide us with an additional tool for the mass determination from the pulsation period. The recovered progenitor mass of SN~2003Z raises a crucial question: what happens to the $9-15~M_{\sun}$ stars which make up $\approx66$\% of all the stars from the $9-25~M_{\sun}$ mass range. Two conceivable suggestions are: (A) all the $9-15~M_{\sun}$ stars explode as low-luminosity, or even fainter, SNe~IIP; (B) low-luminosity SNe~IIP originate only from a narrow range of progenitors around $\sim 16~M_{\sun}$, while the rest of the $9-15~M_{\sun}$ stars produce different varieties of SNe~II. A large sample of SNe~II which is well observed and studied is certainly needed to distinguish between the options A and B. It should be emphasized that the initial peak, the end of the plateau, and the radioactive tail are of crucial importance for the recovery of reliable SN~II parameters. An alternative approach to determine the pre-SN mass exploits archival pre-explosion images of host galaxies and stellar evolution tracks on the HR diagram. Interestingly, using the archival HST images Van Dyk et al. (\cite{VLF_03}) and Maund \& Smartt (\cite{MS_05}) derived an upper limit of $\sim 15~M_{\sun}$ and $\sim 12~M_{\sun}$, respectively, for the progenitor of the low-luminosity SN~1999br. Given a similarity of known low-luminosity SNe~IIP these estimates make the high-mass scenario unlikely. It should be emphasized that a test of the mass determination method based on the pre-explosion images is needed; for instance, the pre-SN light could be partially absorbed in a dusty circumstellar envelope. In this regard it would be of top priority to determine independently the progenitor mass both from the pre-discovery images and from the hydrodynamic simulations. Among SNe~IIP, except for SN~1987A, only SN~2004et and SN~2005cs were detected in the pre-discovery images (Li et al. \cite{LVFC_05}; Maund et al. \cite{MSD_05}; Li et al. \cite{LVF_06}) and have the light curves and spectra appropriate for the hydrodynamic modeling as well (Sahu et al. \cite{SASM_06}; Misra et al. \cite{MPC_07}; Pastorello et al. \cite{PST_06}; Tsvetkov et al. \cite{TVS_06}). These two challenging cases require a detailed study. | 7 | 10 | 0710.1409 |
0710 | 0710.1123_arXiv.txt | Optical and near-infrared spectroscopy of the newly discovered peculiar L dwarf {\name} are presented. Folkes et al.\ identified this source as a high proper motion L9$\pm$1 dwarf based on its strong H$_2$O absorption at 1.4~$\micron$. We find that the optical spectrum of {\namesh} is in fact consistent with that of a normal L4.5 dwarf with notably enhanced FeH absorption at 9896~{\AA}. However, its near-infrared spectrum is unusually blue, with strong H$_2$O and weak CO bands similar in character to several recently identified ``blue L dwarfs''. Using {\namesh} as a case study, and guided by trends in the condensate cloud models of Burrows et al.\ and Marley et al., we find that the observed spectral peculiarities of these sources can be adequately explained by the presence of thin and/or large-grained condensate clouds as compared to normal field L dwarfs. Atypical surface gravities or metallicities alone cannot reproduce the observed peculiarities, although they may be partly responsible for the unusual condensate properties. We also rule out unresolved multiplicity as a cause for the spectral peculiarities of {\namesh}. Our analysis is supported by examination of {\em Spitzer} mid-infrared spectral data from Cushing et al.\ which show that bluer L dwarfs tend to have weaker 10~$\micron$ absorption, a feature tentatively associated with silicate oxide grains. With their unique spectral properties, blue L dwarfs like {\namesh} should prove useful in studying the formation and properties of condensates and condensate clouds in low temperature atmospheres. | L dwarfs comprise one of the two latest-type spectral classes of very low mass stars and brown dwarfs, spanning masses at and below the hydrogen burning minimum mass (see \citealt{kir05} and references therein). They are inexorably linked to the presence and properties of liquid and solid condensates which form in their cool photospheres (e.g., \citealt{tsu96,bur99,ack01,all01,coo03,tsu05}). These condensates significantly influence the spectral energy distributions and photospheric gas abundances of L dwarfs, by removing gaseous TiO and VO from the photosphere and enabling the retention of atomic alkali species (e.g., \citealt{feg96,bur99,lod02}). Weakened {\wat} absorption through backwarming effects (e.g., \citealt{jon97,all01}) and red near-infrared colors ($J-K$ $\approx$ 1.5--2.5; \citealt{kir00}) also result from condensate opacity. In addition, periodic and aperiodic photometric variability observed in several L dwarfs has been associated with surface patchiness in photospheric condensate clouds (e.g., \citealt{bai99,bai01,gel02,moh02}). Condensate abundances at the photosphere appear to reach their zenith amongst the mid- and late-type L dwarfs \citep{kir99,cha00,ack01} before disappearing from the photospheres of cooler T dwarfs \citep{mar96,tsu96b,all01,cus06}. The abundances of photospheric condensates, their grain size distribution, and the radial and surface structure of condensate clouds may vary considerably from source to source, as well as temporally for any one source, and the dependencies of these variations on various physical parameters are only beginning to be explored \citep{hel01,woi03,kna04}. With hundreds of L dwarfs now known,\footnote{A current list is maintained at \url{http://dwarfarchives.org}.} groupings of peculiar L dwarfs -- sources whose spectral properties diverge consistently from the majority of field objects -- are becoming distinguishable. Examples include young, low surface gravity brown dwarfs \citep{mcg04,kir06,all07,cru07} and metal-poor L subdwarfs \citep{me0532,lep1610,giz06,megmos}. There also exists a class of peculiar ``blue'' L dwarfs \citep{cru03,cru07,kna04}, roughly a dozen sources exhibiting normal optical spectra but unusually blue near-infrared colors and strong near-infrared {\wat}, FeH and {\ki} features. Various studies have attributed these peculiarities to subsolar metallicity, high surface gravity, unresolved multiplicity and peculiar cloud properties \citep{giz00,cru03,cru07,mcl03,mewide3,kna04,chi06,fol07}. Any one of these characteristics may impact the presence and character of condensates and condensate clouds in low temperature atmospheres. In an effort to identify new nearby and peculiar L dwarfs, we have been searching for late-type dwarfs using near-infrared imaging data from the Deep Near Infrared Survey of the Southern Sky (DENIS; \citealt{epc97}). One of the objects identified in this program is DENIS~J112639.9$-$500355, a bright source which was concurrently discovered by \citet{fol07} in the SuperCOSMOS Sky Survey \citep[hereafter SSS]{ham01a,ham01b,ham01c} and the Two Micron All Sky Survey (hereafter 2MASS; \citealt{skr06}). It is designated {\name} in that study, and we refer to the source hereafter as {\namesh}. Based on its blue near-infrared colors and deep {\water} absorption bands, \citet{fol07} concluded that {\namesh} is a very late-type L dwarf (L9$\pm$1) which may have unusually patchy or thin condensate clouds. In this article, we critically examine the observational properties of {\namesh} to unravel the origins of its spectral peculiarities, and examine it as a representative of the blue L dwarf subgroup. Our identification of {\namesh} and a slightly revised determination of its proper motion using astrometry from the SSS, DENIS and 2MASS catalogs are described in $\S$~2. Optical and near-infrared spectroscopic observations and their results are described in $\S$~3, along with determination of the optical and near-infrared classifications of {\namesh} and estimates of its distance and space kinematics. In $\S$~4 we analyze the properties of {\namesh} and blue L dwarfs in general, considering metallicity, surface gravity, condensate cloud and unresolved multiplicity effects. We also introduce a new near-infrared {\wat} index that eliminates discrepancies between optical and near-infrared types for these sources. Results are discussed and summarized in $\S$~5. | Our analysis in $\S$~4.2 leads us to conclude that the spectral peculiarities of {\namesh} and other blue L dwarfs have their immediate cause in condensate cloud effects, specifically the presence of thin, patchy or large-grained condensate clouds at the photosphere. Subsolar metallicities and high surface gravities in of themselves cannot reproduce the observed spectral peculiarities of these sources. However, it is clear that these latter physical properties must play a role in determining the cloud characteristics of blue L dwarfs. Lower metallicities reduce the metal species available to form condensates, resulting in less condensate material overall. Higher surface gravities may increase the sedimentation rate of condensate grains, potentially resulting in thinner clouds. The large tangential velocities and absence of {\lii} absorption in the three blue L dwarfs {\namesh}, 2MASS~J1300+1921 and 2MASS~J1721+3344 support the idea that these sources may be relatively old and possibly slightly metal-poor. However, the influence of other physical parameters on condensate cloud properties must also be considered, including rotation rates, vertical upwelling rates (e.g., \citealt{sau06}) and possibly magnetic field strengths. An assessment of how these fundamental physical parameters influence the properties of condensate clouds in low-temperature atmospheres is the subject of ongoing theoretical investigations (e.g., \citealt{hel01,woi03}; M.\ Marley, in preparation). Empirical studies are also necessary, particularly those focused on well-characterized samples of blue (and red) L dwarfs. To this end, Table~\ref{tab_blue} lists all blue L dwarfs currently reported in the literature. We anticipate that this list will grow as near-infrared spectroscopic follow-up of L dwarfs continues. | 7 | 10 | 0710.1123 |
0710 | 0710.1065_arXiv.txt | We use a semi-analytic circumstellar disk model that considers movement of the snow line through evolution of accretion and the central star to investigate how gas giant frequency changes with stellar mass. The snow line distance changes weakly with stellar mass; thus giant planets form over a wide range of spectral types. The probability that a given star has at least one gas giant increases linearly with stellar mass from 0.4\,$M_\odot$ to 3\,$M_\odot$. Stars more massive than 3\,$M_\odot$ evolve quickly to the main-sequence, which pushes the snow line to 10--15\,AU before protoplanets form and limits the range of disk masses that form giant planet cores. If the frequency of gas giants around solar-mass stars is 6\%, we predict occurrence rates of 1\% for 0.4\,$M_\odot$ stars and 10\% for 1.5\,$M_\odot$ stars. This result is largely insensitive to our assumed model parameters. Finally, the movement of the snow line as stars $\gtrsim$2.5\,$M_\odot$ move to the main-sequence may allow the ocean planets suggested by \citeauthor{2004Icar..169..499L} to form without migration. | \label{sec:intro} In the last ten years, the discovery of more than 200 extra-solar planets,\footnote{http://vo.obspm.fr/exoplanetes/encyclo/encycl.html} and more than 200 debris disks,\footnote{http://www.roe.ac.uk/ukatc/research/topics/dust/identification.html } suggests that planet formation is a common and robust process. Planet masses inferred from debris disks range from terrestrial to Jovian, at distances as great as tens of AU from the central star \citep[e.g.][]{2004AJ....127..513K,2005ApJ...619L.187G}. The nature and sensitivity of radial velocity surveys means that most of the planets are $\sim$Jupiter mass gas giants in close orbits around Sun-like stars. However, recent discoveries as diverse as icy $\sim$Neptune-mass planets orbiting M dwarfs \citep[e.g.][]{2005ApJ...634..625R}, and debris disks around A-type stars \citep[e.g.][]{2005ApJ...620.1010R} show that planet formation occurs over a wide range of spectral types. Current theory suggests that planets form in similar ways around all stars. Thus, the increasing diversity of stellar hosts and planetary systems provides an opportunity to test these theories. For this reason, the types of planets most likely to form around stars of differing spectral types has become a renewed area of study \citep[e.g.][]{2005ApJ...626.1045I,2006ApJ...643..501B,2006A&A...458..661K,2006ApJ...650L.139K}, after the idea was first explored by \citeauthor{1988MNRAS.235..193N} nearly 20 years ago \citep{1987MNRAS.224..107N,1988MNRAS.230..551N,1988MNRAS.235..193N}. Theories of Solar System formation generally include the ``snow line,'' where ices condense from the nebular gas. The snow line distance is usually fixed in a disk with a time independent surface density and temperature profile around a main-sequence star \citep[e.g.][]{2005ApJ...626.1045I}. In a more realistic picture, the disk and stellar properties evolve considerably during the 1--10\,Myr pre--main-sequence (PMS) lifetime when planets probably form \citep[e.g.][]{1987Icar...69..249L,1996Icar..124...62P}. As the disk temperature evolves with time, movement of the snow line may therefore influence the properties of theoretical planetary systems \citep[e.g.][]{2006ApJ...650L.139K,2007ApJ...654..606G}. Here, we begin to develop a time dependent model for the formation of gas giant cores that considers the PMS evolution of the star and surrounding accretion disk. We introduce a simple semi-analytic disk model, based on the ``minimum mass solar nebula,'' that links movement of the snow line through evolution of disk accretion and stellar luminosity. In contrast to previous studies \citep[e.g.][]{2005ApJ...626.1045I,2006A&A...458..661K}, our analysis suggests that gas giant formation around stars more massive than the sun is more likely than around less massive stars. We cover the background important to our story in \S \ref{sec:background}, consider the snow line in \S \ref{sec:motivation}, and outline our model in \S \ref{sec:model}. We present our results in \S \ref{sec:results}, and discuss and conclude in \S \ref{sec:discussion} and \S \ref{sec:summary}. | \label{sec:summary} We describe a model for the evolution of the snow line in a planet forming disk, and apply it over a range of stellar masses to derive the probability distribution of gas giants as a function of stellar mass. The two main ingredients for our model are a prescription for movement of the snow line due to accretion and PMS evolution, and rules that determine whether protoplanets are massive enough, and form early enough, to become gas giants. The snow line distance generally moves inward over time. With our prescription for the accretion rate, accretion dominates over irradiation for stars with $M_\star \lesssim 2\,M_\odot$. For $\gtrsim$3\,$M_\odot$ stars, irradiation dominates at times $\gtrsim$1\,Myr as the star moves up to its main-sequence luminosity. The transition is at a few Myr for $\sim$2\,$M_\odot$ stars. Over the wide range of observed accretion rates for any fixed stellar mass, the snow line in some disks may be set entirely by irradiation. The snow line generally sets where the innermost gas giant cores form. In relatively massive disks around intermediate mass stars, rocky cores form interior to the snow line. The location of the outermost core is always set by the gas dissipation timescale. The range of disk masses that form cores, and the radial width of the region in the disk where they form, increase with stellar mass. Lower mass disks produce failed icy cores, which are probably similar to Uranus, Neptune, and the observed ``super-Earths.'' The outward movement of the snow line as stars more massive than the Sun reach the main-sequence, and as the disk becomes optically thin, allows the ocean planets suggested by \citet{2004Icar..169..499L} to form \emph{in situ}. The change in disk temperature is only large enough for these planets to harbour oceans around stars $\gtrsim$2.5\,$M_\odot$. Our model includes several poorly determined parameters, which current and future facilities will investigate. While there are current resolved studies of gaseous disks \citep[e.g.][]{2007arXiv0704.1481B}, the next generation of telescopes such as GMT and ALMA will provide more information on surface density profiles, and how disk properties change with stellar mass and age. These studies will help to constrain input parameters for our model. We have shown that the time dependence of the snow line in part determines where gas giant cores form. This result should motivate future studies of planet formation in disks whose properties change with time. The subsequent evolution of isolated cores is beyond the scope of this paper, but further work that investigates the growth and dynamical evolution of these objects can investigate the diversity of resulting system structures. Given an initial distribution of disk masses, the probability that a star has at least one gas giant increases linearly with stellar mass from 0.4\,$M_\odot$ to 3\,$M_\odot$. If the frequency of gas giants around solar-mass stars is 6\%, we predict an occurrence rate of 1\% (10\%) for 0.4\,$M_\odot$ (1.5\,$M_\odot$) stars. This result is largely insensitive to changes in our model parameters. In contrast to the \citet{2005ApJ...626.1045I} model, where it is hard to form observable gas giants above 1\,$M_\odot$, our model predicts a peak at $\sim$3\,$M_\odot$ because we include disk and PMS evolution in our snow line derivation. However, our model does not include migration, so our prediction applies to observable and currently undetectable gas giants. Though sample numbers are small, it appears that observable gas giant frequency increases with stellar mass across a wide range of host masses \citep{2007arXiv0707.2409J}. Larger samples of stars that host giant planets, particularly low and intermediate-mass stars, will solidify this result. These studies, and the extension of the results to a wider range of semi-major axes, will provide a basis for comparison with our model predictions. | 7 | 10 | 0710.1065 |
0710 | 0710.5271_arXiv.txt | The unveiled main-sequence splitting in $\omega$ Centauri as well as NGC 2808 suggests that matter highly-enriched in He (in terms of its mass fraction $Y\sim0.4$) was produced and made the color of some main-sequence stars bluer in these globular clusters (GCs). The potential production site for the He-rich matter is generally considered to be massive AGB stars that experience the second dredge-up. However, it is found that massive AGB stars provide the matter with $Y\sim 0.35$ at most, while the observed blue-shift requires the presence of $Y\sim 0.4$ matter. Here, we show that extra mixing, which operates in the red giant phase of stars less massive than $\sim2\,M_{\odot}$, could be a mechanism that enhances He content in their envelopes up to $Y\sim 0.4$. The extra mixing is supposed to be induced by red giant encounters with other stars in a collisional system like GCs. The $Y\sim 0.4$ matter released in the AGB phase has alternative fates to (i) escape from a GC or (ii) be captured by kinematically cool stars through encounters. The AGB ejecta in $\omega$ Cen, which follows the latter case, can supply sufficient He to cause the observed blue-shift. Simultaneously, this scheme generates the extreme horizontal branch, as observed in $\omega$ Cen in response to the higher mass loss rates, which is also caused by stellar encounters. | The discovery of a split in the main sequence (MS) of \omcen\ \citep{Bedin04} has opened a new window in the study of Galactic globular clusters (GCs). The observed fact that stars of the blue MS (bMS) in \omcen\ exhibit slightly stronger absorption of iron lines on average, in comparison with the red MS (rMS) \citep{Piotto05} which strongly implies that the origin of bMS is attributable to the enhancement of He inside bMS stars. Subsequently, the recent HST observation has provided us with the second sample of MS splitting, NGC 2808 \citep{Piotto07}. This GC essentially exhibits no dispersion of [Fe/H] and thus the argument of He enhancement in bMS stars has been reinforced. A comparison between the observed color-magnitude diagrams (CMDs) and synthetic population models suggests that bMS He abundance is enhanced to $Y\sim 0.4$ for both of two GCs \citep{Norris04, Lee05, DAntona05, Piotto07}. \citet{Tsujimoto07} have argued that if the surface of a star on the MS is polluted with the $Y=0.4$ matter with the mass of $\sim 0.1 \msun$, the star can move to a position on the observed bMS in the CMD. This relaxes a severe demand on the amount of He possibly supplied from asymptotic giant branch (AGB) stars, in contrast to the idea that the bMS stars are born from pure AGB ejecta consisting of $Y\sim 0.4$. Similarly, \citet{Newsham07} have insisted that surface pollution may explain the high He content of bMS stars. It should be stressed that $Y\sim 0.4$ matter is necessary to realize the bMS stars in these GCs. Since massive AGB stars can enhance the He abundance in their envelopes through the second dredge-up in the early AGB phase, they have been considered a possible production site for the $Y\sim 0.4$ matter \citep{DAntona05}. However, previous studies indicated that the resultant abundance of He in the envelope of any AGB star does not exceed $Y \sim 0.35$ \citep[see e.g.,][]{vandenHoek97}. Therefore, such a He yield from massive AGB stars may not be sufficient to split the MS as observed. In the subsequent section, we will discuss this issue and conclude that it is highly implausible for a massive AGB star to enhance He abundance up to $Y\sim 0.4$ in its envelope. Then, we should search for an alternative mechanism that enhances the He abundance more efficiently than the second dredge-up. We believe that this process, which enhances He abundance, should be accompanied by some other elemental signatures. This is reminiscent of the abundance anomaly of red giants in GCs. The abundance anomalies observed for red giants, such as the abundance variations of CNO elements and the O-Na anticorrelation, are common attributes associated with GCs including \omcen\ \citep[see][]{Gratton04}. One of the proposed mechanisms to explain these anomalies is a deep mixing, so-called extra mixing, on the red giant stage \citep{Sweigart79, Langer93, Charbonnel98, Fujimoto99,Aikawa01,Chaname05,Suda06}. The extra mixing is assumed to be caused by some kind of rotation-induced mixing, driven by shear instability around the border of the He core \citep[see][]{Fujimoto88,Zahn92}. Since such a deep mixing naturally draws He from the core, the extra mixing is a promising mechanism of producing He-rich matter. As a mechanism of mixing, most of the previous works assumed a continuous mixing caused by the meridional circulation, while Fujimoto and his coworkers insisted that a flash of H brought into a degenerate core by the rotation-induced mixing drives the intermittent mixing. In this study, we propose that extra mixing, which occurs in the red giant phase of a star, enhances He abundance in the envelope and ejects He-rich matter when the star evolves to a white dwarf through mass loss. In this scenario, the driving force of the extra mixing is generated through interactions between the red giant and dwarfs in the central region of \omcen. Our detailed calculations reveal that stars with the masses less than $\sim$2 \msun\ can enhance He abundance up to $Y\sim 0.4$, provided that the extra mixing operates in the red giant phase. Based on our scenario, we discuss other conspicuous characteristics of \omcen, the extremely blue horizontal branch (HB), which is observed in some other Galactic GCs. Although factors that determine HB morphology is still an open question and is known as the second parameter problem, the most effective parameter that affects the location of a HB star in the CMD is the mass $M_{\rm env}$ of the envelope. A major factor to determine $M_{\rm env}$ is the mass loss during the red giant phase. We predict that encounters induce a high mass loss rate through a gain of angular momentum and promote the presence of an extremely blue HB population. In the end, two mysteries in \omcen\ of MS splitting and HB morphology are discussed in a unified framework of encounters of dwarfs with AGB ejecta and those of red giants with dwarfs. | We first propose that He-rich matter is synthesized in red giants with the masses of 0.8 - 2 $\msun$, which experience encounters with other stars in GCs, leading to the extra mixing during the RGB phase in their envelopes. On the other hand, massive AGB stars are unable to produce the He abundance as much as $Y \sim 0.4$ because (1) their large envelopes work as a buffer against He enrichment and (2) the amount of dredged-up matter is limited by thinning the He-burning shell, due to the growth of C-O core. Therefore, a $Y\sim 0.4$ matter is a product unique to GCs. In some GCs, this matter is retained in their gravitational potential for a prolonged time and is accreted by kinematically cool stars, although this may not be the case for most GCs. In addition, stellar encounters in GCs induce the extra mixing and increase the mass loss rate during the RGB phase. As a result, such stars are predicted to evolve to significantly blue HB stars. In conclusion, MS splitting is inclined to be accompanied by the existence of an extremely blue HB, as observed in \omcen\ and NGC 2808. It should be, however, noted that our models might make a variation in the final He abundance as a result of extra mixing induced by various amounts of angular momentum transfered through encounters. If the variation is too large, then it will erase the MS splitting and lead to a broad MS. Further investigations of the extra mixing model in terms of the MS splitting are surely awaited. | 7 | 10 | 0710.5271 |
0710 | 0710.0662_arXiv.txt | The study of Wolf-Rayet stars plays an important role in evolutionary theories of massive stars. Among these objects, $\sim 20\%$ are known to be in binary systems and can therefore be used for the mass determination of these stars. Most of these systems are not spatially resolved and spectral lines can be used to constrain the orbital parameters. However, part of the emission may originate in the interaction zone between the stellar winds, modifying the line profiles and thus challenging us to use different models to interpret them. In this work, we analyzed the He{\sevensize II}$\lambda$4686\AA \ + C{\sevensize IV}$\lambda$4658\AA \ blended lines of WR30a (WO4+O5) assuming that part of the emission originate in the wind-wind interaction zone. In fact, this line presents a quiescent base profile, attributed to the WO wind, and a superposed excess, which varies with the orbital phase along the 4.6 day period. Under these assumptions, we were able to fit the excess spectral line profile and central velocity for all phases, except for the longest wavelengths, where a spectral line with constant velocity seems to be present. The fit parameters provide the eccentricity and inclination of the binary orbit, from which it is possible to constrain the stellar masses. | Massive stars are known to drive strong winds, which are responsible for the transfer of a large amount of the stellar mass to the interstellar medium, contributing to the feedback of chemical elements, and to the creation of cloud cavities in which these objects are found. Typically, O stars present mass-loss rates of $\dot{M}_{\rm O} \sim 10^{-6} - 10^{-5}$ M$_\odot$ yr$^{-1}$ and wind velocities of $v_{\rm O} \sim 2000 - 3500$ km s$^{-1}$, while Wolf-Rayet (WR) stars present $\dot{M}_{\rm WR} \sim 10^{-6} - 10^{-4}$ M$_\odot$ yr$^{-1}$ and $v_{\rm WR} \sim 1000 - 4000$ km s$^{-1}$ (Nugis \& Lamers 2000, Lamers 2001). In massive binary systems in which both stars present high mass-loss rates and high-velocity winds, the collision of the winds will occur. Therefore, a contact surface is formed where the momenta of the two winds are equal, surrounded by two shocks. The post-shocked gas, cools as it flows along the contact surface and is responsible for strong free-free emission at X-ray and radio wavelengths. X-rays in massive binary systems are orders of magnitude higher than those observed in single massive stars, and are used as an indication of binarity. UV and optical lines can also indicate the binary nature of a given object, when they present periodic profile variations. If the lines are of photospheric or atmospheric origin, as the stars move along their orbit, they suffer Doppler shifts, which depend on the orbital phase and inclination. However, some massive objects show periodic variable line profiles that cannot be explained under these assumptions. Seggewiss (1974) noted that the binary system WR79 presented two peaks, superimposed to the C{\sevensize III} emission line, which changed their position and intensity with time. Typically, in double line spectroscopic binaries, each peak moves in a different direction, indicating opposite velocity components along the line of sight for each star; however in WR79 both peaks moved in the same direction. L$\rm \ddot{u}$hrs (1997) presented a model in which the two peaks were not produced by the stellar photosphere, but were generated by the flowing gas at the contact surface between the two strong shocks. This model reproduced well the data for WR 79, but failed to reproduce the line profiles of other WR binary systems, among them WR 30a (Bartzakos, Moffat \& Niemela 2001). Falceta-Gon\c calves, Abraham \& Jatenco-Pereira (2006) improved L$\rm \ddot{u}$hrs' model introducing more realistic parameters, as stream turbulence and gas opacity, to account for the line broadening and peak displacement. In the present work, we applied this model to WR 30a, which was classified as a WO4+O5 binary system (Moffat \& Seggewiss 1984, Crowther et al. 1998); its binary nature was reported by Niemela (1995) based on spectral-line radial velocities obtained with a high temporal resolution. Later, Gosset et al. (2001) presented a detailed analysis of the spectra of WR30a for several epochs and were able to confirm the binary hypothesis and to determine the period of $P \sim 4.6$ days. They also noted strong line-profile variations, which made more difficult the determination of the stellar mass ratio and the orbital inclination from standard methods. They concluded that the C{\sevensize IV}$\lambda$4658\AA \ line-profile variations were related to wind-wind collision processes, but did not model the lines under such an assumption. The same conclusion was reached by Bartzakos, Moffat \& Niemela (2001) using the C{\sevensize IV}$\lambda$5801\AA \ line. Also, Paardekooper et al. (2003) presented photometric measurements at $V$ and $B$ bands; the light curves confirmed the period obtained by Gosset et al. (2001), but also showed higher frequency variability in the $V$ band, with timescale of hours. They concluded that this could be due to the strong variability of the C{\sevensize IV}$\lambda$5801\AA \ line, possibly related to the wind-wind interaction. In the present work, we tested the wind-wind shock emission hypothesis on the C{\sevensize IV}$\lambda$4658 \AA \ excess line profile variations measured by Gosset et al. (2001) using the model developed by Falceta-Gon\c calves, Abraham \& Jatenco-Pereira (2006), which is briefly described in Section 2. In Section 3, we show the results obtained for WR 30a and present a brief discussion, followed by the conclusions in Section 4. | In this work we presented an application of the wind-wind shock emission model proposed by Falceta-Gon\c calves et al. (2006) to the line profile variations of WR 30a. In this model, the wind-wind shock structure can be represented by a cone, along which the shocked material flows. This gas will emit radiation, cool down and eventually reach recombination temperatures. The observed emission lines will then suffer Doppler shifts due to the stream velocity component along the line of sight. During the orbital movement, the cone position will change, as well as the radial velocity, which will cause the line profile variations observed in several massive binary systems. Gosset et al. (2001) obtained detailed spectra of WR30a during more than 30 days. They determined the orbital period ($P = 4.6$d), and obtained the radial velocity curve for the O-star. Regarding the WR component, they found that the blended He{\sevensize II}$\lambda$4686\AA \ and C{\sevensize IV}$\lambda$4658\AA \ lines showed a variable excess emission. In the present paper we modeled this variable emission, being able to reproduce the variations except for the red part of the profiles, which seemed to be unchanged in velocity and were probably generated in the stellar wind of the WR star instead of in the shock region. The best-fitting result was obtained for $\beta = 50^\circ$, $i = 20^\circ$, $\sigma = 0.3$ and $v_{\rm flow} = 2200$km s$^{-1}$. Also, correlating the orbital phase with the modeled phase angle, it was possible to determine the orbital eccentricity as $e = 0.2$, similar to the value of 0.0 previously assumed {\it ad hoc} by other authors. Although both values lead to very small differences in the orbital shape, its value is important for the determination of the stellar masses and orbital separation between the stars. Using this eccentricity and orbital inclination to model the radial velocity curve for the O-star, we found $K_{\rm O} = 25$ km s$^{-1}$. If we assume $M_{\rm O} = 40 - 60$ M$_{\odot}$, we find $M_{\rm WR} = 7.5 - 9.7$ M$_{\odot}$, and the orbital major semi-axis $a_{\rm O} \simeq 5.4$ R$_{\odot}$ and $a_{\rm WR} \simeq 30$ R$_{\odot}$. The model showed itself a powerful tool for constraining the wind and orbital parameters of massive binary systems. The fits did not matched the data exactly at all epochs, but considering the difficulties of subtracting the emission excess from the original spectra, the general shape and peak position variations were well reproduced. We must state that this is a very simple approximation, taking into account the complexities found in such systems. Numerical simulations could give a more detailed analysis, and probably more accurate values for the model parameters in future works. | 7 | 10 | 0710.0662 |
0710 | 0710.0381_arXiv.txt | The evolution of the galaxy stellar mass--star formation rate relationship ($M_*-$SFR) provides key constraints on the stellar mass assembly histories of galaxies. For star-forming galaxies, $M_*-$SFR is observed to be fairly tight with a slope close to unity from $z\sim 0\rightarrow 2$, and it evolves downwards roughly independently of $M_*$. Simulations of galaxy formation reproduce these trends, broadly independent of modeling details, owing to the generic dominance of smooth and steady cold accretion in these systems. In contrast, the observed amplitude of the $M_*-$SFR relation evolves markedly differently than in models, indicating either that stellar mass assembly is poorly understood or that observations have been misinterpreted. Stated in terms of a star formation activity parameter $\asf\equiv (M_*/SFR)/(t_{\rm Hubble}-{\rm 1 Gyr})$, models predict a constant $\asf\sim 1$ out to redshifts $z\sim 4+$, while the observed $M_*-$SFR relation indicates that $\asf$ increases by $\sim\times 3$ from $z\sim 2$ until today. The low $\asf$ (i.e. rapid star formation) at high-$z$ not only conflicts with models, but is also difficult to reconcile with other observations of high-$z$ galaxies, such as the small scatter in $M_*-$SFR, the slow evolution of star forming galaxies at $z\sim 2-4$, and the modest passive fractions in mass-selected samples. Systematic biases could significantly affect measurements of $M_*$ and SFR, but detailed considerations suggest that none are obvious candidates to reconcile the discrepancy. A speculative solution is considered in which the stellar initial mass function (IMF) evolves towards more high-mass star formation at earlier epochs. Following Larson (1998), a model is investigated in which the characteristic mass $\hat{M}$ where the IMF turns over increases with redshift. Population synthesis models are used to show that the observed and predicted $M_*-$SFR evolution may be brought into general agreement if $\hat{M}=0.5(1+z)^{2} M_\odot$ out to $z\sim 2$. Such IMF evolution matches recent observations of cosmic stellar mass growth, and the resulting $z=0$ cumulative IMF is similar to the ``paunchy" IMF favored by \citet{far07} to reconcile the observed cosmic star formation history with present-day fossil light measures. | How galaxies build up their stellar mass is a central question in galaxy formation. Within the broadly successful cold dark matter scenario, gas is believed to accrete gravitationally into growing dark matter halos, and then radiative processes enable the gas to decouple from non-baryonic matter and eventually settle into a star-forming disk~\citep[e.g.][]{mo98}. However, many complications arise when testing this scenario against observations, as the stellar assembly of galaxies involves a host of other processes including star formation, kinetic and thermal feedback from various sources, and merger-induced activity, all of which are poorly understood in comparison to halo assembly. A key insight into gas accretion processes is the recent recognition of the importance of ``cold mode" accretion~\citep{kat03,bir03,ker05}, where gas infalling from the intergalactic medium (IGM) does not pass through an accretion shock on its way to forming stars. Simulations and analytic models show that cold mode dominates global accretion at $z\ga 2$~\citep{ker05}, and dominates accretion in all halos with masses $\la 10^{12}$M$_\odot$~\citep{bir03,ker05}. The central features of cold mode accretion are that it is (1) {\it rapid} (2) {\it smooth}, and (3) {\it steady}. It is rapid because it is limited by the free-fall time and not the cooling time; it is smooth because most of the accretion occurs in small lumps and not through major mergers~\citep{mur01,ker05,guo07}; and it is steady because it is governed by the gravitational potential of the slowly growing halo. A consequence of cold mode accretion is that the star formation rate is a fairly steady function of time~\citep[e.g.][]{fin06,fin07}, which results in a tight relationship between the stellar mass $M_*$ and the star formation rate SFR. Hence simulations of star-forming galaxies generically predict a tight relationship with $M_*\propto$SFR that evolves slowly with redshift. In detail, a slope slightly below unity occurs owing to the growth of hot halos around higher-mass galaxies that retards accretion. The stellar mass--star formation rate ($M_*-$SFR) relation for star forming galaxies has now been observed out to $z\sim 2$, thanks to improving multiwavelength surveys. As expected from models, the relationship is seen to be fairly tight from $z\sim 0-2$, with a slope just below unity and a scatter of $\la 0.3$~dex~\citep{noe07a,elb07,dad07}. It also evolves downwards in amplitude to lower redshift in lock step fashion, suggesting a quiescent global quenching mechanism such as a lowering of the ambient cosmic density, again as expected in models. The range of data used to quantify this evolution is impressive: \citet{noe07a} used the AEGIS multi-wavelength survey in the Extended Groth Strip to quantify the $M_*$-SFR relation from $z\sim 1\rightarrow 0$; \citet{elb07} used GOODS at $z\sim 0.8-1.2$ and SDSS spectra at $z\sim 0$; and \citet{dad07} used star-forming BzK-selected galaxies with {\it Spitzer}, X-ray, and radio follow-up to study the relation at $z\sim 1.4-2.5$. In each case, a careful accounting was done of both direct UV and re-radiated infrared photons to measure the total galaxy SFR, paying particular attention to contamination from active galactic nuclei (AGN) emission (discussed in \S\ref{sec:syst}). Multiwavelength data covering the rest-optical were employed to accurately estimate $M_*$. This broad agreement in $M_*-$SFR slope, scatter, and qualitative evolution between model predictions and these data lends support to the idea that cold mode accretion dominates in these galaxies. Yet all is not well for theory. Closer inspection reveals that the {\it amplitude} of the $M_*-$SFR relation evolves with time in a way that is inconsistent with model expectations. This disagreement is fairly generic, as shown in \S\ref{sec:comparison}, arising in both hydrodynamic simulations and semi-analytic models, and is fairly insensitive to feedback implementation. The sense of the disagreement is that going to higher redshifts, the observed SFRs are higher, and/or stellar masses lower, than predicted in current models. The purpose of this paper is to understand the implications of the observed $M_*-$SFR relation for our theoretical view of how galaxies accumulate stellar mass. In \S\ref{sec:tact} it is argued that $M_*-$SFR amplitude evolution implies a typical galaxy star formation history that is difficult to reconcile with not only model expectations, but also other observations of high-$z$ galaxies. Possible systematic effects that may bias the estimation of SFR and $M_*$ are considered in \S\ref{sec:syst}, and it is argued that none of them are obvious candidates to explain the discrepancy. Finally a speculative avenue for reconciliation is considered, namely that the stellar initial mass function (IMF) has a characteristic mass that evolves with redshift (\S\ref{sec:imf}). This model is constrained based on the $M_*$-SFR relation in \S\ref{sec:imfmodel}. The implications for such an evolving IMF are discussed in \S\ref{sec:mchar}. Results are summarized in \S\ref{sec:summary}. A $\Lambda$CDM cosmology with $\Omega=0.25$ and $H_0=70$~km/s/Mpc is assumed throughout for computing cosmic timescales. | \label{sec:summary} Implications of the observed stellar mass--star formation rate correlation are investigated in the context of current theories for stellar mass assembly. The key point, found here and pointed out in previous studies~\citep{dad07,elb07}, is that the amplitude of the $M_*-$SFR relation evolves much more rapidly since $z\sim 2$ in observations that in current galaxy formation models. It is shown here that this is true of both hydrodynamic simulations and semi-analytic models, and is broadly independent of feedback parameters. In contrast, the slope and scatter of the observed $M_*-$SFR relation are in good agreement with models. The tight $M_*-$SFR relation with a slope near unity predicted in models is a generic result owing to the dominance of cold mode accretion, particularly in early galaxies, which produces rapid, smooth and relatively steady infall. The slow amplitude evolution arises because star formation starts at $z\ga 6$ for moderately massive star-forming galaxies and continues at a fairly constant or mildly declining level to low-$z$. The large discrepancy in amplitude evolution when compared to observations from $z\sim 2\rightarrow 0$ hints at some underlying problem either with the models or with the interpretation of data. A convenient parameterization of the problem is through the star formation activity parameter, $\asf\equiv (M_*/{\rm SFR})/(t_H-1\;{\rm Gyr})$, where $t_H$ is the Hubble time. A low value corresponds to a starbursting system, a high value to a passive system, and a value near unity to system forming stars constantly for nearly a Hubble time. In models, $\asf$ remains constant around unity from $z=0-4$, whereas in observations it rises steadily from $\la 0.2$ at $z\sim 2$ to close to unity at $z\sim 0$. Several ad hoc modifications to the theoretical picture of stellar mass assembly are considered in order to match $\asf$ evolution, but each one is found to be in conflict with other observations of high-redshift galaxies. Bursts seem unlikely given the low scatter in $M_*-$SFR. Delaying galaxy formation to match $\asf$ results in such a low redshift for the onset of star formation, it is in conflict with observations of star forming galaxies at earlier epochs. Hiding a large population of galaxies as passive runs into difficulty when compared to direct observations of the passive galaxy fraction in mass-selected samples at $z\sim 2$. Having an exponentially growing phase of star formation or ``staged" galaxy formation are perhaps the most plausible solutions from an observational viewpoint, but it is difficult to understand theoretically how such scenarios can arise within hierarchical structure formation. Hence if observed star-forming galaxies are quiescently forming stars and represent the majority of galaxies at $z\sim 2$, as other observations seem to suggest, then it is not easy to accomodate the low $\asf$ inferred from the $M_*-$SFR amplitude. Systematic uncertainties in $M_*$ or SFR determinations could bias results progressively more at high-$z$ in order to mimic evolution in $\asf$. Various currently debated sources are considered, such as the contribution to near-IR light from TP-AGB stars, extinction corrections, assumptions about star formation history, AGN contamination, and PAH emission calibration. It may be possible to concoct scenarios whereby several of these effects combine to mimic $\asf$ evolution, but there are no suggestions from local observations that such scenarios are to be expected. Hence if systematic effects are to explain the low $\asf$ at high-$z$, it would imply significant and unexpected changes in tracers of star formation and stellar mass between now and high redshifts. A solution is proposed that the stellar initial mass function becomes increasingly bottom-light to higher redshifts. Several lines of arguments are presented that vaguely or circumstantially favor such evolution, though no smoking gun signatures are currently known. A simple model of IMF evolution is constructed, based on the ansatz by \citet{lar98} that the minimum temperature of molecular clouds is reflected in the characteristic mass of star formation ($\hat{M}_{\rm IMF}$) where the mass contribution per logarithmic mass bin is maximized. The minimum temperature may increase with redshift owing to a hotter cosmic microwave background, more vigorous star formation activity, or lower metallicities within early galactic ISMs. In order to reconcile the observed $\asf$ evolution with theoretical expectations of no evolution, an evolving IMF of the form \begin{eqnarray*} \frac{dN}{d\log{M}} &\propto& M^{-0.3}\;\; {\rm for}\; M<\hat{M}_{\rm IMF},\\ &\propto& M^{-1.3}\;\; {\rm for}\; M>\hat{M}_{\rm IMF};\\ \hat{M}_{\rm IMF}&=&0.5 (1+z)^{2} M_\odot \end{eqnarray*} is proposed. The exponent of $\hat{M}_{\rm IMF}$ evolution is constrained by requiring no $\asf$ evolution, through careful modeling with the PEGASE.2 population synthesis code. While the exact form of IMF evolution is not well constrained by present observations, what is required is that the IMF has progressively more high-mass stars compared to low-mass at earlier epochs, and that this IMF applies to the majority of star forming galaxies at any epoch. It is worth noting that this evolving IMF is only constrained out to $z\sim 2$ from the $M_*-$SFR relation, though it yields predictions that are consistent with other observations out to $z\sim 4$. Extrapolating such evolution to higher redshifts is dangerous, since no observational constraints exist and the precise cause of IMF evolution is not understood. Implications of such an evolving IMF are investigated. By leaving the high-mass end of the IMF unchanged, recent successes in understanding the connections between high-mass star formation, feedback, and metal enrichment are broadly preserved. The cosmic stellar mass accumulation rate would be altered compared to what is inferred from cosmic star formation history measurements using a standard IMF. It is shown that an evolving IMF is at face value in better agreement with direct measures of cosmic stellar mass assembly~\citep{per07,els07,wil07}. Furthermore, the evolving IMF goes towards relieving the generic tension between present-day fossil light measures versus observations of the cosmic star formation history. In particular, the paunchy IMF favored by \citet{far07} in order to reconcile the observed cosmic star formation history, present-day $K$-band luminosity density, and extragalactic background light constraints, is qualitatively similar to the cumulative IMF of all stars formed by today in the evolving IMF case. Individually, each argument has sufficient uncertainties to cast doubt on whether a radical solution such as an evolving IMF is necessary. But taken together, the $M_*-$SFR relation adds to a growing body of circumstantial evidence that the ratio of high-mass to low-mass stars formed is higher at earlier epochs. It is by no means clear that IMF variations are the only viable solution to the $M_*-$SFR dilemma. The claim here is only that an evolving IMF is an equally (un)likely solution as invoking unknown systematic effects or carefully crafted star formation histories in order to explain $M_*-$SFR evolution. It is hoped that this work will spur further efforts, both observational and theoretical, to investigate this important issue. Future plans include performing a more careful analysis of the evolving IMF in terms of observable quantities, in order to accurately quantify the impact of such IMF evolution on the interpretation of UV, near-IR, and mid-IR light. Also, it is feasible to incorporate such an evolving IMF directly into simulation runs to properly account for gas recycling in such a scenario. Finally, including some feedback mechanism to truncate star formation in massive systems may impact the $M_*-$SFR relation in some way, so this will be incorporated into the hydro simulations. Observationally, pushing SFR and $M_*$ determinations to higher redshifts is key; a continued drop in $\asf$ to $z\sim 3$ would rapidly solidify the discrepancies with current models, and would also generate stronger conflicts with other observations of high-$z$ galaxies. Assessing the AGN contribution and extinction uncertainties from high-$z$ systems is critical for accurately quantifying the light from high-mass star formation, for instance through the use of more direct star formation indicators such as Paschen-$\alpha$. Pushing observations further into the mid-IR such as to 70$\mu$, past the PAH bands at $z\sim 2$, would mitigate PAH calibration uncertainties in SFR estimates. Obtaining a large sample of spectra for typical high-$z$ star-forming systems \citep[like cB58;][]{pet00} would more accurately constrain the SED than broad-band data. All of these programs push current technological capabilities to their limits and perhaps beyond, but are being planned as facilities continue their rapid improvement. In summary, owing to the robust form of star formation histories in current galaxy formation models, the $M_*$-SFR relation represents a key test of our understanding of stellar mass assembly. Current models reproduce the observed slope and scatter remarkably well, but broadly fail this test in terms of amplitude evolution. Whether this reflects some fundamental lack of physical insight, or else some missing ingredient such as an evolving IMF, is an issue whose resolution will have a significant impact on our understanding of galaxy formation. | 7 | 10 | 0710.0381 |
0710 | 0710.5047_arXiv.txt | {We present results of an analysis of the optical spectrum of the post-AGB star HD\,56126 (IRAS\,07134\,+\,1005) based on observations made with the echelle spectrographs of the 6-m telescope with spectral resolutions of R\,=\,25000 and 60000 at 4012--8790\,\AA. The profiles of strongest lines (HI; FeII, YII, BaII absorptions, etc.) formed in the expanding atmosphere at the base of the stellar wind have complex and variable shapes. To study the kinematics of the atmosphere, the velocities of individual features in these profiles must be measured. Differential line shifts of up to $\rm V_r$\,=\,15$\div$30\,km/s have been detected from the lines of metals and molecular features. The star's atmosphere simultaneously contains both expanding layers and layers falling onto the star. A comparison of the data for different times demonstrates that both the radial velocity and the velocity pattern in whole are variable. The position of the molecular spectrum is stable, implying stability of the expansion velocity of the circumstellar envelope around HD\,56126 detected in observations in the C$_2$ and NaI lines.} \authorrunning{Klochkova \& Chentsov} \titlerunning{HD\,56126: structure of the atmosphere and envelope} | We studied the kinematic parameters of the atmosphere of the star HD\,56126 (SAO\,96709), which belongs to the asymptotic giant branch (the post-AGB stage). In this short evolution stage, stars are in the process of their transition from the AGB to becoming a planetary nebula; for this reason, they are generally known as proto-planetary nebulae (PPNe). In the Hertzsprung--Russell diagram, post-AGB stars evolve toward the left from the AGB, maintaining nearly constant luminosity while becoming hotter. As descendants of AGB stars, these objects can be used to trace variations in the physical conditions and chemical parameters of the stellar material due to changes in the sources of energy release in the stellar interior, the ejection of the envelope, and mixing. HD\,56126 is the optical component of the IR source IRAS\,07134+1005, which has a double peaked spectral energy distribution (SED), typical for PPNe. In addition to this anomalous SED, associated with the circumstellar dust envelope, the star also displays other characteristics of this type of object [\cite{Kwok}]: the optical component of the PPN is a F5\,Iab supergiant outside the galactic plane (b=+9$\lefteqn{.}^m$99); the central star is surrounded by an extended nebula whose angular size exceeds 4$^{\rm ''}$, according to Hubble Space Telescope observations [\cite{Ueta}] (the largest known for this type of PPN); and the optical spectrum displays H$\alpha$ emission and absorption with a variable line profile [\cite{Oud1994}]. In addition to these general PPN characteristics noted by Kwok [\cite{Kwok}], subsequent studies of HD\,56126 and the associated IR source revealed several important peculiarities expected for this evolution stage: a large excess of carbon and $s$-process elements [\cite{Kloch1995}] and an emission feature at $\lambda$\,=\,21\,$\mu$ in the IR--spectrum. In the small subgroup of PPN having this emission, the presence of this 21\,$\mu$ feature was found to correlate with observational manifestations of the products of nucleosynthesis (the excess of carbon and heavy metals in the outer layers of the atmosphere) [\cite{Kloch1998, Decin}]. Thus, HD\,56126 displays all the properties ex pected for a post-AGB object, indicating the importance of detailed studies of its optical spectrum with high spectral resolution in a broad wavelength range. Fortunately, HD\,56126 is fairly bright (B\,=\,9$\lefteqn{.}^m$11, V = 8.27m), making it the most convenient carbon enriched PPN star for high resolution spectroscopy. The main moments of our study are spectral features identification and comparison the spectra of HD\,56126 and the standard $\alpha$\,Per (Sp\,=\,F5\,Iab); search for profile variability for spectral features; analyses the radial velocities $\rm V_r$ to search for differential shifts; and studying the variability of the radial velocity. In Section~2, we briefly describe the techniques used for the observations, processing, and analysis of the spectral data. The main conclusions concerning the spectrum of the star are presented in Section~3.1, while Sections~3.2 and 3.3 present our radial velocity measurements derived using various features of the spectrum, and discuss the temporal behavior of the radial-velocity pattern. Our results are summarized in Section~4. {\footnotesize \begin{table}[t] \caption{Moments of observations and values of heliocentric radial velocitity $\rm V_r$ measured. Column~4 contains $\rm V_r$ values averaged over weak lines (with depths R$_\lambda$ close to the continuum level, R$_\lambda \rightarrow$0). Velocities corresponding to the positions of the strongest components are presented for FeII(42), H$\alpha$, and D2\,NaI, with values determined from the weakest components given in parantheses. The two velocities in italics in column~5 are determined from the IR--oxygen triplet OI\,7773\,\AA. Uncertain values are marked by a colon.} \medskip \begin{tabular}{ll c| c c c l l l @{\quad} | l l l} \hline \small Date &\small Spectro-- &\small $\Delta\lambda$ & \multicolumn{7}{c}{\small $V_r$} \\ \cline{4-12} &\small graph &\small \AA{} & R$_\lambda \rightarrow $0 & \small Fe{\sc ii} &\small H$\beta$ & \quad \small H$\alpha$ & \small D\,Na{\sc i}& \small C$_2$ &\multicolumn{3}{c}{\small intestellar} \\ \hline \quad 1& 2& 3&4&5&6&\quad7& 8&9& 10& 11 & 12\\ \hline 12.01.93&Lynx&5560--8790& 88.8 &{\it 91} & &78 (100:)&77 &79: & & & \\ 10.03.93&Lynx&5560--8790& 89.0 &{\it 93} & &71 (43:) &75: &76: & & & \\ 04.03.99&Lynx&5050--6640& 85.9 & 77 & &76 (43:) &78 &77.1& & & \\ 20.11.02&NES &4560--5995& 89.6 &95 (80:) &89 & &74.9 (89) &77.2&12.0 &23.5 & 30.8\\ 21.02.03&NES &5150--6660& 88.8 & 96: & &88 (112:)&75.6 (89) &77.1&12 &24 & 31 \\ 12.04.03&NES &5270--6760& 88.4 & & &82 (103:)&75.4 (89:)& &13 &23 & 30.5 \\ 14.11.03&NES &4518--6000& 85.3 &96 (86:) &97 & &75.0 (87:)&76.9&12.5 & & \\ 10.01.04&NES &5270--6760& 86.7 & & &54: (78:)&75.6 (86:)& &13.0 &23.5 & 31 \\ 09.03.04&NES &5275--6767& 89.8 & & &58 (74:) &76.1 (89) & &13 &24 & 31 \\ 12.11.05&NES &4010--5460& 82.5 &97 (77:) &98 & & &77.5& & & \\ \hline \end{tabular} \end{table} } | Our high resolution spectra of HD\,56126 have revealed variability of various line profiles: along with the previously known H$\alpha$ profile variability, the profiles of strongest lines (such as BaII, YII, and FeII) also proved to be variable. The broad spectral range encompassed by our spectra enabled us to measure radial velocities using spectral features that form at various depths in the stellar atmosphere and circumstellar envelope. We were able to distinguish the C$_2$ absorption molecular bands (and their structure), as well as bands identified with diffuse interstellar bands. We found substantial differential shifts in lines of different intensities within the same spectrum, i.e., appreciable $\rm V_r$(R) dependences, as well as variation of these shifts with time. This indicates the need for velocity measurements based on a large set of lines in an extended spectral interval. We conclude that both expanding layers and matter falling onto the star exist simultaneously in the stellar atmosphere. The position of the molecular spectrum is stable, indicating stability of the expansion velocity of the envelope. | 7 | 10 | 0710.5047 |
0710 | 0710.2451_arXiv.txt | Recently the numerical simulations of the process of reionization of the universe at $z>6$ have made a qualitative leap forward, reaching sufficient sizes and dynamic range to determine the characteristic scales of this process. This allowed making the first realistic predictions for a variety of observational signatures. We discuss recent results from large-scale radiative transfer and structure formation simulations on the observability of high-redshift Ly-$\alpha$ sources. We also briefly discuss the dependence of the characteristic scales and topology of the ionized and neutral patches on the reionization parameters. | The observations of high-redshift QSO's \citep{2001AJ....122.2833F,2001AJ....122.2850B} and large-scale CMB polarization \citep{2007ApJS..170..377S} indicate that the intergalactic medium has been completely ionized by redshift $z\sim6$ through an extended process. The most probable cause was the ionizing radiation of the First Stars and QSO's. Currently these are the two main direct observational constraints on this epoch. This scarcity of observational data is set to change dramatically in the next few years, however. A number of large observational projects are currently under way, e.g. observations at the redshifted 21-cm line of hydrogen \citep[e.g.][]{1997ApJ...475..429M,2000ApJ...528..597T, 2002ApJ...572L.123I,2006MNRAS.372..679M,2006PhR...433..181F}, detection of small-scale CMB anisotropies due to the kinetic Sunyaev-Zel'dovich (kSZ) effect \citep[e.g.][]{2000ApJ...529...12H,2003ApJ...598..756S,kSZ}, and surveys of high-redshift Ly-$\alpha$ emitters and studies of the IGM absorption \citep[e.g.][]{2003AJ....125.1006R,2004ApJ...604L..13S, 2006NewAR..50...94B}. The planning and success of these experiments relies critically upon understanding the large-scale geometry of reionization, i.e. the size- and spatial distribution of the ionized and neutral patches. This is best derived by large-scale simulations, although a number of semi-analytical models exist as well \citep[e.g.][]{2004ApJ...613....1F}. Recently we presented the first large-scale, high-resolution radiative transfer simulations of cosmic reionization \citep{2006MNRAS.369.1625I,% 2007MNRAS.376..534I} and applied those to derive a range of reionization observables \citep{2006MNRAS.372..679M,kSZ,pol21,cmbpol,wmap3,2007arXiv0708.3846I}. Here we summarize recent results on the characteristic scales and topology of reionization and implications of our simulations for the observability of high-redshift Ly-$\alpha$ sources. | 7 | 10 | 0710.2451 |
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0710 | 0710.4936_arXiv.txt | We interpret the recent gravitational lensing observations of Jee et al. \cite{Jee} as first evidence for a {\it caustic} ring of dark matter in a galaxy cluster. A caustic ring unavoidably forms when a cold collisionless flow falls with net overall rotation in and out of a gravitational potential well. Evidence for caustic rings of dark matter was previously found in the Milky Way and other isolated spiral galaxies. We argue that galaxy clusters have at least one and possibly two or three caustic rings. We calculate the column density profile of a caustic ring in a cluster and show that it is consistent with the observations of Jee et al. | Using strong and weak gravitational lensing methods, Jee et al. \cite{Jee} constructed a column density map of the central region of the galaxy cluster Cl 0024+1654. The map shows a ring of dark matter of radius $\simeq$ 400 kpc, width $\sim$ 150 kpc and maximum column density $\simeq 58 {M_\odot \over {\rm pc}^2}$. Remarkably, the overdensity in dark matter is not accompanied by an analogous structure in x-ray emitting gas or luminous matter. In this regard, Jee et al. discovered a new instance where dark and ordinary matter have dramatically different spatial distributions. The well-known observations of the ``Bullet Cluster" 1E0657-56 \cite{bul} provided an earlier example. Jee et al. interpret the dark matter ring in Cl 0024+1654 as the product of a near head-on collision, along the line of sight, of two subclusters. They performed a simulation of the response of the dark matter particles to the time-varying gravitational field and found that, after the collision has occurred, the dark matter particles move outward and form shell-like structures which appear as a ring when projected along the collision axis \cite{Jee}. The interpretation of Jee et al. fits with independent lines of evidence that Cl 0024+1654, an apparently relaxed cluster, is a collision of two subclusters. However, it is shown in ref. \cite{ZuH} that the observed dark matter ring is reproduced only for highly fine-tuned, and hence unlikely, initial velocity distributions. The purpose of our paper is to propose an alternative interpretation, to wit that Jee et al. have observed a caustic ring formed by the in and out flow of dark matter particles falling onto the cluster for the first time. We show that the observed ring is explained assuming only that the dark matter falling onto the cluster has net overall rotation, with angular momentum vector close to the line of sight, and velocity dispersion less than 60 km/s. When cold collisionless dark matter falls from all directions into a smooth gravitational potential well, the phase space distribution of the dark matter particles is characterized everywhere by a set of discrete flows \cite{Ips}. The flows form outer and inner caustics. The outer caustics are formed by outflows where they turn around before falling back in. Each outer caustic is a fold catastrophe ($A_2$) located on a topological sphere surrounding the potential well. The inner caustics \cite{crdm,sing} are formed near where the particles with the most angular momentum in a given inflow reach their closest approach to the center before going back out. The catastrophe structure of the inner caustics \cite{inner} depends on the angular momentum distribution of the infalling particles. If that angular momentum distribution is characterized by net overall rotation, the inner caustics are rings (closed tubes) whose cross-section is a section of the elliptic umbilic catastrophe ($D_{-4}$) \cite{sing}. These statements are valid independently of any assumptions of symmetry, self-similarity, or anything else. It is argued in ref.~\cite{rob} that discrete flows and caustics are a generic and robust property of {\it galactic} halos if the dark matter is collisionless and cold. The radii of the outer caustic spheres are predicted by the self-similar infall model \cite{FG,B} of halo formation. The radii of the inner caustic rings are predicted \cite{crdm}, in terms of a single parameter $j_{\rm max}$, after the model is generalized \cite{STW} to allow angular momentum for the infalling particles. Evidence for inner caustic rings distributed according to the predictions of the self-similar infall model has been found in the Milky Way \cite{crdm,milk,mon} and in other isolated spiral galaxies \cite{crdm,Kinn}. The resolution of most present numerical simulations is inadequate to see discrete flows and caustics. However such features are seen in dedicated simulations which increase the number of particles in the relevant regions of phase space \cite{Stiff,simca}. They should also become apparent in fully general simulations of structure formation through the use of special techniques \cite{White}. | We conclude that the dark matter ring observed by Jee et al. has properties consistent with a caustic ring of dark matter. The column density agrees both in shape and overall amplitude. At present the data are too imprecise to infer with confidence the nature of the observed ring. However, our proposal makes a distinct prediction for the column density profile across the ring, as illustrated in Fig.1. This may be tested by future observations. In this regard, let us emphasize that the ${1 \over x} I_a({2(x-a) \over s})$ profile applies to each azimuth. If the caustic ring interpretation is confirmed, an important corollary is that dark matter falls onto Cl 0024+1654 with net overall rotation. ~~ We thank Igor Tkachev for making available his numerical codes to solve the equations of the self-similar infall model. We thank Leanne Duffy for useful discussions. This work was supported in part by the U.S. Department of Energy under contract DE-FG02-97ER41029. P.S. gratefully acknowledges the hospitality of the Aspen Center of Physics while working on this project. | 7 | 10 | 0710.4936 |
0710 | 0710.3727_arXiv.txt | The energetics and emission mechanism of GRBs are not well understood. Here we demonstrate that the instantaneous peak flux or equivalent isotropic peak luminosity, $L_{iso}$ ergs s$^{-1}$, rather than the integrated fluence or equivalent isotropic energy, $E_{iso}$ ergs, underpins the known high-energy correlations. Using new spectral/temporal parameters calculated for 101 bursts with redshifts from {\em BATSE}, {\em BeppoSAX}, {\em HETE-II} and {\em Swift} we describe a parameter space which characterises the apparently diverse properties of the prompt emission. We show that a source frame characteristic-photon-energy/peak luminosity ratio, $K_{z}$, can be constructed which is constant within a factor of 2 for all bursts whatever their duration, spectrum, luminosity and the instrumentation used to detect them. The new parameterization embodies the Amati relation but indicates that some correlation between $E_{peak}$ and $E_{iso}$ follows as a direct mathematical inference from the Band function and that a simple transformation of $E_{iso}$ to $L_{iso}$ yields a universal high energy correlation for GRBs. The existence of $K_{z}$ indicates that the mechanism responsible for the prompt emission from all GRBs is probably predominantly thermal. | The energetics of the central engine which powers the explosion responsible for a GRB are both intriguing and fundamental to our understanding of these cosmic events. The isotropic energy outflow at source, estimated using the integrated gamma-ray fluence, is enormous, up to $E_{iso}\sim10^{54}$ ergs, and even if the outflow is collimated in jets the total energy involved is still huge, $E_{\gamma}\sim10^{51}$ ergs. The possibility that the explosion taps a standard energy resevoir has been pursued by many authors following the initial suggestion from Frail et al. (2001). If this total energy available were, indeed, roughly constant (or predictable through other means) and we could reliably estimate the collimation, then GRBs could be used as a cosmological probe to very high redshifts, Bloom et al. (2003), Ghirlanda et al. (2004). Early on it was noted that, based on analysis of {\em BATSE} data, there was a correlation between $E_{p}$, the peak of $E.F(E)$ where $F(E)$ ergs cm$^{-2}$ keV$^{-1}$ is the observed spectrum, and the fluence (Mallozzi et al. 1995, Lloyd et al. 2000). When redshifts became available for long bursts the isotropic energy, $E_{iso}$, could be estimated from the fluence and the peak energy could be transformed into the source frame, $E_{pz}$, the so-called Amati relation, a correlation between $E_{iso}$ and $E_{pz}$ in the sense that more energetic bursts have a higher $E_{pz}$, was discovered using data from {\em BeppoSAX}, (Amati et al. 2002). This correlation has subsequently been confirmed and extended although there remain many significant outliers, including all short bursts. The physical origin of the correlation may be associated with the emission mechanisms operating in the fireball but the theoretical details are far from settled (see the discussion by Amati (2006) and references therein). More recently a tighter correlation between $E_{iso}$, $E_{pz}$ and the jet break time, $t_{break}$, measured in the optical afterglow has been reported (Ghirlanda et al. 2004). This is explained in terms of a modification to the Amati relation in which $E_{iso}$ is corrected to a true collimated energy, $E_{\gamma}$, using an estimate of the collimation angle derived from $t_{break}$. The details of the collimation correction depend on the density and density profile of the circumburst medium, Nava et al. (2006) and references therein. Multivariable regression analysis was performed by Liang \& Zhang (2005) to derive a model-independent relationship, $E_{iso}\propto E_{pz}^{1.94}t_{zbreak}^{-1.24}$, indicating that the rest-frame break time of the optical afterglow, $t_{zbreak}$ was indeed correlated with the prompt emission parameters. Other studies have concentrated on the properties of the isotropic peak (maximum) luminosity, $L_{iso}$ ergs s$^{-1}$, measured over some short time scale $\approx1$ s, rather than the time integrated isotropic energy, $E_{iso}$. Yonetoku et al. (2004) noted a correlation between $L_{iso}$ and $E_{pz}$ for 16 GRBs with firm redshifts. A correlation between $L_{iso}$ and the spectral lag was first identified by Norris et al. (2000) and explained in terms of the evolution of $E_{peak}$ with time. The shocked material responsible for the gamma-ray emission is expected to cool at a rate proportional to the gamma-ray luminosity and it has been suggested that $E_{peak}$ traces the cooling (Schaefer 2004). A similar correlation between $L_{iso}$ and the variability of the GRB ($V$) was described by Reichart et al. (2001). The origin of the $L_{iso}-V$ relation is likely to be related to the physics of the relativistic shocks and the bulk Lorentz factor of the outflow. It could be that high $\Gamma_{outflow}$ results in high $L_{iso}$ and $V$ while lower luminosity and variability are expected if $\Gamma_{outflow}$ is low (see, for example, M\'{e}sz\'{a}ros et al. 2002). A rather bizzare correlation involving $L_{iso}$, $E_{pz}$ and variability was found by Firmani et al. (2006). They employed the ``high signal'' time, $T_{45}$, as formulated by Reichart et al. (2001) in their study of variability, and showed that $L_{iso}\propto E_{pz}^{1.62}T_{45}^{-0.49}$ for 19 GRBs with a spread much narrower than that of the Amati relation. There is currently no explanation for such a correlation although it may be connected with the spectral lag and variability correlations and the Amati relation. The correlation between $E_{iso}$ and $E_{pz}$ supplemented by additional empirical information can be used in pseudo redshift indicators, for example Atteia (2003), Pelangeon \& Atteia (2006), but the intrinsic spread in the correlation and uncertainty about the underlying physical interpretation introduce errors, typically of a factor $\sim2$. It may be possible to reduce the errors by simultaneous application of several independent luminosity/energy correlations, and extension of the Hubble Diagram to high redshifts using GRBs has been attempted, see for example Schaefer (2007). However, it is not clear that the correlations briefly described above are truly independent and there may be some underlying principle or mechanism which connects them all together. Recently, and more controversially, Butler et al. (2007) have raised serious doubts about the validity of these correlations suggesting that it is likely that they are introduced by observational/instrumental bias and have nothing to do with the physical properties of the GRBs and hence they conclude that GRBs are probably useless as cosmological probes. Here we take a new look at the source frame spectral and temporal properties of a large number of GRBs for which we have redshifts in order to try and understand what really correlates with what and whether or not this can provide useful intrinsic information about the GRBs and what drives them. In this analysis we include the short-duration GRBs which may share a similar emission mechanism with long bursts despite probably having different progenitors. | The equivalent isotropic energy, $E_{iso}$ ergs, of a GRB can be expressed as the product of two source frame terms, a characteristic photon energy, $E_{wz}$ keV, calculated from the shape of the spectrum across the range 1-10000 keV and the energy density at the peak of the $E.F_{z}(E)$ spectrum, $Q_{pz}$ ergs keV$^{-1}$. The correlation trend between $E_{wz}$ and $Q_{pz}$ gives rise to the Amati relation. By stacking the samples of a GRB light curve into descending order we can construct a rate profile. The functional form of such rate profiles is common to the vast majority of bursts. Fitting the profile gives us a luminosity time, $T_{Lz}$ s, a measure of the burst duration which can be used to convert the energy density at the peak to a luminosity density at peak, $Q_{pz}/T_{Lz}$ ergs keV$^{-1}$ s$^{-1}$. We can calculate the peak equivalent isotropic luminosity as a product $L_{iso}=E_{wz}Q_{pz}/T_{Lz}=E_{iso}/T_{Lz}$ ergs s$^{-1}$. $E_{wz}$ is a characteristic photon energy or a measure of the colour or hardness of the burst and $Q_{pz}/T_{Lz}$ is a measure of the instantaneous peak brightness. We have gathered and analysed sufficient spectral and temporal data from 101 bursts to produce the relation between $E_{wz}$ vs. $Q_{pz}/T_{Lz}$ and $E_{wz}$ vs. $L_{iso}$, shown in Figure \ref{fig11}, which constitutes the closest thing we have to an intrinsic colour-magnitude diagram for the peak emission from GRBs, $E_{wz}\propto L_{iso}^{0.25}$. All bursts are clustered such that we can construct a intrinsic colour-magnitude quasi constant $K_{z}$, which is a function of the source frame characteristic photon energy/peak luminosity ratio given by Equation \ref{eq15}. The range of equivalent isotropic energy that drives the expanding fireball is very large, 6 orders of magnitude (Figure \ref{fig3}), but the instantaneous hardness/brightness of the peak emission covers a very small intrinsic dynamic range, $\approx4$. The existence and form of $K_{z}$ indicates that the physical mechanism for the Gamma-ray production at the photosphere of the fireball is common to all bursts and is probably thermal although many other possibilities are not ruled out. If the prompt spectra are dominated by thermal photons the scatter in $K_{z}$ may be attributed to variations in the size and/or Lorentz factor of the fireball. XRFs have low $\Gamma_{0}$ and/or large radii. Short bursts have high $\Gamma_{0}$ and/or small radii. The relation between $T_{Lz}$ vs. $Q_{pz}$ clearly separates short from long, but both classes have the same instantaneous peak hardness/brightness. | 7 | 10 | 0710.3727 |
0710 | 0710.5667_arXiv.txt | We present a review of the standard paradigm for giant planet formation, the core accretion theory. After an overview of the basic concepts of this model, results of the original implementation are discussed. Then, recent improvements and extensions, like the inclusion of planetary migration and the resulting effects are discussed. It is shown that these improvement solve the ``timescale problem''. Finally, it is shown that by means of generating synthetic populations of (extrasolar) planets, core accretion models are able to reproduce in a statistically significant way the actually observed planetary population. | Our current understanding of planet formation is based on several centuries of observations of the planets of our own Solar System, 12 years of extrasolar planets detection, and several decades of observations of young stellar systems. These studies have let to the general concept that after the collapse of a dense gas cloud, a protostar surrounded by a protoplanetary disk was formed. In this disk, solids started to coagulate from fine dust and grew further by mutual collision to form planetesimals (provided the bottleneck by bodies roughly one meter in size can be overcome), then protoplanets, and ultimately the actual planets. Some of the protoplanets managed to accrete a massive gaseous envelope onto their core. This is the very rough outline of the core accretion model. | The core accretion paradigm explains in an unified way the formation of giant and terrestrial planets, so that there is no need for a special mechanism for giant planets. Since the first core accretion models, significant improvement and extensions were made. Such improved and extended core accretion models can form giant planets well within observed disk lifetimes, so that there is no need for a faster formation mechanism. They have also reached a degree of maturity that allows quantitative tests with observations, of both the giant planets of our own Solar System, and, by means of population synthesis, of the extrasolar planet population. In the latter case, the whole population of detected planets can be used to constrain the models, which excludes model fine tuning for a specific case, and fully exploits the observational investment. As shown by statistical tests, extended core accretion models can reproduce many observed properties and correlations in the extrasolar planet population in a quantitative significant way with one synthetic population at one time. This means that accretion models can now be used to predict future observations, so that theory can feed back on the design of future instruments | 7 | 10 | 0710.5667 |
0710 | 0710.2047_arXiv.txt | We review current state of neutron star cooling theory and discuss the prospects to constrain the equation of state, neutrino emission and superfluid properties of neutron star cores by comparing the cooling theory with observations of thermal radiation from isolated neutron stars. | The equation of state (EOS) of superdense matter in neutron star cores is still a mystery. It is not clear if it is soft, moderate or stiff; if the matter contains nucleons/hyperons, or exotic components. In the absence of good practical theory of supranuclear matter the problem cannot be solved on purely theoretical basis, but it can be solved by comparing theoretical models with observations of neutron stars. The attempts to solve this long-standing problem by different methods are numerous (e.g., Refs.\ \cite{nsb1,lp07}). Here we discuss current results obtained from studies of cooling isolated neutron stars. The first papers on neutron star cooling appeared with an advent of X-ray astronomy, before the discovery of neutron stars. Their authors tried to prove that neutron stars cool not too fast and can be discovered as sources of thermal surface X-ray radiation. The first estimates of thermal emission from cooling neutron stars were most probably done by Stabler \cite{stabler60} in 1960. Four years later Chiu \cite{chiu64} made similar estimates and analyzed the possibility to discover neutron stars from their thermal emission. First, simplified calculations of neutron star cooling were done in 1964 and 1965 \cite{morton64, cs64, bw65b}. The foundation of the strict cooling theory was laid in 1966 by Tsuruta and Cameron \cite{tc66}, one year before the discovery of pulsars. We review the current state of the cooling theory. More details can be found in recent review papers \cite{yp04,pgw06}. | The theory of cooling neutron stars of ages $10^2-10^6$ yr mostly tests the neutrino emission properties of the neutron star core. Its main results are as follows. (1) Neutrino emission in the outer core (i.e., in the core of a low-mass star) is a factor of 30--100 lower than the modified Urca emission in a nonsuperfluid star. (2) Neutrino emission in the inner core (of a massive star) is at least a factor of 30--100 higher than the modified Urca emission. It can be enhanced by direct Urca process in nucleon/hyperon inner core or by the presence of pion or kaon condensate, or quark matter. (3) The scenario with open direct Urca process predicts the existence (Fig.\ \ref{exotica}) of massive isolated neutron stars which are much colder than those observed now. In the scenario with pion condensate, the massive stars should be warmer (than those with open direct Urca) but colder than the observed ones. In the scenario with kaon condensate the massive stars should be even warmer but slightly colder than the observed sources. A discovery of cold cooling neutron stars would be crucial to constrain the level of enhanced neutrino emission in the inner core. (4) Observations of cooling neutron stars can be analyzed together with observations of SXTs in quiescent states. The data on SXTs indicate the existence of very cold neutron stars (first of all, SAX J1808.4--3658) which cool via direct Urca process, but the data and interpretation require additional confirmation. (5) A transition from slow neutrino emission in the outer core to enhanced emission in the inner core has to be smooth. Current observations of cooling neutron stars and SXTs do not constrain the parameters of this transition. A firm measurement of masses of cooling or accreting stars would help to impose such constraints. (6) New observations and reliable practical theories of dense matter are vitally important to tune the cooling theory as an instrument for exploring physical properties of neutron star interiors and neutron star parameters. The tuning will imply a careful analysis of many cooling regulators. \begin{theacknowledgments} This work was partially supported by the Russian Foundation for Basic Research (grants 05-02-16245 and 05-02-22003), by FASI-Rosnauka (grant NSh 9879.2006.2), and by the Joint Institute for Nuclear Astrophysics (grant NSF PHY 0216783). \end{theacknowledgments} | 7 | 10 | 0710.2047 |
0710 | 0710.5498_arXiv.txt | We have observed nearly 200~FGK stars at 24~and 70~microns with the {\em Spitzer} Space Telescope. We identify excess infrared emission, including a number of cases where the observed flux is more than 10~times brighter than the predicted photospheric flux, and interpret these signatures as evidence of debris disks in those systems. We combine this sample of FGK stars with similar published results to produce a sample of more than 350~main sequence AFGKM stars. The incidence of debris disks is 4.2$^{+2.0}_{-1.1}$\% at 24~microns for a sample of 213~Sun-like (FG) stars and 16.4$^{+2.8}_{-2.9}$\% at 70~microns for 225~Sun-like (FG) stars. We find that the excess rates for A, F, G, and K stars are statistically indistinguishable, but with a suggestion of decreasing excess rate toward the later spectral types; this may be an age effect. The lack of strong trend among FGK stars of comparable ages is surprising, given the factor of~50 change in stellar luminosity across this spectral range. We also find % that the incidence of debris disks declines very slowly beyond ages of 1~billion years. | Planetary system formation must be studied through various indirect means due to the relative faintness of distant planets and the long timescales over which planetary system evolution takes place. One powerful technique has advanced substantially in the era of the {\em Spitzer} Space Telescope: investigations of dusty debris disks around mature, main sequence stars. Debris disks arise from populations of planetesimals that remain from the era of planet formation; the analogs in our Solar System are the asteroid belt and the Kuiper Belt. Any system that possesses a debris disk necessarily has progressed toward forming a planetary system to some degree. The many small bodies that inhabit a debris disk can, on occasion, collide, producing a shower of fragments that grind each other down to dust particles. These dust grains can be heated by the central star to temperatures $\sim$100~K, where they can be detected at wavelengths of 10--100~microns. Data from the IRAS and ISO satellites were used to identify and characterize debris disks \citep[e.g.,][]{aumann84,decin00,spangler01,habing01,decin}. The Multiband Imaging Photometer for {\em Spitzer} (MIPS; \citet{mips}) offers substantially improved sensitivities at 24~and 70~microns and therefore can be used to advance the study of debris disks and measure the fraction of stars that possess colliding swarms of remnant planetesimals. A number of these surveys have been carried out using MIPS. \citet{astars} and \citet{astars2} observed hundreds of A~stars and found that the number of A~stars showing thermal infrared excess suggestive of collisionally-produced dust decreases as a function of stellar age, from 50\% or more at ages just past the gas dissipation age of 10~Myr to 30\% at 500~Myr. The excess rates are lower for older and lower mass stars. \citet{bryden} found that the excess rate is around 15\% for stars of mass and age similar to our Sun (for a sample of 69~stars). \citet{gautier} did not detect any 24~or 70~micron excesses suggestive of debris disks in a sample of 60~field (old) M~stars (with only 13~strongly detected at 70~microns). Finally, multiplicity appears to play a significant role in modulating debris disks, as \citet{binaries} found that the debris disk rate for A~and F~binaries was higher than that for single stars of similar spectral types. We present here a survey for excesses around almost 200~F, G, and K~stars, with ages and masses similar to that of our Sun. (Some of these data were published in \citet{bryden}.) We characterize the excesses that we find and present a few systems of particular interest. We create a larger sample of FGK~stars by adding 75~stars from a similar survey, and derive excess rates as a function of spectral type and as a function of age. Finally, we discuss the implications of our results for planetary system formation. | \subsection{Metallicity and excesses} Four stars in our FGK sample do not have published metallicities. For the 189~stars with known metallicities in Table~\ref{targetinfo}, the mean metallicity is -0.08$\pm$0.22, with a median of~-0.06. The mean metallicity of stars with excesses is -0.11$\pm$0.19 (median is~-0.09). The mean metallicity of stars with no excesses is -0.08$\pm$0.22 (median is~-0.05). There is no difference between the metallicity of the population of stars with excesses and the population of stars with no excesses. This lack of correlation has been discussed previously (e.g., B06, \citet{chastpf}). \subsection{Excess rates across spectral type \label{spectral}} Many parameters affecting infrared excesses change with spectral type, including the importance of grain loss mechanisms (winds, Poynting-Robertson drag, photon pressure); stellar luminosity; and the locations of key temperatures (e.g., the ice line) in the systems. Significant surveys for debris disks across spectral types A--M~have now been published, and we use those data to look for systematic trends across spectral types. It is well known that excess rates decrease with stellar age \citep{habing01,spangler01,astars,astars2,siegler}. This dependence must be avoided in testing for changes with spectral type. We take the oldest ($\ge$600~Myr) A~stars from the \citet{astars2} sample as our representative sample from that spectral type. For our F, G, and K~samples we take the union of the data presented here and the data in \citet{chastpf}. The targets and data reduction presented in \citet{chastpf} are quite similar to the selections and techniques we have employed here, which allows us to merge the two samples relatively seamlessly. We calculate the excess ratios for these F, G, and K samples and take the M~stars excess rates from \citet{gautier}. This compilation is presented in Table~\ref{excesssum} and Figure~\ref{spexcess}. Within the error bars, the excess rates for the A, F, G, and K~subsamples are essentially indistinguishable. However, there is a suggestion of a trend of decreasing 70~\micron\ excess rates with later spectral types (Figure~\ref{spexcess}). We note, however, that the mean age for the populations increases with later spectral types. It is possible that we are instead detecting a time-related effect, although the decay timescales identified by \citet{astars}, \citet{astars2}, and \citet{siegler} of hundreds of millions of years should long since have diminished all disks at ages of billions of years. We discuss this possibility in Section~\ref{ages}. \citet{chastpf} remarked that the excess rates for K~stars appeared to be lower than that for F and G stars, finding zero excesses among 23~stars later than K2 (and zero excesses among a larger combined sample of 61~K1--M6~stars). We include the \citet{chastpf} data in our analysis here, and find that, formally, the excess rate for K~stars is not significantly different than the excess rates for earlier (F and G) stars. As \citet{chastpf} note, and we confirm, none of the 6~K~stars with excesses in our larger sample are later than K2. Part of the motivation for assembling the F~stars program (PID~30211) was as a control sample for the binary star program presented in \citet{binaries}. In that study, 69~A3--F8 binary star systems were found, overall, to have relatively high excess rates: 9\% at 24~microns and 40\% at 70~microns. It is clear from our results here (see Figure~\ref{spexcess}) that the (single) F~stars excess rate is equal to or lower than the (single) A~stars excess rate. The binaries excess rates remain significantly high compared to the control sample of A~and F~stars. Figure~\ref{fd} shows that, in general, there is no trend of dust distance as a function of stellar effective temperature (spectral type). This may suggest that the processes that drive planetesimal formation do not depend strongly on a single critical temperature, as would be the case in the ``ice line'' model, where protoplanetary disk surface densities increase across certain temperature boundaries. However, our method for calculating dust distances may be too crude to see this effect. We note that all of the (minimum) dust distances given in Table~\ref{excesstable} would fall within the planetary realm of our Solar System ($<$30~AU). We are not observing disks that are far outside of the potential planetary realm of these systems. We stated above that there is no particular trend for either fractional luminosity or dust distance with spectral type, but there is an important caveat: no disks with large dust distances or relatively small fractional luminosities were identified among the latest stars in our sample. This dearth may simply allude to the fact that later stars are cooler. Dust at 25~AU around a K1 (5000~K) star would have a temperature around 50~K; this dust would have its peak emission near 70~microns, and cooler (more distant) dust would have its peak emission longward. MIPS 70~micron observations of such a dust population, or a cooler one, would not readily show the presence of this excess. The lack of distant disks around K stars may therefore be an observational bias. Similarly, low fractional luminosity disks would be more difficult to detect around K~stars than around earlier stars, so the lack of low fractional luminosity disks for later stars may also be due to observational bias. These observational biases may be corrected with sufficiently sensitive measurements at $\sim$100~microns (e.g., with Herschel). To further address the question of whether there is any detectable trend of excess rate as a function of spectral type using existing published data, we compared the incidence of 70~$\mu$m excesses across these 5~spectral types (A, F, G, K, M), a total sample size of more than 350~stars. To avoid uncontrolled selection effects, we confined the comparison to stars that would have been detected at 70~microns at a level of at least 2:1 on the photosphere. We then applied a number of tests. First, we computed weighted average values of R70 (observed flux over predicted flux) as in \citet{gautier}. We find a value of $\sim$5 for the A stars, but we discount this large average because of the small size of the sample (27~stars). The values for the F, G, K, and M stars are 1.16, 1.23, 1.06, and 1.025, respectively. Errors are difficult to estimate because the excesses are not normally distributed. We also computed straight averages of R70 for these same samples. Here, we calculate 4, 2.6, 1.8, 1.4, and 1.1 for the A, F, G, K, and M stars, respectively. Again, the values need to be interpreted with caution because error estimation is difficult. To circumvent the difficulties in determining errors, we binned the excesses into intervals of 0.5~in excess ratio and used the K-S test to determine if the resulting distributions were likely to have been drawn from the identical parent distribution. For each stellar type, we tested the relevant distribution against the distribution for all the stars, excluding the contribution of the spectral type in question. The result was a set of probabilities of 0.03, 0.8, 0.3, 0.3, and 0.01 that the types A, F, G, K, and M, respectively, were drawn from the same parent distribution as the other types. (For this test, a claim of a significant difference requires a probability of 0.05 or less that the samples are from the same distribution. Values of~0.3 imply that the the G~and K~stars are 1$\sigma$~different from their respective control samples.) From this suite of tests, we conclude that the incidence of excesses is different for old A stars and for M stars from that of the rest of our sample. We further deduce that there is a possibility of a difference appearing in the K~stars from their lower average excess ratio. The distributions for F and G stars appear to be indistinguishable with our data. We employ this conclusion in the creation of an FG ``supersample'' (Section~\ref{supersample}). The higher excesses indicated for the old A stars could be an age effect, since the sample is by necessity significantly younger than the later types (which we selected in general to be $>$ 1 Gyr in age). In fact, \citet{gorlova} and \citet{siegler} compare 24$\mu$m excesses from young A and solar-like stars and find that the incidence is quite similar at a given age. Age effects would presumably cause an apparent decrease of excess incidence with later type among the F, G, and K stars because F stars will evolve off the main sequence quickly enough to bias our sample toward younger objects (Section~\ref{ages}). In the end, this discussion may still be suffering from a relatively small number of K~stars sampled. An ongoing {\em Spitzer} program (PID~30490) to survey nearby stars that were not observed in other programs --- a sample that includes $\sim$400~K~stars --- should help unravel these statistics. Nonetheless, given our result, it is unlikely that future surveys will find a strong trend among F, G, and K stars of comparable ages --- a range of spectral types that spans more than a factor of~50 in stellar luminosity. This behavior is counter to our expectations and is a challenge to models of debris disk evolution. \subsection{Effects of age \label{ages}} The lack of strong dependence on spectral type lets us combine data on various types to study the evolution with stellar age. Figure~\ref{fgkage} shows individual R24 and $\chi_{70}$ determinations for the stars in our FGK~sample whose ages are known, as a function of system age. There is no correlation apparent for these individual sources, so we look to binned data in a larger combined sample for evidence of trends. We once again take the union of the data presented here and that presented in \citet{chastpf}, but this time we include only spectral types F0--K5 from the \citet{chastpf} sample (that is, we exclude the latest K~stars). This is because the data we present in this paper covers the range F0--K5, and we want the best match to our combined sample. There are 10~stars in the F0--K5 TPF/SIM subsample that have no known ages, and 10~additional stars with ages less than 1~billion years. Of these~20, 2~systems have excesses (10\%). Since this excess ratio is not significantly different from that of the overall FGK sample, and since the PID~30211 F~stars sample is controlled for age but the PID~41 sample is not specifically controlled for age, we make no attempt to correct the larger sample for age. Figure~\ref{ageexcess} shows excess rate as a function of age for this combined sample. To zeroth order, there is no trend as a function of age: a constant excess rate of $\sim$20\% adequately fits the data, being consistent at 1$\sigma$ with the 10~Gyr data point and at 1.5$\sigma$ with the 8~Gyr data point. Furthermore, the 10~Gyr bin has only 7~targets in it, and the 8~Gyr bin has only 33~targets in it (still a relatively small number). On the other hand, we note that the data shown in Figure~\ref{ageexcess} is suggestive of an excess rate that decreases with time through at least 8~Gyr. In this scenario, the 10~Gyr bin would be highly anomalous, although we note that this bin is clearly affected by small number statistics (2~excesses out of 7~stars). In other words, there may be a real trend and a real evolution of planetary systems and debris disks even on the billion-year timescale. However, this may again be a manifestation of the (potential) observational bias shown in Figure~\ref{fd}, as follows. The fraction of stars in a given age bin that are K~stars increases for the later age bins. If K~stars truly have fewer excesses (or fewer detectable excesses, according to observational biases) than other spectral types, Figures~\ref{ageexcess} and~\ref{agesfd} might indeed be showing a real decrease with increasing age, but caused not by a long-timescale evolution of planetary systems but by the increasing dominance of excess-deficient K~stars at the oldest ages. The data present in our larger sample cannot distinguish between the competing possibilities of K~stars preferentially lacking (detectable) disks, or of old stars increasingly lacking disks. It also remains to be seen whether the high excess rate for the oldest bin in Figure~\ref{ageexcess} is anything more than a small number statistics anomaly. The overall high rate of excess incidence in our samples (17\%) indicates that the 400~Myr decay timescale the drives the evolution of A~star debris disks \citep{astars2} cannot drive the evolution of debris disks in our $>$1~Gyr Sun-like sample. Instead, Sun-like stars appear to have a relatively constant incidence of 15\%--20\% that is not strongly dependent on age, but may be weakly dependent on age through a very long timescale decrease. \subsection{Debris disks around Sun-like stars \label{supersample}} Because there is no difference in excess rate between F and G stars, we can combine our sample of F0--G9~stars into a single population of ``Sun-like'' stars. To this sample of 169~stars % we add the 56~F~and G~stars from \citet{chastpf} to create an even larger sample of 225~Sun-like stars (213~stars at 24~microns). We refer to this merged sample of 213~and 225~stars as the ``Sun-like supersample.'' The excess rates for this supersample are 4.2$^{+2.0}_{-1.1}$\% at 24~microns % and 16.4$^{+2.8}_{-2.9}$\% at 70~microns % (Table~\ref{excesssum}). With this large supersample, we can now state the debris disk incidence rate for Sun-like stars with quite good confidence (small error bars). \subsection{Implications for planetary system formation} The majority of the debris disk systems that we present here have excesses at 70~microns only, suggesting temperatures $\lesssim$100~K and therefore an inner edge to the disks. These dusty debris disks are likely produced by collisions within a swarm of planetesimals akin to the asteroid belt or Kuiper Belt in our Solar System. We interpret the cool temperatures we derive for the dust in these disks as evidence of inner disk holes where the surface density of dust is much smaller than in the planetesimal ring, and potentially zero. This architecture is strongly reminiscent of our Solar System, where planets sculpt the edges of the planetesimal and dust belts. It may be that many of the systems we discuss here similarly have planets sculpting their dust distributions. In several cases, the known planets may indeed be the ones sculpting the inner edges of the disks (Tables~\ref{colortemps} and~\ref{excesstable}). Additionally, Lawler et al.\ (in prep.) have found, using IRS spectra, that the incidence of detectable levels of warm dust may be higher than the 4.2\% we find here, indicating that dust in a region analogous to our Solar System's asteroid belt may also be somewhat common. \citet{bryden07} show that the properties of disks around planet-bearing stars are unlikely to be similar to the properties of disks around stars without known planets. Briefly, the detection rates between these two populations are similar, but the planet-bearing stars are generally farther away and in more confused regions of the sky. They conclude that disks around planet-bearing stars are dustier (i.e., more massive), and probably more common, than disks around stars without known planets. Additionally, \citet{binaries} recently showed that the excess rate for binary A-F~stars is 9\% and 40\% at 24~and 70~\micron, respectively, significantly higher than our results for FGK stars and for Sun-like stars. We see in Figure~\ref{spexcess} that these excess rates are relatively high, compared to the large sample of single stars we present here. Combining these two studies, we find strong evidence that the presence of additional massive bodies (whether stellar or planetary companions) in stellar systems appears to promote higher excess rates. This effect may simply be dynamical --- more massive bodies means more stirring, more collisions, and more dust --- but it is not clear that such a process could remain effective for billions of years. Further work is necessary to explain these results. Figure~\ref{spexcess} shows that the transition from the high excess rate around A~stars to the more modest excess rate around Sun-like stars is gradual. This gradual decay is likely an age effect (consider the mean ages of the samples given in Figure~\ref{spexcess}). Our data show that excess rates for FGK stars decline slowly with stellar age, and it is not clear whether we are detecting a billion-year tail of debris disk evolution, a dependence on spectral type, and/or an observational bias. Young ($<$1~Gyr) debris disks decay on 100--400~million year timescales \citep{astars,astars2,gorlova,siegler}. The presence of excesses around 16\% of Sun-like stars at billion year ages indicates that these old debris disks must be driven by a different evolutionary process than debris disks around those younger stars. One interpretation, using our Solar System as an analogy, is that the younger systems are still active in a Late Heavy Bombardment kind of dynamical upheaval. After a billion years, such large scale processes likely have ceased in all but the most unusual systems. Debris disks are then produced from collisions within remaining planetesimal belts (e.g., our asteroid belt) that have been dynamically excited by the previous eon's dynamical stirring. However, there is no trend of fractional luminosity with age (Figure~\ref{agesfd}), although there is a lack of high fractional luminosity disks at old ages. The presence of a debris disk indicates that planetary system formation progressed at least to the planetesimal stage in a given system. There is no apparent dependence on spectral type across Sun-like stars for fractional luminosity (i.e., dust mass), dust distance, or even excess rate. This indicates that the processes that give rise to debris disk --- planetesimal formation, dynamical stirring that produces collisions --- must equally be insensitive to stellar parameters (temperature, mass, luminosity). This may argue that planetary system formation is quite robust --- able to occur in many different conditions. While none of the debris disks we observed are very similar to our own Solar System, there can be no question that the process of planetary system formation is quite common. | 7 | 10 | 0710.5498 |
0710 | 0710.4048.txt | {}{}{}{}{} % 5 {} token are mandatory % \abstract % context heading (optional) % {} % % % aims heading (mandatory) % {For Seyfert galaxies, the AGN unification model provides a simple and well established explanation of the Type~1/Type~2 dichotomy through orientation based effects. The generalization of this unification model to the higher luminosity AGNs that are the quasars remains a key question. The recent detection of Type~2 Radio-Quiet quasars seems to support such an extension. We propose to further test this scenario.} % {Type~1 but that they are seen at smaller inclination preventing us to see the broad % emission lines present in the spectrum of Type~1 AGN.} % % % methods heading (mandatory) % pola + deconvo we search ... {On the basis of a compilation of quasar host galaxy position angles consisting of previously published data and of new measurements performed using HST Archive images, we investigate the possible existence of a correlation between the linear polarization position angle and the host galaxy/extended emission position angle of quasars.} %% % % results heading (mandatory) {We find that the orientation of the rest-frame UV/blue extended emission is correlated to the direction of the quasar polarization. For Type~1 quasars, the polarization is aligned with the extended UV/blue emission while these two quantities are perpendicular in Type~2 objects. This result is independent of the quasar radio-loudness. We interpret this (anti-)alignment effect in terms of scattering in a two-component polar+equatorial model which applies to both Type~1 and Type~2 objects. Moreover the orientation of the polarization --and then of the UV/blue scattered light-- does not appear correlated to the major axis of the stellar component of the host galaxy measured from near-IR images.} % % conclusions heading (optional), leave it empty if necessary {} % These observations can be simply explained in the framework of an % extension to quasars of the unification model of AGN where the polarization is produced in two scattering regions. %% | The study of quasars shows that we can classify them among various categories. Radio-Loud quasars (RLQ) are distinguished from the Radio-Quiet quasars (RQQ) according to their radio power (Kellerman et al. \cite{ke89}), and the Type~1/Type~2 objects from the presence or absence of broad emission lines in their spectrum (Lawrence \cite{law87}). One can then wonder whether there exists a common link between all these objects, i.e. are the physical processes at the origin of all quasars the same? An interesting way to tackle this question consists in the use of the linear optical polarization. Polarimetry, in combination with spectroscopy, has led to major advances in the development of a unified scheme for AGN (see Antonucci \cite{anto93} for a review). The discovery of hidden broad emission lines in the polarized spectrum of the Type~2 Seyfert galaxy NGC1068 led to consider these objects as intrinsically identical to Type 1 Seyfert, the edge-on orientation of a dusty torus blocking the direct view of the central engine and the broad emission line region (Antonucci \& Miller \cite{anto85}). The question we investigate in this paper relates to the possible existence of a correlation between the optical polarization position angle $\theta_{Pola}$\footnote{The polarization position angle $\theta_{Pola}$ is defined as the position angle of the maximal elongation of the electric vector in the plane of polarization, measured in degrees East of North.} and the orientation of the host galaxy/extended emission $PA_{host}$\footnote{The orientation of the host galaxy $PA_{host}$ is characterized by the position angle of its major axis projected onto the plane of the sky, measured in degrees East of North.} in the case of RLQs and RQQs. The relation between quasars and their host galaxies may play a fundamental role in our understanding of the AGN phenomenon and in determining the importance of the feedback of AGN on their hosts. Such a correlation has already been studied in the case of the less powerful AGN that are the Seyfert galaxies. Thompson \& Martin (\cite{toma88}) found a tendency for Seyfert 1 to have their polarization angle aligned with the major axis of their host galaxy, an observation that they interpreted as due to dichroic extinction by aligned dust grains in the Seyfert host galaxy. In the case of the more powerful AGN that are quasars, Berriman et al. (\cite{beri90}) investigated this question. They determined by hand the $PA_{host}$ of 24 PG quasars from ground based images and computed the acute angle $\Delta \theta$ between $\theta_{Pola}$ and $PA_{host}$. They observed that while more objects seem to appear at small values ($\Delta \theta \leq 45\degr$), this effect was only marginally statistically significant. Investigating this problem with ground based data for Type~1 quasars is hampered by the inability of separating the faint host galaxy/extended emission\footnote{In the following, we use indifferently the term host galaxy or extended emission to refer to all the extended emission around the central source including stars, ionized gas or scattered light.} from the blinding light of the quasar nucleus. Because of its high angular resolution and stable Point Spread Function (PSF), the Hubble Space Telescope (HST) permits to properly remove the contribution of the powerful quasar nucleus thus allowing the investigation of the host galaxy parameters. A large number of quasar host observing programs were carried out with the HST leading to a statistically useful sample of quasar host galaxy parameters (e.g. Bahcall et al. \cite{ba97}; Dunlop et al. \cite{du03} and many more see Sects.~\ref{publidata} and \ref{newosdat}). Our aim is to investigate the $\theta_{Pola}/PA_{host}$ relation on the basis of high resolution HST quasar images. We will use either position angles given in the literature or $PA_{host}$ we measured ourselves from HST Archive observations (hereafter called ``\emph{new $PA_{host}$ data}"). The layout of this paper is as follows. In Sect.~\ref{publidata} we introduce the samples of quasars used in this study which possess a $PA_{host}$ given in the literature. In Sect.~\ref{tres}, we outline the samples with good imaging data but for which no $PA_{host}$ were published, and we summarize our data analysis process and the approach followed to model the HST Archive images and to derive the host galaxy parameters. In Sect.\ref{rapola} we briefly describe the polarization and radio data. Then in Sect.~\ref{statos} we present the statistical analysis of the sample and the results obtained. In Sect.~\ref{discu} we discuss the results and compare them to former studies. Finally, our conclusions are summarized in Sect.~\ref{conclu}. %__________________________________________________________________ | \label{conclu} Using host galaxy position angles ($PA_{host}$) determined from high resolution optical/near-IR images of quasars and data from the literature, we investigate the possible existence of a correlation between the host morphology and the polarization direction in the case of Radio-Quiet and Radio-Loud quasars. We can summarize our results as follows : \begin{enumerate} \item We find an alignment between the direction of the linear polarization and the rest-frame UV/blue major axis of the host galaxy of Type~1 quasars. In the case of Type~2 objects, it is well established that the extended UV/blue light is correlated to the observed polarization, as these blue regions are thought to be dominated by scattering. Our results suggest that such an extended UV/blue scattering region is also present in Type~1 quasars. \item We do not find such an alignment effect with the near-IR host morphology. This suggests that the morphology of the extended UV/blue emission is not related to the morphology of the stellar component of the host galaxy which dominates in the near-IR. \item We observe the same $PA_{host} - \theta_{Pola}$ behavior for either Radio-Loud or Radio-Quiet objects. This observation supports the idea that the UV/blue continuum is not entirely due to star formation processes triggered by the radio-jet. \item The observed correlation fits a unification model where the Type~1/Type~2 dichotomy is essentially determined by orientation effects assuming the two component scattering model of Smith et al. (\cite{sm04,sm05}). Indeed, depending on the viewing angle to the quasar, the polarization would predominantly arise from either the equatorial or the polar scattering region giving rise to the observed behavior. \end{enumerate} In order to strengthen the conclusions and to further investigate the correlations presented in this paper, new observations of quasars in the rest-frame UV/blue domain are needed, especially of Radio-Quiet objects for which few observations at $\lambda_{rest} \le 5000$ \AA~are available. The detection of the extended blue/UV continuum region and the measurement of its polarization would help to test our interpretation. | 7 | 10 | 0710.4048 |
0710 | 0710.5721_arXiv.txt | The equation of state for radiation is derived in a canonical formulation of the electromagnetic field. This allows one to include correction terms expected from canonical quantum gravity and to infer implications to the universe evolution in radiation dominated epochs. Corrections implied by quantum geometry can be interpreted in physically appealing ways, relating to the conformal invariance of the classical equations. | \label{sec:INTRODUCTION} In theoretical cosmology, many insights can already be gained from spatially isotropic Friedmann--Robertson--Walker models \begin{equation} \md s^2= -\md\tau^2+a(\tau)^2\left(\frac{\md r^2}{1-kr^2}+r^2(\md\vartheta^2+ \sin^2\vartheta\md\varphi^2)\right) \end{equation} with $k=0$ or $\pm 1$. The matter content in such a highly symmetric space-time can only be of the form of a perfect fluid with stress-energy tensor $T_{ab}=\rho u_au_b+P (g_{ab}+u_au_b)$ where $\rho$ is the energy density of the fluid, $P$ its pressure and $u^a$ the 4-velocity vector field of isotropic co-moving observers. Once an equation of state $P=P(\rho)$ is specified to characterize the matter ingredients, the continuity equation $\dot{\rho}+3H(\rho+P)=0$ with the Hubble parameter $H=\dot{a}/a$ allows one to determine the behavior of $\rho(a)$ in which energy density changes during the expansion or contraction of the universe. This function, in turn, enters the Friedmann equation $H^2+k/a^2=8\pi G\rho/3$ and allows one to derive solutions for $a(\tau)$. In general, one would expect the equation of state $P=P(\rho)$ to be non-linear which would make an explicit solution of the continuity and Friedmann equations difficult. It is thus quite fortunate that in many cases linear equations of state $P=w\rho$ with $w$ constant are sufficient to describe the main matter contributions encountered in cosmology at least phenomenologically. The influence of compact objects on cosmological scales is, for instance, described well by the simple dust equation of state $P(\rho)=0$. Relativistic matter, mainly electromagnetic radiation, satisfies the linear equation of state $P=\frac{1}{3}\rho$. The latter example is an exact equation describing the Maxwell field, rather than an approximation for large scale cosmology. It is thus, at first sight, rather surprising that the dynamics of electromagnetic waves in a universe can be summarized in such a simple equation of state irrespective of details of the field configuration. The result follows in the standard way from the trace-freedom of the electromagnetic stress-energy tensor and is thus related to the conformal symmetry of Maxwell's equations. That the availability of such a simple equation of state is very special for a matter field can be seen by taking the example of a scalar field $\phi$ with potential $V(\phi)$. In this case, we have an energy density $\rho=\frac{1}{2}\dot{\phi}^2+V(\phi)$ and pressure $P=\frac{1}{2}\dot{\phi}^2-V(\phi)$. Unless the scalar is free and massless, $V(\phi)=0$ for which we have a stiff fluid $P=\rho$, there is no simple relation between pressure and energy density independently of a specific solution. Any conformal symmetry such as that of elecromagnetism might be broken by quantum effects especially when quantum gravity with its new scale provided by the Planck length is taken into account. The coupling of the electromagnetic field to geometry will then change, and exact conformal symmetries can easily be violated. Accordingly, one expects corrections from quantum gravity to the radiation equation of state and corresponding effects in the universe evolution during radiation dominated epochs. In loop quantum cosmology \cite{LivRev} equations of state of matter fields are in general modified by perturbative corrections at large scales and non-perturbative ones on small scales \cite{InvScale}. This has mainly been studied so far for a scalar field for which quantum modifications can be so strong that negative pressure results independently of the chosen potential \cite{Inflation}. The main reason is the fact that the isotropic scalar field Hamiltonian $H_{\phi}=\frac{1}{2}a^{-3}p_{\phi}^2+ a^3V(\phi)$, where $p_{\phi}$ is the momentum of $\phi$, contains an inverse power of the scale factor $a$. For quantum gravity, this factor has to be quantized, too. Using the methods of \cite{QSDV}, it turns out that inverse powers receive strong loop quantum corrections at small length scales \cite{InvScale}. Accordingly, such modifications play a role for effective equations describing the universe after the big bang (or even during the quantum transition through the big bang singularity). During later stages, modifications are expected to decrease in size, but they might still be relevant due to sometimes tight constraints on evolution parameters. An extension to the usual matter ingredients of cosmology with linear equations of state is, however, difficult since the modification is based on quantizations of the fundamental field Hamiltonians. Equations of state are obtained from fundamental Hamiltonians after an analysis of the matter field equations, which can be difficult in general especially when quantum effects are taken into account. The only exception is the dust case since it implies a constant Hamiltonian (the total mass of dust) which is straightforwardly quantized without any corrections. Thus, although the dust energy density is proportional to $a^{-3}$ and metric dependent in a way which involves the inverse, it does not receive any modification since the Hamiltonian, i.e.\ total energy $a^3\rho$, is the essential object to be quantized. For radiation with $\rho\propto a^{-4}$ the expectation is not clear since the total energy does behave like an inverse power of $a$, but this follows only after an indirect analysis of the field dynamics. It is not the solution $\rho(a)\propto a^{-4}$ of the continuity equation which is quantized but the original field Hamiltonian from which the equation of state has to be derived first. One thus has to go back to the fundamental Maxwell Hamiltonian, derive energy density and pressure and see how quantum effects change the equation of state. If this is completed, one may attempt to solve the continuity equation to obtain corrections to $\rho(a)$. We will derive such corrections in this article, using the canonical quantization given by loop quantum gravity \cite{Rov,ALRev,ThomasRev}. Candidates for Hamiltonian operators of the Maxwell field have been proposed \cite{QSDV} which show several sources of correction terms. To derive corrections to the equation of state, however, we need to perform the usual calculation in a Hamiltonian formulation. Thus, we first present the canonical formulation for the free classical Maxwell field to rederive the standard result for the equation of state parameter $w$ without reference to an action or the stress-energy tensor. Appropriate modifications to the matter Hamiltonian $H_{M}$ are then made to derive possible loop quantum gravity corrections to the equation of state $w$. We will show that one case of corrections results again in a linear equation of state, albeit in a corrected way which depends on the basic discreteness scale of quantum gravity. In this case we are able to express, as in the classical case, the full field dynamics in terms of a simple modified $w$, and to solve explicitly for $\rho(a)$. Our derivation takes into account inhomogeneous field configurations and presents the first modified equation of state obtained for a realistic matter source in loop quantum gravity. | \label{sec:DISCUSSIONS} We have derived here the equation of state of the Maxwell field in a canonical form, including corrections expected from loop quantum gravity. In the canonical derivation, the reason for a linear equation of state, which is trace-freedom in the Lagrangean derivation, is the fact that the same metric dependent factor $q_{ab}/\sqrt{q}$ multiplies both terms in the Hamiltonian. The Maxwell Hamiltonian is thus simply rescaled if the metric is conformally transformed, which explains the conformal invariance of Maxwell's equations. This is special for the Maxwell field and different from, e.g., a scalar field with a non-vanishing potential. The same fact allows one to quantize the Hamiltonian in a way which affects both the electric and magnetic term in the same way, at least as far as the metric dependence is concerned. One then obtains a single correction function $\alpha=\beta$ which only corrects the metric dependence of the total scale of the Hamiltonian. In this sense, conformal invariance is preserved even after quantization. (But this would not be the case if a quantization is used which results in $\alpha\not=\beta$.) This preservation of the form of the Hamiltonian explains why we are still able to derive an equation of state independently of the specific field dynamics and that it remains linear. However, the classical value $w=\frac{1}{3}$ is corrected due to quantum effects in the space-time structure. This modification is also understandable from a Lagrangean perspective, together with basic information from the loop quantization. Employing trace freedom of the stress-energy tensor to derive the equation of state, we have to use the inverse metric in $g^{ab}T_{ab}$. But from loop quantum gravity we know that, when quantized, not all components of the inverse metric agree with inverse operators of the quantization. For the scale factor of an isotropic metric, for instance, we have $\widehat{a^{-1}}\not=\mbox{``$\hat{a}^{-1}$''}$ since the right hand side is not even defined \cite{InvScale}. While the left hand side is defined through identities such as (\ref{poissonbracketofvolume}), it satisfies $\widehat{a^{-1}}\hat{a}\not=1$ and thus shows deviations from the classical expectation $a^{-1}a=1$ on small scales which were captured here in correction functions. As derived in detail, this implies scale dependent modifications to the equation of state parameter $w_{\rm eff}$. The result can also be interpreted in more physical terms. The classical behavior $\rho(a)\propto a^{-4}$ can be understood as a combination of a dilution factor $a^{-3}$ and an additional redshift factor $a^{-1}$ for radiation in an expanding universe. As we have seen, this is corrected to $\alpha(a)a^{-4}$ where $\alpha(a)$ corrects the metric factor $q_{ab}/\sqrt{q}\sim a^{-1}\delta_{ab}$. Since this is only a single inverse power of $a$ for an isotropic solution, we can interpret the result as saying that only redshift receives corrections due to quantum effects on electromagnetic propagation. The dilution factor due to expansion is unmodified, except that the background evolution $a(t)$ itself receives corrections. This agrees with the result for dust, which is only diluted and has an unmodified equation of state even after quantization \footnote{But it disagrees with \cite{Metamorph} both for dust and radiation, where a direct quantization of energy densities exclusively for isotropic fields was attempted.}. Unlike dust, for radiation one has to refer to the inhomogeneous field and its quantum Hamiltonian to derive a reliable equation of state, as presented here. | 7 | 10 | 0710.5721 |
0710 | 0710.5517_arXiv.txt | In a large region of the supersymmetry parameter space, the annihilation cross section for neutralino dark matter is strongly dependent on the relative velocity of the incoming particles. We explore the consequences of this velocity dependence in the context of indirect detection of dark matter from the galactic center. We find that the increase in the annihilation cross section at high velocities leads to a flattening of the halo density profile near the galactic center and an enhancement of the annihilation signal. | 7 | 10 | 0710.5517 |
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0710 | 0710.2930_arXiv.txt | Hot Jupiters are new laboratories for the physics of giant planet atmospheres. Subject to unusual forcing conditions, the circulation regime on these planets may be unlike anything known in the Solar System. Characterizing the atmospheric circulation of hot Jupiters is necessary for reliable interpretation of the multifaceted data currently being collected on these planets. We discuss several fundamental concepts of atmospheric dynamics that are likely central to obtaining a solid understanding of these fascinating atmospheres. A particular effort is made to compare the various modeling approaches employed so far to address this challenging problem. | An exploding body of observations promises to unveil the meteorology of giant planets around other stars. Because of their high temperatures, short orbital periods, and likelihood of transiting their stars, the hot Jupiters are yielding their secrets most easily, and we now have constraints on radii, composition, albedo, dayside temperature structure, and even day-night temperature distributions for a variety of planets \citep[e.g.,][] {knutson-etal-2007b, cowan-etal-2007, harrington-etal-2006, harrington-etal-2007, charbonneau-etal-2002, charbonneau-etal-2005, charbonneau-etal-2007, deming-etal-2005, deming-etal-2006}. A knowledge of atmospheric dynamics is crucial for interpreting these observations. Understanding the atmospheric circulation of {\it any} planet is a difficult task, and hot Jupiters are no exception. The difficulty results from the nonlinearity of fluid motion and from the complex interaction between radiation, fluid flow, and cloud microphysics. Turbulence, convection, atmospheric waves, vortices, and jet streams can all interact in a complex manner across a range of temporal and spatial scales. Furthermore, radiative transfer and dynamics are coupled and cannot be understood in isolation. For example, the equator-to-pole temperature contrasts on terrestrial planets depend not only on the rate of latitudinal energy transport by the circulation but also on the way the atmospheric radiation field is affected by the advected temperature field. Cloud microphysics complicates the problem even more, because the large-scale radiative properties of clouds depend on the microphysics of the cloud particles (particle number density, shape, size distribution, and absorption/scattering properties). We thus have a non-linear, coupled radiation-hydrodynamics problem potentially involving interactions over 14 orders of magnitude, from the cloud-particle scale ($\sim1\,\mu$m) to the global scale ($\sim10^8\,$m for hot Jupiters). It is worth emphasizing here the important difference between the complexity of planetary atmospheres and the relative simplicity of stars, exemplified by the main sequence in the HR diagram. While specifying a few global parameters---such as mass, composition, and age---is typically sufficient to understand the key observable properties of stars, this is generally not the case for planetary atmospheres. It is possible that a significant observable diversity exists even among a group of extrasolar planets which share similar global attributes. Understanding such complexity is a challenge for extrasolar planetary science. A proven method for dealing with the complexity of planetary atmospheres is the concept of a {\it model hierarchy.} Even a hypothetical computer model that included all relevant processes and perfectly simulated an atmosphere would not, by itself, guarantee an {\it understanding} of the atmospheric behavior any more than would the observations of the real atmosphere \citep{held-2005}. This is because, in complex numerical simulations, it is often unclear how and why the simulation produces a specific behavior. To discern which processes cause which outcomes, it is important to compare models with a range of complexities (in which various processes are turned on or off) to build a hierarchical understanding. For example, east-west jet streams occur in the atmospheres of all the planets in our Solar System. The study of these jets involves a wide range of models, all of which have important lessons to teach. The most idealized are pure 2D models that investigate jet formation in the simplest possible context of horizontal, 2D turbulence interacting with planetary rotation \citep[e.g.,][]{williams-1978, yoden-yamada-1993, cho-polvani-1996b, huang-robinson-1998, sukoriansky-etal-2007}. Despite the fact that these simplified models exclude vertical structure, thermodynamics, radiation, and clouds and provide only a crude parameterization of turbulent stirring, they produce jets with several similarities to those on the planets. Next are one-layer ``shallow-water-type'' models that allow the fluid thickness to vary, which introduces buoyancy waves, alters the vortex interaction lengths, and hence changes the details of jet formation \citep{cho-polvani-1996a, cho-polvani-1996b, scott-polvani-2007, showman-2007}. 3D models with simplified forcing allow the investigation of jet vertical structure, the interaction of heat transport and jet formation, and the 3D stability of jets to various instabilities \citep{cho-etal-2001, williams-1979, williams-2003a, schneider-2006, lian-showman-2007}. Such models suggest, for example, that Jupiter and Saturn's superrotating\footnote{Superrotation is defined simply as a positive (eastward) longitudinally averaged wind velocity at the equator, so that the atmospheric gas rotates faster than the planet.} equatorial jets may require 3D dynamics. Finally are full 3D general-circulation models (GCMs) that include realistic radiative transfer, representations of clouds, and other effects necessary for detailed predictions and comparisons with observations of specific planets. The comparison of simple models with more complex models provides insights not easily obtainable from one type of model alone. This lesson applies equally to hot Jupiters: a hierarchy of models ranging from simple to complex will be necessary to build a robust understanding. Here we review our current understanding of atmospheric circulation on hot Jupiters. We first describe basic aspects of relevant theory from atmospheric dynamics; this is followed by a detailed comparison of the equations, forcing methods, and results obtained by the different groups attempting to model the atmospheric circulation on these fascinating objects. | The study of hot Jupiter atmospheres is maturing. While data of increasing quality are being collected, atmospheric models are also being refined to help build a robust, hierarchical understanding of these unusual atmospheres. Indeed, one of the main motivations for studying these atmospheres, which is also a main source of difficulty, is the unusual physical regime that characterizes them. Hence, hot Jupiters represent new laboratories for studying the complex physics of giant planet atmospheres. In this way, they offer the promise of extending the boundary of comparative planetology well beyond the solar-system planets. | 7 | 10 | 0710.2930 |
0710 | 0710.4511_arXiv.txt | We present a novel technique to phase-lock two lasers with controllable frequency difference. In our setup, one sideband of a current modulated Vertical-Cavity Surface-Emitting Laser (VCSEL) is phase locked to the master laser by injection seeding, while another sideband of the VCSEL is used to phase lock the slave laser. The slave laser is therefore locked in phase with the master laser, with a frequency difference tunable up to about 35 GHz. The sideband suppression rate of the slave laser is more than 30dB at 30 $\mu$W seed power. The heterodyne spectrum between master and slave has a linewidth of less than 1 Hz. A narrow linewidth spectrum of coherent population trapping in rubidium is achieved using such beams. | 7 | 10 | 0710.4511 |
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0710 | 0710.3041_arXiv.txt | A perturbative analysis is used to investigate the effect of rotation on the instability of a steady accretion shock (SASI) in a simple toy-model, in view of better understanding supernova explosions in which the collapsing core contains angular momentum. A cylindrical geometry is chosen for the sake of simplicity. Even when the centrifugal force is very small, rotation can have a strong effect on the non-axisymmetric modes of SASI by increasing the growth rate of the spiral modes rotating in the same direction as the steady flow. Counter-rotating spiral modes are significantly damped, while axisymmetric modes are hardly affected by rotation. The growth rates of spiral modes have a nearly linear dependence on the specific angular momentum of the flow. The fundamental one-armed spiral mode ($m=1$) is favoured for small rotation rates, whereas stronger rotation rates favour the mode $m=2$. A WKB analysis of higher harmonics indicates that the efficiency of the advective-acoustic cycles associated to spiral modes is strongly affected by rotation in the same manner as low frequency modes, whereas the purely acoustic cycles are stable. These results suggest that the linear phase of SASI in rotating core-collapse supernovae naturally selects a spiral mode rotating in the same direction of the flow, as observed in the 3D numerical simulations of Blondin \& Mezzacappa (2007). This emphasizes the need for a 3D approach of rotating core-collapse, before conclusions on the explosion mechanisms and pulsar kicks can be drawn. | } Despite extensive studies, the explosion mechanism of core-collapse supernovae is still elusive. According to the delayed explosion scenario, the shock is first stalled at a distance of a few hundred kilometers, and then revived after neutrinos diffuse out of the proto neutron star. Unfortunately, numerical simulations suggest that neutrino heating may not be efficient enough, at least in spherical symmetry \citep{lieb05}. Recent studies have shown that the spherical stalled shock is unstable against non radial perturbations with a low degree $l=1,2$ even if the flow is convectively stable. This result was demonstrated using axisymmetric numerical simulations \citep{blo03,blo06,sch06,sch08,ohn06} and linear stability analyses \citep{gal05,fog07,yam07}. Some numerical simulations have shown that this hydrodynamical instability, often called SASI, may assist the revival of the shock and trigger a successful explosion, powered either by neutrino heating \citep{mar07} or by acoustic waves \citep{bur06}. Some observed properties of young neutron stars may also be the consequences of SASI, such as their distribution of velocities \citep{sch04, sch06} or their spin \citep{blo07a,blo07b}. Until now, most studies of SASI have assumed that the unperturbed flow is purely radial and not rotating. Since the angular momentum of massive stars is likely to be large \citep{heg05}, it is desirable to understand how the properties of SASI are affected by rotation. In this Letter, the effect of rotation on the linear stage of SASI is investigated using a perturbative analysis in order to shed light on one of the surprising results observed by \cite{blo07a} in their 3D numerical simulations: the development of SASI seems to systematically favour a spiral mode rotating in the same direction as the accretion flow. As a consequence of momentum conservation, this mode diminishes and may even reverse the angular momentum acquired by the proto-neutron star from the stationary flow. Incidentally, the present linear study does not address another surprising result of \cite{blo07a}, that a spiral mode of SASI always dominate the axisymmetric mode even without rotation. Following an approach similar to \citet{fog07} (hereafter FGSJ07), we first compute the eigenfrequencies by solving accurately a boundary value problem between the shock surface and the accretor surface; in a second step, we use the same WKB method as in FGSJ07 to measure the stability of purely acoustic and advective-acoustic cycles in this region. This approach is different from \cite{lam07}, which is based on the approximate derivation of a dispersion relation. Rather than the complexity of describing the non-spherical shape of a shock deformed by rotation \citep{yam05}, we have chosen, as a first step, to solve the much simpler problem of a cylindrical accretion shock. This flow is simple enough to allow for a complete coverage of the parameter space and a physical insight of the main effects of rotation on SASI. Once characterized, these effects can be transposed into the more complex geometry of a rotating stellar core. | } \subsection{Instability mechanism} \label{mechanism} As underlined in Sect.~3, the dynamical effect of the centrifugal force on the stationary flow is modest. We anticipated in Sect.~2 that the only linear effect of angular momentum is a Doppler shift of the eigenfrequency $\omega'=\omega-m\Omega(r)$, where $\Omega(r)$ is the local rotation frequency. This leaves the axisymmetric mode $m=0$ unaffected and explains the relative insensitivity of its growth rate with respect to the rotation rate, at least for moderate angular momentum. The strong effect of rotation on the growth rate of the spiral modes can thus be traced back to this Doppler shifted frequency. What is the mechanism of the instability ? As seen in the previous section, the destabilizing role of rotation does not seem related to the presence or absence of a corotation radius, thus discarding a Papaloizou-Pringle mechanism \citep{gol85}. Two possibilities have been proposed for the mechanism of SASI without rotation; one is the advective-acoustic mechanism \citep{fog00,fog01,fog02} and the other is the purely acoustic mechanism \citep{blo06}. Up to now, there is no satisfactory direct argument for the mechanism of the modes with a long wavelength. FGSJ07 used a WKB approximation to prove that the instability of the modes with a short wavelength is due to an advective-acoustic mechanism and extrapolated this conclusion to the modes with a long wavelength, which are the most unstable. This method, recalled in Appendix~B, is based on the identification of acoustic waves and advected waves at a radius immediately below the shock surface, and the measurement of their coupling coefficients, above this radius due to the shock, and below this radius due to the flow gradients. These coupling processes are responsible for the existence of two cycles, namely a purely acoustic cycle characterized by an efficiency ${\cal R}$, and an advective-acoustic cycle characterized by an efficiency ${\cal Q}$. By using the same method, the present study does not address directly the instability mechanism of long wavelength modes. However, the WKB approximation enables us to describe, in a conclusive manner, the instability mechanism of short wavelength modes affected by rotation. First we checked that when the shock distance is increased ($r_{\rm sh}=20r_*$), the overtones are also unstable and their growth rate is an oscillatory function of the frequency similar to Fig.~7 of FGSJ07. The effect of rotation on the advective-acoustic cycle is illustrated by Fig~3, for the spiral modes $m=\pm1$ corresponding to the $10$-th overtone, as a function of the rotation rate. The cycle efficiency ${\cal Q}$ is strongly amplified by rotation if $m>0$, while strongly damped if $m<0$. The stabilization of the counter-rotating spiral coincides with a marginally stable cycle ${\cal Q}\sim1$. The calculation of the amplification factor ${\cal R}$ of perturbations during each purely acoustic cycle indicates its stability (${\cal R}<1$). Contrary to the expectation of \citet{lam07} (see next subsection), rotation clearly favours the spiral mode of the advective-acoustic cycle. This consequence of rotation established unambiguously for short wavelength perturbations is identical to the influence of rotation on the fundamental mode of SASI: we consider this a new hint that the advective-acoustic mechanism can be extrapolated to low frequencies. The detailed analysis of the consequences of the Doppler shifted frequency on the increase of the advective-acoustic efficiency ${\cal Q}$ will be presented elsewhere (\cite{yam08}, in preparation). \subsection{Comments on the Results of Laming (2007)} \label{comment} The effect of rotation on the growth rate of SASI, established in Sect. 3 in a cylindrical geometry, is qualitatively similar to the effect conjectured by \cite{lam07} (hereafter L07). Nevertheless, their investigation about the instability mechanism led them to a different interpretation of the roles of the acoustic and advective-acoustic cycles. We must point out a fundamental difference between the method of L07 and ours: by using a WKB approximation, we have carefuly defined the range of validity of our method, namely short wavelength modes. This guarantees that the advective-acoustic interpretation of the instability mechanism is physical and robust, at least in some parameter range. In contrast, the existence of a purely acoustic instability is still a conjecture because the domain of validity of the method used by L07 is ambiguous: their analytical derivation of a dispersion relation when advection is included requires to neglect terms of order $(v_r /\omega r)$ while terms of order ${\cal M}$ are retained. This approximation is not supported by the results of their Fig.~2, which indicates that $(v_{r,{\rm sh}}/\omega r_{\rm sh})$ is comparable to or larger than ${\cal M}_{\rm sh}$ for the modes $l=0$ and $l=1$. An accurate description of this acoustic mode, even in a simplified set up, would be useful to gain confidence in its possible existence. In addition to the question of the validity of the approximations used by L07, we find that our results invalidate their reasoning concerning the instability mechanism. They proposed that the advective-acoustic mechanism would be essential if $r_{\rm sh}/r_*\ge 10$, whereas a purely acoustic unstable process would be dominant for small shock radii, and they argued that rotation is a key ingredient to discriminate between the two mechanisms. When rotation is included, its effect on SASI has been attributed by L07 to a purely acoustic mechanism, despite the results of their Table 3. However, their view that rotation cannot possibly enhance the growth of the advective-acoustic cycle is clearly incorrect, at least for the short wavelength modes (our Fig.~\ref{fig4}). \subsection{Consequences of rotation on supernova explosions} \label{consequence} The perturbative study of a simple cylindrical configuration has enabled us to cover a large parameter space of shock radii and rotation rates, in order to (i) demonstrate the linear selection of non-axisymmetric modes, (ii) establish a correlation between the preferred direction of the spiral SASI and the rotation of the collapsing core, (iii) identify the advective-acoustic mechanism at work for short wavelength spiral perturbations. The fact that rotation favours a spiral mode $m=1,2$ in a cylindrical flow seems directly connected to the property observed by \cite{blo07a} in their 3D simulations including rotation. Tracing back the main influence of rotation to the local Doppler shifted frequency $\omega-m\Omega$, we may indeed expect a similar destabilization of the spiral modes with a positive value of $m$, a stabilization of the counter-rotating ones, and a comparatively weak influence on the axisymmetric modes. Even a moderate amount of angular momentum results in a shortening of the growth time of SASI through the destabilization of a non-axisymmetric mode. The promising consequences of SASI on both the explosion mechanisms and the pulsar kick could thus be considerably modified, since they were established on the basis of axisymmetric numerical simulations \citep{bur06,bur07,mar07,sch04,sch06}. Our study suggests that the effect of rotation on the linear phase of SASI can be safely neglected only for slowly rotating progenitors with a specific angular momentum $L\ll 2\pi \cdot 10^{14}$ cm$^2/$s. Although a fast growth of SASI might be helpful to an early shock revival, the dynamical effects of a spiral mode $m=1$, and even $m=2$, on the possible explosion mechanisms are not known yet. If the direction of the kick were determined by the geometry of the most unstable $l=1$ SASI mode, our perturbative approach would suggest a kick-spin misalignment. The strength of the equatorial kick may be diminished by the domination of a symmetric mode $m=2$. It is worth noting however that the relationship between the timescale of the most unstable SASI mode and the onset of explosion is not straightforward, and should be evaluated by future 3D numerical simulations. Our linear approach modestly aims at guiding our intuition for the interpretation of these simulations. \appendix | 7 | 10 | 0710.3041 |
0710 | 0710.3780_arXiv.txt | Imaging data from the Sloan Digital Sky Survey are used to characterize the population of galaxies in groups and clusters detected with the MaxBCG algorithm. We investigate the dependence of Brightest Cluster Galaxy (BCG) luminosity, and the distributions of satellite galaxy luminosity and satellite color, on cluster properties over the redshift range $0.1 \le z \le 0.3$. The size of the dataset allows us to make measurements in many bins of cluster richness, radius and redshift. We find that, within \rtwo\ of clusters with mass above $3 \times 10^{13}$\Msun, the luminosity function of both red and blue satellites is only weakly dependent on richness. We further find that the shape of the satellite luminosity function does not depend on cluster-centric distance for magnitudes brighter than \Mi\ $= -19$. However, the mix of faint red and blue galaxies changes dramatically. The satellite red fraction is dependent on cluster-centric distance, galaxy luminosity and cluster mass, and also increases by $\sim$5\% between redshifts 0.28 and 0.2, independent of richness. We find that BCG luminosity is tightly correlated with cluster richness, scaling as $L_{BCG} \sim M_{200}^{0.3}$, and has a Gaussian distribution at fixed richness, with $\sigma_{log L} \sim 0.17$ for massive clusters. The ratios of BCG luminosity to total cluster luminosity and characteristic satellite luminosity scale strongly with cluster richness: in richer systems, BCGs contribute a smaller fraction of the total light, but are brighter compared to typical satellites. This study demonstrates the power of cross-correlation techniques for measuring galaxy populations in purely photometric data. | \label{sec:intro} Clusters of galaxies are important systems for studying both galaxy evolution and cosmology. Used as laboratories with well-defined environments, these massive objects are a tool for investigating processes that influence galaxies' physical characteristics. Used as tracers of the underlying mass distribution, they are a tool for investigating the evolution of structure and the nature of dark energy in the Universe. These objectives are closely linked: cosmological studies require accurate knowledge of cluster selection and redshift- and mass-observable relations, and these facts are directly related the evolutionary properties of the cluster galaxy population. In the context of galaxy evolution, the high-density environment of galaxy clusters is a particularly interesting place to examine the galaxy population. Several studies have suggested substantial galaxy transformation in such environments. Galaxy morphology, star-formation rate, and luminosity have long been known to depend on cluster properties and to depart significantly from the cosmological average \citep[\eg,][]{Hubble26, Abell62, Oemler74, Dressler80}. Historically, the cluster galaxy content has been quantified by the luminosity function (LF) and type fraction (such as the late-type or blue fraction). Measurement of these quantities as a function of cluster mass, redshift, and distance from the cluster center provides insight into the underlying physical mechanisms responsible for these trends. While the LF is primarily a decreasing function of luminosity, galaxies are bimodal in color and spectral type, with red, early-type galaxies displaying little ongoing star formation, and blue, late-type galaxies exhibiting signs of recent star formation \citep[\eg,][]{Strateva01, Baldry04, Bell04, Menanteau06, Blanton05b}. This bimodality was in place by $z\sim1$ at the latest and may provide a signpost of galaxy transformation \citep[\eg,][]{Faber07}. Although this bimodality persists in all environments, the fraction of galaxies in each class (a.k.a. the red and blue fractions) changes systematically with local density; this trend is the so-called morphology-density relationship \citep{Oemler74,Dressler80,Dressler97,Smith05}. In recent large galaxy surveys this work has been extended to show that a wide range of galaxy properties, including morphology, star-formation rate, and color, depend on local density \citep{Gomez03,Balogh04a,Balogh04b,Hogg04,Kauffmann04,Tovmassian04,Blanton05b,Christlein05,Croton05,Rojas05,Cooper06,Mandelbaum06,Cooper07}. The advent of large galaxy surveys has made it possible to place observational constraints on both the type fraction and the LF as a function of cluster mass (the conditional luminosity function, CLF), and to further investigate how these quantities depend on other variables. In the Sloan Digital Sky Survey \citep*[][SDSS]{York00}, \citet{Goto02} and \citet{Hansen05} examined the LF as a function of cluster richness and cluster-centric distance for systems found in the SDSS Early Data Release, and \citet{Weinmann06b} measured the CLF measured from a group catalog derived from the spectroscopic sample of SDSS DR2. Using the 2dF Galaxy Redshift Survey, \citet{depropris03} and \citet{Robotham06} compared the LF in high- and low-mass systems; a similar study was performed with a sample of 93 X-ray selected clusters \citep{LMS04}. The dependence of the LF on galaxy color has been recently investigated \citep{PopessoLF} as has the galaxy type fraction \citep{Goto03,depropris04} for clusters in these large surveys. The type fraction of galaxies in clusters depends on both cluster richness and redshift: the fraction of star-forming galaxies at fixed local density is larger at higher redshift, an effect known as the Butcher--Oemler effect \citep{BO78,ButcherOemler}. This effect is now well-documented, if not entirely well-explained, in clusters over a wide range of masses by studies of the blue fraction, or its converse, the red fraction \citep{Rakos95,Margoniner00,Ellingson01,KodamaBower01,Margoniner01,depropris04,Martinez06,Gerke07}. Other indicators of galaxy state, including galaxy morphology and emission line strength, also show a trend with redshift \citep{AS93,ODB97,Balogh97,Couch98,vanDokkum00,Fasano00,Lubin02,Goto03,Treu03,Wilman05,Poggianti06,Desai07,vanderwel07}. However, few samples to date have had both large numbers of systems and well-understood mass proxies, so the dynamical range and mass resolution of these previous works has been somewhat limited. Another characteristic of galaxy clusters is the presence of a highly-luminous galaxy near the cluster center --- the Brightest Cluster Galaxy (BCG). In addition to being extraordinarily luminous, BCGs differ in a number of ways from other cluster members: they tend to have extended light profiles \citep{Matthews64, Tonry87, Schombert88, Gonzalez00, Gonzalez03}, larger size at fixed luminosity than other early types \citep[][and references therein]{Bernardi07} and may contain a larger fraction of dark matter than typical galaxies \citep[\eg,][]{Mandelbaum06,Anja07}. Also, while the traditional fitting function to the LF of \citet{Schechter76} provides a good fit for satellite galaxies, BCGs follow a different distribution, causing the so-called ``bright end bump'' \citep[\eg,][]{Hansen05}. This difference between the BCG and other satellites in a cluster is also manifested in the luminosity gap statistic, the difference in luminosity between the BCG and the next brightest cluster member, which may be indicative of the special accretion history of BCGs \citep{Ostriker75, Tremaine77, Loh06}. BCGs are also distinct from other galaxies with similar mass that are not at the center of cluster-sized potential wells \citep{Anja07}. Indeed, the properties of BCGs seem to be closely linked to properties of their host clusters \citep{Sandage73,Sch83,Schombert88, Edge91, Brough02, Degrandi04,Brough05,Loh06}, including to the masses of their parent halos \citep{LinMohr04, Mandelbaum06, ZCZ07}. The outer light profile of BCGs often merges smoothly with the diffuse intra-cluster light (ICL), suggesting again a coupling between the BCG and the host halo \citep{Gonzalez05}. Generally the BCG+ICL light is closely linked with cluster mass \citep{Zibetti05,Conroy07ICL,Purcell07,Gonzalez07}. As BCGs have properties different from those of the rest of the cluster members, BCGs and satellites are often analyzed separately, and we follow this convention here. It is expected that the properties of cluster galaxies are closely tied to the merging and accretion history of their parent dark matter halos. Current models based on Cold Dark Matter (CDM) suggest that the stars in BCGs were formed in dense peaks quite early but that the BCGs were assembled in a series of galaxy merging events that continue until relatively recent times \citep[e.g.][]{AS98,Dubinski98,Gao04a,BK06,DeLucia07}. Satellite galaxies in clusters are now generally understood to be hosted by smaller dark matter halos that have merged into the parent halo. Several features of the observed cluster galaxy populations can be understood based on the assembly history of the parent halo \citep[\eg,][]{Poggianti06, Iro07}. One such characterization is the core distinction between central galaxies and satellites: the central concentrations of mass and light in massive halos continue to build up while the growth of satellite systems is halted upon accretion. BCG luminosity, and the correlations between this luminosity and both cluster mass and satellite properties, can thus provide insight into the assembly histories of clusters. The merger history of dark matter halos alone is not enough to account for the bimodality in galaxy properties and its dependence on redshift and environment. Several physical processes have been proposed to transform star-forming galaxies into the typical cluster galaxies on the red sequence. Although the relative strengths of these processes are still hotly debated, a consensus is emerging that the main transformation mechanisms are related to the mass of the host halo and whether (and for how long) the galaxy has been a satellite within a larger system. Among the suggested transformation mechanisms, some are expected to be most effective in rich clusters, such as ram-pressure stripping \citep{GunnGott72}, interaction with the cluster potential \citep{ByrdValtonen90} and high-velocity close encounters \citep[``harassment;''][]{Moore96}. However, studies of very poor systems have shown that the environmental dependence of galaxy properties is not limited to the richest objects \citep[\eg,][]{Zabludoff98,Weinmann06a,Gerke07}. Processes that can operate efficiently in low velocity dispersion systems therefore must also play a role in shaping the galaxy population: \eg, galaxy mergers \citep{TT72} and ``strangulation,'' a cutoff to gas accretion onto galaxy disks by stripping or AGN feedback \citep{Larson80,BNM00,Croton06}. It is likely that some combination of these effects is at work. Distinguishing their relative significance requires precisely quantifying cluster galaxy properties over a wide range of masses and as a function of cluster-centric distance. These data will provide information on both the assembly histories of clusters and on the physical mechanisms that trigger and quench star formation. Models of galaxy evolution in a cosmological context most readily predict galaxy properties as a function of halo mass rather than cluster observables. In order to make these comparisons, a reliable mass--observable relationship is a prerequisite. In addition, there is consensus that the luminosity function and type fraction both depend on a number of variables, complicating detailed comparison between clusters of different mass. For example, the cluster galaxy LF depends on cluster-centric distance, and the size of the bound regions of clusters scales with mass. In order to make physically meaningful comparisons between LFs of different mass clusters, an aperture scaled to the bound region is therefore preferable to a fixed metric aperture. With recent extensive surveys providing well-calibrated cluster catalogs spanning a wide range in mass, it is possible to examine in detail the dependence of the cluster galaxy population on several cluster and galaxy properties simultaneously. Large, homogeneous photometric surveys such as the SDSS provide rich data with which to characterize the cluster galaxy population. These data have been used to define large, robust, clean samples of galaxy clusters with accurate photometric redshifts. These samples are sizable enough to split on several variables allowing detailed statistical exploration of the galaxy populations in clusters. Currently, the largest sample of clusters available is the \maxbcg\ catalog from the Sloan Digital Sky Survey \citep{Koester07a}. The selection effects of the cluster-finding algorithm are well understood \citep{Koester07b, Rozo07a}, and there are a number of studies exploring the mass--richness relationship for these objects \citep{Rozo07b,Becker07,Sheldon07a, Johnston07b,Rykoff07}. In this work, we convert cluster richness to cluster mass using weak lensing measurements of the \maxbcg\ mass--richness relation \citep{Sheldon07a,Johnston07b}. The quality and quantity of these data allow for detailed measurements of a variety of cluster galaxy properties as a function of cluster mass and radius. For satellite galaxies we then measure the luminosity function of all, red, and blue satellites conditional on both mass and cluster radius, and investigate the dependence of the red fraction of satellites on cluster mass, redshift, galaxy luminosity and distance from cluster center; for BCGs we quantify the dependence on cluster mass of both the BCG luminosity and the relationship between the BCG luminosity and satellite galaxy luminosities. Although this sample of clusters extends only to $z = 0.3$, these objects provide a valuable low-redshift baseline with which higher redshift samples may be compared. In this paper we use a statistical background-subtraction technique to measure mean galaxy properties over a wide range of color and luminosity in \maxbcg\ clusters. We average the signal from many clusters binned by cluster properties, and statistically subtract the contribution from random galaxies along the line of sight. This method of cross-correlating clusters with the galaxy population provides very precise statistical measurements, and allows us to study blue and low luminosity galaxies that are indistinguishable from the background in individual clusters. We test the background-correction algorithm by running the full analysis on realistic mock catalogs, and find that we are able to robustly recover 3D cluster properties using these methods. The statistical techniques presented here for background correction of photometric data will be directly applicable to future large multi-band imaging programs, including the Dark Energy Survey \citep[DES\footnote{\tt http://www.darkenergysurvey.org},][]{DES} and the Large Synoptic Survey Telescope \citep[LSST\footnote{{\tt http://www.lsst.org}},][]{LSST} and the Panoramic Survey Telescope \& Rapid Response System \citep[Pan-STARRS\footnote{{\tt http://pan-starrs.ifa.hawaii.edu}},][]{PANSTARRS} and will be essential for leveraging these data to provide the desired insights into cosmology and galaxy evolution. The paper is organized as follows: in \S\ \ref{sec:data} we describe the SDSS and simulation data used; we present the stacking and background-correction method in \S\ \ref{sec:methods}. Our primary results are given in \S\ \ref{sec:results}: \S\ \ref{sec:sats} presents the luminosity and color characteristics of the satellite population, while \S\ \ref{sec:BCGs} discusses the BCG population. A summary and discussion of the implications of the results is given in \S\ \ref{sec:conclusion}. The notation used for cluster-related variables in previous \maxbcg\ work includes defining \Ntwo\ and $L_{200}$ as the counts and $i$-band luminosity of red-sequence galaxies within the measurement aperture of the cluster finder and with L $> 0.4L_*$. We note that as the aperture for cluster finding was determined with a previous definition of richness, it is not strictly the true value of \rtwo\ for these systems (in fact, it is larger), and only red galaxies are included in these definitions. In this work we will refer to the total excess luminosity associated with the light from galaxies of {\em all} colors above a luminosity threshold and within the measured \rtwo\ of these systems as $L_{200}$. In addition, we follow the standard convention of using $R$ to to denote projected, 2D radii and $r$ to refer to deprojected, 3D radii. Where necessary for computing distances, we assume a flat, LCDM cosmology with H$_{0} = 100h$ km s$^{-1}$ Mpc$^{-1}$, $h = 0.7$, and matter density $\Omega_m = 0.3$. | \label{sec:conclusion} In this study, we have examined the properties of cluster-associated galaxies, separating BCGs from satellites and focusing on trends in galaxy color and luminosity as a function of cluster richness, distance from cluster center, and redshift. We use clusters in the SDSS \maxbcg\ sample, the largest set of clusters identified to date. Employing photometric data alone, we apply cross-correlation background-correction techniques to characterize the cluster-associated galaxy population, within $5\times$\rtwo\ and brighter than \Mi\ $< -19$, around \nclust\ systems spanning more than two decades in mass in the redshift range $0.1 \le z \le 0.3$. Our principle results are as follows. \begin{enumerate} \item The luminosity function of satellites within \rtwo\ as a function of cluster mass for systems with mass greater than $3 \times 10^{13}$\Msun\ shows remarkable uniformity for \Mi\ $< -19$. The characteristic satellite luminosity \Lstar\ is only weakly dependent on cluster richness. \item The shape of the luminosity function of satellites brighter than \Mi\ $< -19$ does not change with cluster-centric radius. However, the color-separated luminosity functions of satellites as a function of \rad\ and of \Ntwo\ show that the mix of sub-\Lstar\ red and blue galaxies changes dramatically as a function of radius. In contrast, the relative number of red and blue galaxies at the bright end is roughly constant with radius. \item The average color of satellite galaxies is redder near cluster centers, but this trend is largely a reflection of the changing ratio of red to blue galaxies mentioned above. This effect is a very weak function of cluster mass over the range investigated here. \item The fraction of red galaxies increases with cluster mass over the full range explored, although only weakly for \Mtwo $\gae 10^{14}$\Msun. This fraction decreases with cluster-centric distance until $2\times$\rtwo, decreases with luminosity for $L<$ \Lstar, and is constant for $r>2\times$\rtwo\ and $L>$ \Lstar. \item The fraction of cluster galaxies within \rtwo\ that are red increases by $\sim$5\% during the 0.8 Gyr between redshift $z=0.28$ to $z=0.2$; this change is roughly independent of mass over the range investigated. \item The luminosity of BCGs and the ratios of BCG luminosity to total cluster luminosity and to characteristic satellite luminosity are all correlated strongly with cluster mass, and we have quantified each of these scalings over the mass range $3 \times 10^{13}$\Msun\ to $9 \times 10^{14}$\Msun. The BCG luminosity has a Gaussian distribution at fixed cluster richness, with dispersion $\sigma_{logL} \sim 0.17$ for clusters with \Mtwo\ $ > 10^{14}$\Msun. \end{enumerate} While these results are in general agreement with previous observational work, due to the volume probed we are able to investigate a wider mass range, extending the statistics to higher mass than previous samples. We are also able to split the sample into finer bins for several variables than has previously been possible. The \maxbcg\ selection function, which is well-understood for most of the richness range used here, is less well quantified for clusters with \Ntwo\ $< 10$. This low richness set of clusters may be less complete and pure, which could result in the significant difference in LF shape as compared to higher \Ntwo\ systems, or contribute to the drop-off in \fred\ for low richness systems. However, not all cluster galaxy properties change significantly at \Ntwo\ $= 10$ (\eg, \Lbcg/\Ltwo) and there is no break at a particular \Ntwo\ in either the mass--observable relationship \citep{Johnston07b} or mass-to-light ratios \citep{Sheldon07b} as measured by lensing. Interestingly, it is the quantities that most closely trace the total mass of the systems (\ie, \Lbcg, \Ltwo and the lensing signal) that are smoothly scaling over the full richness range, while quantities that are related to the mix of galaxies within clusters (\ie, the LF and \fred) are the ones that change more dramatically. We hypothesize that, at these low richnesses, \maxbcg\ is finding legitimate low-mass systems, but that the selection priors demanding the close proximity of only a few red galaxies result in finding only systems with the observed mix of galaxies and not necessarily {\em all} systems of this low mass. Further investigation using lower mass threshold mock catalogs is needed to understand in detail the selection of low-richness systems. Our results can shed light both on the processes that build up the galaxy population in clusters and distinguish central galaxies from satellites, as well as the processes that are responsible for the galaxy transformation from blue to red. With respect to the former, the results presented here fit well within the basic picture of galaxy formation in CDM: that galaxy properties are likely linked to the formation history of a cluster's dark matter halo and its substructures. Although detailed comparisons are beyond the scope of this work, our results qualitatively match both HOD constraints from clustering statistics as well as models based on matching the abundance of halos and subhalos to galaxies. The HOD framework provides a way to examine the galaxy population in both the observed Universe and in models of galaxy formation and evolution. Without needing to identify specific groups or clusters in the data, halo model interpretations of the statistics of luminosity-dependent galaxy clustering result in specific predictions for trends of both central and satellite galaxies as a function of halo mass that are in reasonable agreement with the findings presented here, especially for the relationship between central and satellite luminosities \citep[\eg,][]{Berlind03,Skibba07}. Models which link the properties of galaxies directly to their halos and subhalos also agree broadly with several of the results presented here \citep[\eg,][]{ Conroy06, VO07}. Detailed predictions for several cluster statistics from such a model will be presented in \citet{Iro07}. Our results on scatter in the BCG luminosity at fixed cluster richness very likely provide an upper limit on scatter in central galaxy luminosity at fixed mass, as the scatter between halo mass and cluster richness should act to increase this scatter. The fact that this scatter is already fairly small provides further support for the tight coupling between halo mass and galaxy luminosity that is the basis of these models. In addition to processes that cause physical changes to satellites resulting from interactions with the cluster gas, cluster potential, or other satellites, presumably some of the satellites are lost due to being accreted onto the BCG. In detail, this process likely results in the stellar component of the disrupted galaxy joining both the BCG and ICL \citep{Conroy07ICL}, but nonetheless should result in a BCG population closely linked to both halo mass and satellite population, as is observed. Indeed, \Lbcg/\Ltwo\ and \Lbcg/\Lstarsat\ must be intimately related to the processes responsible for BCG growth. That BCGs get brighter as a function of cluster mass faster than do typical satellies may be further evidence that BCGs are different (in their merger history) than typical satellites. Understanding the timescales and mechanisms for galaxies to transform from star-forming galaxies onto the red sequence is one of the primary current challenges for galaxy formation theories. There are several processes that can operate within clusters to shape the population of the cluster galaxies, such as ram pressure stripping, harassment and strangulation, that directly influence the galaxies' gas content and thus their subsequent star formation \citep[for a recent review, see][]{DeLucia06}. Clearly many of the processes responsible for this transformation are related either to the mass of the host halo or how long the galaxies have been satellites. The relative strengths of various effects are still rather uncertain, however. What fraction of galaxies become red while they are central galaxies? Do the processes happen only for satellites in a certain mass range, or do they happen equally for all satellite galaxies? On what timescales do these processes operate? Our measurements of how the red fraction scales with cluster mass, radius and redshift will be instrumental in answering these questions. We address a few of these issues here. Ram pressure stripping predicts that \fred\ will be inversely correlated with cluster-centric distance, and larger for both brighter galaxies and more massive halos. However, our results indicate that \fred\ is essentially independent of galaxy luminosity at fixed \Ntwo\ for galaxies brighter than \Lstar\ and that the radial trend in \fred\ is not any more pronounced in high \Ntwo\ systems. These results indicate that ram pressure stripping cannot be the dominant process at work to transform the galaxy population. The harassment scenario predicts that \fred\ will be anticorrelated with cluster mass, as is observed; however, if this mechanism were dominant, then at fixed cluster mass \fred\ would be expected to be larger for less luminous galaxies \citep{Weinmann06a}, contradicting the observed trends. Strangulation, where star formation in infalling galaxies is halted because no further gas accretion is allowed, makes several predictions that are in good agreement with our observations. For example, the model presented by \citet{Diaferio01} predicts: that the mean satellite color gets bluer as a function of cluster-centric distance until reaching a plateau at the field value around 2--3\rtwo; that the mean satellite color depends on halo mass only for $M \lae 5 \times 10^{13}$\Msun; and that the incidence of blue galaxies in clusters increases at higher redshift. However, this model also predicts that both bulge- and disk-dominated satellies will get redder toward the central regions of clusters, a trend that we do not see in our red/blue sample split (or course, our color-separated subsamples are not directly comparable to their morphology-separated samples). Note that the predictions of such a model will in detail depend on its implementation. Our observation that the satellite red fraction changes in the same way with redshift regardless of cluster richness is significant evidence that the timescale of the physics responsible for quenching is the same in systems over the full mass range examined. The difference in cosmic time between the median redshifts of our two samples is approximately the dynamical timescale, over which we observe a $\sim$5\% change in \fred, and this observation should allow for useful constraints in studies on the details of quenching mechanisms. A consensus is emerging from a variety of modeling efforts that star formation proceeds most efficiently in a mass range around $L_*$, and is less efficient in more massive halos and for satellite galaxies. Simple models based on this type of assumption can produce many of the rough trends that we have seen here; for example, that the mean color of galaxies will not change significantly as a function of halo mass for the range of masses we investigate here \citep[\eg,][]{Diaferio01}, and that the galaxy type fraction should be a weak function of halo mass for $M$ \gae $10^{13}$\Msun\ \citep{Berlind03,Zheng05,Cooray05}. Furthermore, the general concept of galaxies falling in, quenching, and fading matches well with the observed radial trends of LF and red fraction. However, most of the detailed semi-analytic modeling efforts have had some trouble matching in detail observations of the color distribution of galaxies and how it changes with environment \citep[see \eg,][]{Coil07}. Unresolved issues include the rates at which star formation is triggered or shut off, and accordingly red galaxies tend to be overproduced. To understand in detail the physical mechanisms responsible for the quenching of star formation as clusters are assembled, further work is needed to accurately reproduce the observed trends in the cluster galaxy population. The present findings set a local-universe target for modeling results, and provide some guidance for the relative importance of some of the germane effects. A substantial effort over the next decade will be devoted to large-scale, multi-band photometric surveys, including DES, Pan-STARRS, and LSST. Although the primary science driver of many of these projects is to investigate the nature of dark energy, the resulting data are likely to also provide strong constraints on the processes of galaxy evolution. The results presented here provide a low-redshift baseline against which current and future high-redshift samples may be compared. From a technical standpoint, these data are informative to next-generation cluster-finding techniques, and are useful input for creating the mock catalogs necessary for interpreting cluster surveys. Furthermore, the techniques presented here, which use photometric data alone, are directly applicable to these upcoming imaging surveys, and will thus enable detailed studies of the galaxy population at significantly higher redshifts without extensive spectroscopy. | 7 | 10 | 0710.3780 |
0710 | 0710.1099_arXiv.txt | The reliability of quiet Sun magnetic field diagnostics based on the \ion{Fe}{1} lines at 6302 \AA \, has been questioned by recent work. We present here the results of a thorough study of high-resolution multi-line observations taken with the new spectro-polarimeter SPINOR, comprising the 5250 and 6302 \AA \, spectral domains. The observations were analyzed using several inversion algorithms, including Milne-Eddington, LTE with 1 and 2 components, and MISMA codes. We find that the line-ratio technique applied to the 5250~\AA \, lines is not sufficiently reliable to provide a direct magnetic diagnostic in the presence of thermal fluctuations and variable line broadening. In general, one needs to resort to inversion algorithms, ideally with realistic magneto-hydrodynamical constrains. When this is done, the 5250~\AA \, lines do not seem to provide any significant advantage over those at 6302~\AA . In fact, our results point towards a better performance with the latter (in the presence of turbulent line broadening). In any case, for very weak flux concentrations, neither spectral region alone provides sufficient constraints to fully disentangle the intrinsic field strengths. Instead, we advocate for a combined analysis of both spectral ranges, which yields a better determination of the quiet Sun magnetic properties. Finally, we propose the use of two other \ion{Fe}{1} lines (at 4122 and 9000~\AA ) with identical line opacities that seem to work much better than the others. | \label{sec:intro} The empirical investigation of quiet Sun\footnote{In this work, we use the term ``quiet Sun'' to refer to the solar surface away from sunspots and active regions.} magnetism is a very important but extremely challenging problem. A large (probably dominant) fraction of the solar magnetic flux resides in magnetic accumulations outside active regions, forming network and inter-network patches (e.g., \citeNP{SNSA02}). It is difficult to obtain conclusive observations of these structures, mainly because of two reasons. First, the size of the magnetic concentrations is much smaller than the spatial resolution capability of modern spectro-polarimetric instrumentation. Estimates obtained with inversion codes yield typical values for the filling factor of the resolution element between $\sim$1\% and 30\% . The interpretation of the polarization signal becomes non-trivial in these conditions and one needs to make use of detailed inversion codes to infer the magnetic field in the atmosphere. Second, the observed signals are extremely weak (typically below $\sim$1\% of the average continuum intensity), demanding both high sensitivity and high resolution. Linear polarization is rarely observed in visible lines, so one is usually left with Stokes~$I$ and~$V$ alone. \citeN{S73} proposed to use the pair of \ion{Fe}{1} lines at 5247 and~5250~\AA \, which have very similar excitation potentials and oscillator strengths (and, therefore, very similar opacities) but different Land\'e factors, to determine the intrinsic field strength directly from the Stokes~$V$ line ratio. That work led to the subsequent popularization of this spectral region for further studies of unresolved solar magnetic structures. Later, the pair of \ion{Fe}{1} lines at 6302~\AA \, became the primary target of the Advanced Stokes Polarimeter (ASP, \citeNP{ELT+92}), mainly due to their lower sensitivity to temperature fluctuations. The success of the ASP has contributed largely to the currently widespread use of the 6302~\AA \, lines by the solar community. Recent advances in infrared spectro-polarimetric instrumentation now permit the routine observation of another very interesting pair of \ion{Fe}{1} lines, namely those at 15648 and 15653~\AA \, (hereafter, the 1.56~$\mu$m lines). Examples are the works of \citeN{LR99}; \citeN{KCS+03}. The large Land\' e factors of these lines, combined with their very long wavelengths, result in an extraordinary Zeeman sensitivity. Their Stokes~$V$ profiles exhibit patterns where the $\sigma$-components are completely split for fields stronger than $\sim$400~G at typical photospheric conditions. They also produce stronger linear polarization. On the downside, this spectral range is accesible to very few polarimeters. Furthermore, the 1.56~$\mu$m lines are rather weak in comparison with the above-mentioned visible lines. Unfortunately, the picture revealed by the new infrared data often differs drastically from what was being inferred from the 6302~\AA \, observations (e.g., \citeNP{LR99}; \citeNP{KCS+03}; \citeNP{SNSA02}; \citeNP{SNL04}; \citeNP{DCSAK03}), particularly in the inter-network. \citeN{SNSA03} proposed that the discrepancy in the field strengths inferred from the visible and infrared lines may be explained by magnetic inhomogeneities within the resolution element (typically 1\arcsec). If multiple field strengths coexist in the observed pixel, then the infrared lines will be more sensitive to the weaker fields of the distribution whereas the visible lines will provide information on the stronger fields (see also the discussion about polarimetric signal increase in the 1.56 $\mu$m lines with weakening fields in \citeNP{SAL00}). This conjecture has been tested recently by \citeN{DCSAK06} who modeled simultaenous observations of visible and infrared lines using unresolved magnetic inhomogeneities. A recent paper describing numerical simulations by \citeN{MGCRC06} casts some doubts on the results obtained using the 6302~\AA \, lines. Our motivation for the present work is to resolve this issue by observing simultaneously the quiet Sun at 5250 and 6302~\AA . We know that unresolved magnetic structure might result in different field determinations in the visible and the infrared, but the lines analyzed in this work are close enough in wavelengths and Zeeman sensitivities that one would expect to obtain the same results for both spectral regions. | \label{sec:conc} The ratio of Stokes~$V$ amplitudes at 5250 and 5247~\AA \, is a very good indicator of the intrinsic field strength in the absence of line broadening, e.g. due to turbulence. However, line broadening tends to smear out spectral features and reduce the Stokes~$V$ amplitudes. This reduction is not the same for both lines, depending on the profile shape. If the broadening could somehow be held constant, one would obtain a line-ratio calibration with very low scatter. However, if the broadening is allowed to fluctuate, even with amplitudes as small as 1~km~s$^{-1}$, the scatter becomes very large. Fluctuations in the thermal conditions of the atmosphere further complicate the analysis. This paper is not intended to question the historical merits of the line-ratio technique, which led researchers to learn that fields seen in the quiet Sun at low spatial resolution are mostly of kG strength with small filling factors. However, it is important to know its limitations. Otherwise, the interpretation of data such as those in Figure~\ref{fig:mapratios} could be misleading. Before this work, most of the authors were under the impression that measuring the line ratio of the 5250~\AA \, lines would always provide an accurate determination of the intrinsic field strength. With very high-resolution observations, such as those expected from the Advanced Technology Solar Telescope (ATST, \citeNP{KRK+03}) or the Hinode satellite, there is some hope that most of the turbulent velocity fields may be resolved. In that case, the turbulent broadening would be negligible and the line-ratio technique would be more robust. However, even with the highest possible spatial resolution, velocity and temperature fluctuations along the line of sight will still produce turbulent broadening. From the study presented here we conclude that, away from active region flux concentrations, it is not straightforward to measure intrinsic field strengths from either 5250 or 6302~\AA \, observations taken separately. Weak-flux internetwork observations would be even more challenging, as demonstrated recently by \citeN{MG07}. Surprisingly enough, the 6302~\AA \, pair of \ion{Fe}{1} lines is more robust than the 5250~\AA \, lines in the sense that it is indeed possible to discriminate between weak and strong field solutions if one is able to rule out a thermal stratification with temperatures that increase outwards. Even so, this is only possible when one employs an inversion code that has sufficient MHD constrains (an example is the MISMA implementation used here) to reduce the space of possible solutions. The longitudinal flux density obtained from inversions of the 6302~\AA \, lines is better determined than those obtained with 5250~\AA . This happens regardless of the inversion method employed, although using a code like LILIA provides better results than a simpler one such as MELANIE. The best fits to average network profiles correspond to strong kG fields, as one would expect. An interesting conclusion of this study is that it is possible to obtain reliable results by inverting simultaneous observations at both 5250 and 6302~\AA . Obviously this would be possible with relatively sophisticated algorithms (e.g., LTE inversions) but not with simple Milne-Eddington inversions. The combination of two other \ion{Fe}{1} lines, namely those at 4122 and 9000~\AA , seems to provide a much more robust determination of the quiet Sun magnetic fields. Unfortunately, these lines are very distant in wavelength and few spectro-polarimeters are capable of observing them simultaneously. Examples of instrument with this capability are the currently operational SPINOR and THEMIS, as well as the planned ATST and GREGOR. Depending on the evolution time scales of the structures analyzed it may be possible for some other instruments to observe the blue and red lines alternatively. | 7 | 10 | 0710.1099 |
0710 | 0710.3325_arXiv.txt | {} {In this paper we study the possibility of testing $CPT$ symmetry with Cosmic Microwave Background (CMB) measurements.} {Working with an effective lagrangian of the photon with $CPT$ violation ${\cal L} \sim p_{\mu}A_{\nu}\tilde F^{\mu\nu}$ which causes the polarization vectors of the propagating CMB photons rotated, we determine the rotation angle $\Delta\alpha$ using the BOOMERanG 2003 and the WMAP3 angular power spectra.} {In this analysis we have included the newly released $TC$ and $GC$ ($l<450$) information of WMAP3 and found $\Delta\alpha=-6.2\pm3.8$ deg at $68\%$ confidence level.} {This result increases slightly the significance for the $CPT$ violation obtained in our previous paper (Feng et al. 2006) $\Delta\alpha=-6.0 \pm 4.0$ deg (1$\sigma$). Furthermore we examine the constraint on the rotation angle from the simulated polarization data with Planck precision. Our results show that the future Planck measurement will be sensitive to $\Delta \alpha$ at the level of $0.057$ deg and able to test the $CPT$ symmetry with a higher precision.} | \label{Introduction} In the standard model of particle physics $CPT$ is a fundamental symmetry. Probing its violation is an important way to search for the new physics beyond the standard model. Up to now, $CPT$ symmetry has passed a number of high precision experimental tests and no definite signal of its violation has been observed in the laboratory. So, the present $CPT$ violating effects, if exist, should be very small to be amenable to the experimental limits. The $CPT$ symmetry could be dynamically violated in the expanding universe \cite{Li:2007}. The cosmological $CPT$ violation mechanism investigated in the literature \cite{Li:2001st,Li:2002wd,Li:2004hh,Feng:2004mq,Li:2007} has an interesting feature that the $CPT$ violating effects at present time are too small to be detected by the laboratory experiments but large enough in the early universe to account for the generation of matter-antimatter asymmetry \cite{Li:2001st,Li:2002wd,Li:2004hh,Li:2007}. And more importantly, this type of $CPT$ violating effects could be accumulated to be observable in the cosmological experiments \cite{Feng:2004mq,Li:2007,Feng:2006dp}. With the accumulation of high quality observational data, especially those from the CMB experiments, cosmological observation becomes a powerful way to test the $CPT$ symmetry. Here we study the CMB polarizations and $CPT$ violation in the photon sector with an effective lagrangian \cite{Carroll:1989vb,Carroll:1990zs}: \begin{equation}\label{Lagrangian} \mathcal{L} = -\frac{1}{4}F_{\mu\nu}F^{\mu\nu}+\mathcal{L}_{cs}~, \end{equation} where $\mathcal{L}_{cs}\sim p_{\mu}A_{\nu}\tilde F^{\mu\nu}$ is a Chern-Simons term, $p_{\mu}$ is an external vector and $\tilde F^{\mu\nu}=(1/2)\epsilon^{\mu\nu\rho\sigma}F_{\rho\sigma}$ is the dual of the electromagnetic tensor. This Lagrangian is not gauge invariant, but the action is gauge independent if $\partial_{\nu}p_{\mu}=\partial_{\mu}p_{\nu}$. This may be possible if $p_{\mu}$ is constant in spacetime or the gradient of a scalar field in the quintessential baryo-/leptogenesis \cite{Li:2002wd,Li:2001st,quin_baryogenesis} or the gradient of a function of the Ricci scalar in gravitational baryo-/leptogenesis \cite{Li:2004hh,R}. The Chern-Simons term violates Lorentz and $CPT$ symmetries when the background value of $p_{\mu}$ is nonzero. One of the physical consequences of the Chern-Simons term is the rotation of the polarization direction of electromagnetic waves propagating over large distances \cite{Carroll:1989vb}. From the Lagrangian (\ref{Lagrangian}), we can directly obtain the equation of motion for the electromagnetic field: \begin{equation}\label{maxwell1} \nabla_{\mu}(\nabla^{\mu}A^{\nu}-\nabla^{\nu}A^{\mu})=-p_{\mu} \epsilon^{\mu\nu\rho\sigma}(\nabla_{\rho}A_{\sigma}-\nabla_{\sigma}A_{\rho})~. \end{equation} After imposing Lorentz gauge condition $\nabla_{\mu}A^{\mu}=0$, it becomes: \begin{equation}\label{maxwell2} \nabla_{\mu}\nabla^{\mu}A^{\nu}+R^{\nu}_{\mu}A^{\mu}=-p_{\mu} \epsilon^{\mu\nu\rho\sigma}(\nabla_{\rho}A_{\sigma}-\nabla_{\sigma}A_{\rho})~, \end{equation} where $R^{\nu}_{\mu}$ is the Ricci tensor. With the geometric optics approximation, the solution to the equation of motion is expected to be: $A^{\mu}={\rm Re}[(a^{\mu}+\epsilon b^{\mu}+\epsilon^2 c^{\mu}+...)e^{iS/\epsilon}]$, where $\epsilon$ is a small number. With this ansatz, one can easily see that the Lorentz gauge condition implies $k_{\mu}a^{\mu}=0$, where the wave vector $k_{\mu}\equiv \nabla_{\mu}S$ is orthogonal to the surfaces of constant phase and represents the direction which photons travel along with. The vector $a^{\mu}$ is the product of a scalar amplitude $A$ and a normalized polarization vector $\varepsilon^{\mu}$, $a^{\mu}=A\varepsilon^{\mu}$, with $\varepsilon_{\mu}\varepsilon^{\mu}=1$. Hence in the Lorentz gauge, the wave vector $k_{\mu}$ is orthogonal to the polarization vector $\varepsilon^{\mu}$. Substituting this solution into the modified Maxwell equation (\ref{maxwell2}) and neglecting the Ricci tensor we have at the leading order of $\epsilon$ the equation is $k_{\mu}k^{\mu}=0$. It indicates that photons still propagate along the null geodesics. The effect of Chern-Simons term appears at the next order, $k^{\mu}\nabla_{\mu}\varepsilon^{\nu}=-p_{\mu}\epsilon^{\mu\nu\rho\sigma} k_{\rho}\varepsilon_{\sigma}$. We can see that the Chern-Simons term makes $k^{\mu}\nabla_{\mu}\varepsilon^{\nu}$ not vanished. This means that the polarization vector $\varepsilon^{\nu}$ is not parallel transported along the light-ray. It rotates as the photon propagates in spacetime. We consider here the spacetime described by spatially flat Friedmann-Robertson-Walker (FRW) metric. The null geodesics equation is $(k^0)^2-k^ik^i=0$. We assume that photons propagate along the positive direction of $x$ axis, \emph{i.e.} $k^{\mu}=(k^0,k^1,0,0)$ and $k^1=k^0$. Gauge invariance guarantees that the polarization vector of the photon has only two independent components which are orthogonal to the propagating direction. So, we are only interested in the changes of the components of the polarization vector, $\varepsilon^2$ and $\varepsilon^3$. Assuming $p_{\mu}=p_0$ to be a non-vanishing constant, we obtain the following equations: $d\varepsilon^2/d\lambda+\mathcal{H}k^0\varepsilon^2=p_0 k^0\varepsilon^3$, $d\varepsilon^3/d\lambda+\mathcal{H}k^0\varepsilon^3= - p_0 k^0\varepsilon^2$, where we have defined the affine parameter $\lambda$ which measures the distance along the light-ray, $k^{\mu}\equiv dx^{\mu}/d\lambda$, and the reduced expansion rate $\mathcal{H}\equiv \dot a/a$. The polarization angle is defined as $\chi\equiv\arctan{(\varepsilon^3}/{\varepsilon^2})$. It is easy to find that the rotation angle is \begin{equation}\label{kmu1} \Delta\chi\equiv \chi_0-\chi_z=-\int^{\eta_0}_{\eta_z} ~p_0 ~d\eta=\int^{t_z}_{t_0} p_0 ~\frac{dt}{a}~, \end{equation} where the subscript $z$ is the redshift of the source when the light was emitted. For CMB photons, the source is the last scattering surface with $z\simeq 1100$ \footnote{Besides the CMB photons which come from the last scattering surface, we might observed the different CMB photons which travelled different distances as well. However, we are only interested in linear perturbations of CMB photons in this paper. CMB polarizations are already linear phenomena. They are not existent at the zeroth order. When calculating the variations of polarizations due to $CPT$ violation in the perturbation theory up to linear order, we may ignore the fluctuations of the travelling distances of CMB photons. Otherwise, these fluctuations combined with polarizations would give higher order result which are beyond the scope of this paper. So, in this paper, we only consider linear perturbations and we can assume that each CMB photon detected by us travelled the same distance.}. The subscript $0$ indicates the present time. As we know, a vector rotated by an angle $\Delta\chi$ in a fixed coordinates frame is equivalent to a fixed vector observed in a coordinates frame which is rotated by $-\Delta\chi$. So, with the notion of coordinates frame rotation, the rotation angle is \begin{equation}\label{kmu2} \Delta\alpha=-\Delta\chi=\int^{t_0}_{t_z} ~ p_0 ~ \frac{dt}{a}~ = p_0 r_z~. \end{equation} with $r_z$ being the comoving distance of the light source away from us. This phenomena is known as ``cosmological birefringence". This rotation angle $\Delta\alpha$ can be obtained by observing polarized radiation from distant sources such as radio galaxies, quasars and CMB. The Stokes parameters $Q$ and $U$ of the CMB polarization can be decomposed into a gradient-like ($G$) and a curl-like ($C$) component \cite{Kamionkowski:1996ks}. For the standard theory of CMB, the $TC$ and $GC$ cross-correlation power spectra vanish. With the existence of cosmological birefringence, the polarization vector of each photon is rotated by an angle $\Delta\alpha$, and one would observe nonzero $TC$ and $GC$ correlations, even if they are zero at the last scattering surface. Denoting the rotated quantities with a prime, one gets \cite{Feng:2004mq,Lue:1998mq}: \begin{eqnarray}\label{modify} C_{l}^{'TC} &=& C_{l}^{TG}\sin(2\Delta\alpha)~, \nonumber\\ C_{l}^{'GC} &=& \frac{1}{2}(C_{l}^{GG}-C_{l}^{CC})\sin(4\Delta\alpha)~,\nonumber\\ C_{l}^{'TG} &=& C_{l}^{TG}\cos(2\Delta\alpha)~,\nonumber\\ C_{l}^{'GG} &=& C_{l}^{GG}\cos^2(2\Delta\alpha) + C_{l}^{CC}\sin^2(2\Delta\alpha)~,\nonumber\\ C_{l}^{'CC} &=& C_{l}^{CC}\cos^2(2\Delta\alpha) + C_{l}^{GG}\sin^2(2\Delta\alpha)~, \end{eqnarray} while the temperature power spectrum $TT$ remains unchanged. | 7 | 10 | 0710.3325 |
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0710 | 0710.3113_arXiv.txt | {In April $2006$ a $4$-channel acoustic antenna has been put in long-term operation on Lake Baikal. The detector was installed at a depth of about $100$ m on the instrumentation string of Baikal Neutrino Telescope NT200+. This detector may be regarded as a prototype of subunit for a future underwater acoustic neutrino telescope. We describe the design of acoustic detector and present first results obtained from data analysis.} \begin{document} | \begin{figure*}[t] \begin{center} \includegraphics [width=0.98\textwidth]{icrc0639_fig01.eps} \end{center} \caption{Schematic view of underwater 4-channel digital device for detection of acoustic signals from high energy neutrinos.} \label{fig1} \end{figure*} The large scale neutrino telescopes currently under operation (NT200+ in Lake Baikal, AMANDA/IceCube at the South Pole and ANTARES in the Mediterranean) detect the Cherenkov light emitted in water or ice by relativistic charged particles produced via neutrino interactions with matter. Back in 1957, G.A. Askaryan has shown that a high-energy particle cascade in water should also produce an acoustic signal \cite{Askarian-1957}. The absorption length for acoustic waves with a frequency about 30 kHz (the peak frequency of acoustic signals from a shower) in sea water is at least an order of magnitude larger than that of Cherenkov radiation, in the fresh Baikal water this ratio is even close to 100 \cite{Clay}. Therefore acoustic pulses can be detected from considerably larger distances than Cherenkov radiation, and the acoustic method appears to be attractive for the detection of ultra high-energy neutrinos \cite{Askarian-1977}. However, the technology of acoustic detection in high-energy physics is much worse developed than optical methods. Since several years, however, an increasing number of feasibility studies on acoustic particle detection are performed \cite{ARENA}. In order to test the possibility of acoustic detection of high-energy neutrinos in Lake Baikal, the Baikal collaboration started with an in-situ study of acoustic noise which constitutes the background for the acoustic neutrino detection in the lake. For the purpose of noise measurement, an autonomous hydro-acoustic recorder with two input channels has been developed. We have performed a series of hydro-acoustic measurements in Lake Baikal in order to investigate the the background properties \cite{Zeuthen-1, Akustika-noise}. It turned out that at stationary and homogeneous meteorological conditions the integral noise power in the frequency range $20$-$50$ kHz can reach levels as low as about $1$ mPa. At the same time, short acoustic pulses with different amplitudes and shapes including bipolar ones have been observed. The latter should be considered as a background for acoustic neutrino detection. However, the overwhelming majority of the short pulses have probably been generated by quasi-local sources or are due to interference of noise sound waves coming from a layer near the surface. Taking into account these properties of the noise, we conclude that the most promising way to detect acoustic particle signals is to deploy a net of rather compact acoustic antennas at relatively shallow depths (for example about $100$--$200$ m for Lake Baikal) and monitoring the water volume top-down. It is also necessary to suppress signals from the surface by caps made of a sound-absorbing material and mounted on top of the antennas. | The results of the experiment have demonstrated the feasibility of the proposed acoustic pulse detection technique in searching signals from cascade showers. Although the Baikal water temperature is close to the temperature of its maximal density, the absence of strong acoustic noise sources in the lake's deep zone, and the very low absorption of sound in freshwater may result in neutrino detection in Lake Baikal with a threshold as low as $10^{18}$ -- $10^{19}$ eV. This motivates further activities towards a large-scale acoustic neutrino detector in Lake Baikal. | 7 | 10 | 0710.3113 |
0710 | 0710.5503_arXiv.txt | We present relations of the black hole mass and the optical luminosity with the velocity dispersion and the luminosity of the \nev\ and the \oiv\ high-ionization lines in the mid-infrared (MIR) for 28 reverberation-mapped active galactic nuclei. We used high-resolution \spi\ Infrared Spectrograph and {\it Infrared Space Observatory} Short Wavelength Spectrometer data to fit the profiles of these MIR emission lines that originate from the narrow-line region of the nucleus. We find that the lines are often resolved and that the velocity dispersion of \nev\ and \oiv\ follows a relation similar to that between the black hole mass and the bulge stellar velocity dispersion found for local galaxies. The luminosity of the \nev\ and the \oiv\ lines in these sources is correlated with that of the optical 5100$\ang$ continuum and with the black hole mass. Our results provide a means to derive black hole properties in various types of active galactic nuclei, including highly obscured systems. | \label{sec:intro} Following the relation between the black hole (BH) mass, \mbh, and the bulge stellar velocity dispersion, $\sigma_*$, (\citealt{ferrarese}; \citealt{gebhardt}; \citealt{tremaine}), a similar relation was found for the velocity dispersion $\sigma$ of the \oiii\ 5007\ang\ line (\citealt{nelson00}; \citealt{greene}) that originates from the narrow-line region (NLR) gas of the active galactic nucleus (AGN). The NLR gas is thought to be mostly gravitationally bound to the bulge in high-luminosity AGNs (\citealt{nelson96}; \citealt{greene}; \citealt{laor07}), unlike the broad-line region gas that is virialized due its proximity to the BH (\citealt{peterson99}). This relation is useful for systems in which the stellar absorption lines cannot be observed, e.g., because of dilution of the stellar light by the AGN continuum. In addition to this relation, \mbh\ and the luminosity $L$ of the 5100 \ang\ optical continuum, \lopt , are correlated in AGNs with BH measurements (\citealt{kaspi00}). There are two main advantages to expanding such relations in the mid-infrared (MIR). The first is that the \ion{Ne}{5} and \ion{O}{4} ions emitting at 14.32 and 25.89 \micron\ cannot be easily excited by star-forming regions since they have ionization potentials $\chi$ of 97.12 and 54.93 eV. The second reason is the low obscuration, which allows for the results to be applied to type 2 AGNs. In this Letter, we investigate for relations between \mbh\ and \lopt\ with MIR NLR line velocity dispersions and luminosities. | \label{sec:origin} The relations between \snev,\soiv\ and \mbh\ indicate that the NLR gas kinematics are primarily determined by the potential of the bulge, as it is also believed based on the \oiii\ line profiles (\citealt{whittle92}; \citealt{nelson96}; \citealt{greene}). However, \snev\ and \soiv\ are on average 67 and 62 \kms\ higher than $\sigma_*$, and 51 and 32 \kms\ higher than \soiii . Such discrepancies could be attributed to an increase of the line width with increasing $\chi$, which is often found for optical NLR lines (e.g. \citealt{osterbrock}; \citealt{whittle85b}). It is possible that ions with high ionization potentials have high velocity dispersions because they are located in NLR clouds that are close to the BH sphere of influence. Since the \ion{O}{4} and, mostly, the \ion{Ne}{5} ions are predominantly excited by AGNs, the MIR \msigma\ relation can be applied even in systems that undergo starbursts. It can also be applied in type 2 AGNs since the lines suffer little from extinction. Moreover, \snev\ and \soiv\ can be used as surrogates for $\sigma_*$ in environments where measuring $\sigma_*$ is hard, such as in bright QSOs. However, the relation does not necessarily hold for AGNs with strong winds and jets such as luminous radio sources (\citealt{whittle92}), and for AGNs with high Eddington rates, \eedd\ (\citealt{greene}). Such systems can be recognized by the asymmetric wings or the high kurtosis of their line profiles (Whittle 1985a; 1992). The luminosities of the broad lines (\citealt{baldwin}) and the narrow lines (Fig.~\ref{fig:ff}) scale with \lopt\ and with the bolometric AGN luminosity, \lbol , since they depend on the absolute accretion rate onto the central object. Specifically, \lbol\ is equal to \xopt \lopt, \xnev \lnev, and \xoiv \loiv , where $x$ is a wavelength-dependent bolometric correction factor. For \xopt $=$9 (\citealt{kaspi00}; \citealt{marconi04}) and for the median values of \lopt/\lnev\ and \lopt/\loiv\ in our sample, we find that \xnev $=$13000 and \xoiv $=$4500. That the relation does not have a slope of unity could indicate that \xnev\ and \xoiv\ depend on $L$ in a manner different than \xopt\ does. \cite{netzer06} found that the equivalent width of \oiii\ decreases with increasing \lopt. This behaviour could be attributed to a different geometric distribution of NLR clouds around the AGN at different luminosities. Whether the \nev\ or \oiv\ luminosity of a source can be used to derive its \mbh\ partly depends on its Eddington rate. To illustrate how changes of \eedd\ affect the relation between the luminosity of a MIR line and \mbh , we overplot lines of constant \eedd\ in Figure~\ref{fig:ml}. The more quiescent a source is, the more its position is shifted to the top left corner of this diagram. Equation~(\ref{rel3}) is valid within the \eedd\ range 0.003$-$0.6 that our sources span, and implies that the mass of a BH determines its luminosity output. Its advantage is that it can be applied to sources with obscured optical continua or underestimated \oiii\ luminosities, such as type 2 AGNs (\citealt{netzer06}). However its scatter is likely to be larger than that presented in \S~\ref{sec:mir_opt} since the Eddington rates of the reverberation-mapped AGNs are not necessarily representative of those of all local AGNs. We conclude that the MIR NLR gas kinematics trace \mbh\ in local reverberation-mapped AGNs in a manner similar to stellar kinematics. The \nev\ and \oiv\ line widths can be used to estimate \mbh. The calibration of \lnev\ and \loiv\ to \lopt\ provides a new method to compute bolometric luminosities, black hole masses, and Eddington rates in various types of AGNs, including highly obscured systems, albeit with a large scatter.\\ \\ This work was based on observations made with the {\it Spitzer Space Telescope}, % and was supported by NASA through an award issued by JPL/Caltech. | 7 | 10 | 0710.5503 |
0710 | 0710.2115_arXiv.txt | With a goal toward deriving the physical conditions in external galaxies, we present a survey of the formaldehyde emission in a sample of starburst systems. By extending a technique used to derive the spatial density in star formation regions in our own Galaxy, we show how the relative intensity of the $1_{10}-1_{11}$ and $2_{11}-2_{12}$ K-doublet transitions of H$_2$CO can provide an accurate densitometer for the active star formation environments found in starburst galaxies. Relying upon an assumed kinetic temperature and co-spatial emission and absorption from both H$_2$CO transitions, our technique is applied to a sample of nineteen IR-bright galaxies which exhibit various forms of starburst activity. In the five galaxies of our sample where both H$_2$CO transitions were detected we have derived spatial densities. We also use H$_2$CO to estimate the dense gas mass in our starburst galaxy sample, finding similar mass estimates for the dense gas forming stars in these objects as derived using other dense gas tracers. A related trend can be seen when one compares $L_{IR}$ to our derived $n(H_2)$ for the five galaxies within which we have derived spatial densities. Even though our number statistics are small, there appears to be a trend toward higher spatial density for galaxies with higher infrared luminosity. This is likely another representation of the $L_{IR}$-$M_{dense}$ correlation. | \label{intro} Studies of the distribution of Carbon Monoxide (CO) emission in external galaxies (\cf\ \cite{Young1991}) have pointed to the presence of large quantities of molecular material in these systems. These studies have yielded a detailed picture of the molecular mass in many external galaxies. But, because emission from the abundant CO molecule is generally dominated by radiative transfer effects, such as high optical depth, it is not a reliable monitor of the physical conditions, such as spatial density and kinetic temperature, quantities necessary to assess the possibility of star formation. Emission from less-abundant, higher-dipole moment molecules are better-suited to the task of deriving the spatial density and kinetic temperature of the dense gas in our and external galaxies. For this reason, emission line studies from a variety of molecules have been made toward mainly nearby galaxies (see \cite{Mauersberger1989} (CS), \cite{Gao2004a} (HCN), \cite{Nguyen1992} (HCO$^+$), \cite{Mauersberger1990} and \cite{Meier2005} (HC$_3$N), \cite{Mauersberger2003} (NH$_3$), or \cite{Henkel1991} for a review). The most extensive sets of measurements of molecular line emission in external galaxies has been done using the J=1-0 transitions of CO \citep{Helfer2003} and HCN \citep{Gao2004a}. Since the J=1-0 transitions of CO and HCN are good tracers of the more generally distributed and the denser gas, respectively, but do not provide comprehensive information about the individual physical conditions of the dense, potentially star-forming gas, another molecule must be observed for this purpose. Formaldehyde (H$_2$CO) has proven to be a reliable density and kinetic temperature probe in Galactic molecular clouds. Existing measurements of the H$_2$CO $1_{10}-1_{11}$ and $2_{11}-2_{12}$ emission in a wide variety of galaxies by \cite{Baan1986}, \cite{Baan1990}, \cite{Baan1993}, and \cite{Araya2004} have mainly concentrated on measurements of the $1_{10}-1_{11}$ transition. One of our goals with the present study was to obtain a uniform set of measurements of both K-doublet transitions with which the physical conditions, specifically the spatial density, in the extragalactic context could be derived. Using the unique density selectivity of the K-doublet transitions of H$_2$CO we have measured the spatial density in a sample of galaxies exhibiting starburst phenomena and/or high infrared luminosity. In \S\ref{H2coProbe} we discuss the specific properties of the H$_2$CO molecule which make it a good probe of spatial density. \S\ref{Observations} presents our observation summary; \S\ref{Results} our H$_2$CO, OH, H111$\alpha$, and continuum emission measurement results; \S\ref{Analysis} analyses of our H$_2$CO, OH, and H111$\alpha$ measurements, including Large Velocity Gradient (LVG) model fits to and dense gas mass calculations based on our H$_2$CO measurements. | \label{Conclusions} Using measurements of the $1_{10}-1_{11}$ and $2_{11}-2_{12}$ K-doublet transitions of H$_2$CO we have derived accurate \textit{measurements} of the spatial density (n(H$_2$)) in a sample of starburst galaxies. The derived densities range from $10^{4.7}$-$10^{5.7}$~cm$^{-3}$, consistent with the suggestion that the high infrared brightness of these galaxies is driven by extreme star formation activity. We believe that these spatial density measurements are the most accurate measurements of this important physical quantity in starburst galaxies made to-date. We have also used our H$_2$CO measurements to derive a measure of the dense gas mass which ranges from $0.06-77\times10^8 M_\odot$, generally consistent with previous measurements, mainly from studies of the HCN emission in these galaxies. The linear correlation between the IR luminosity ($L_{IR}$) and the dense gas mass ($M_{dense}$) found in larger samples of the HCN emission in starburst galaxies is also apparent in our H$_2$CO measurements. This further supports the suggestion that active star formation in IR-bright galaxies is driven by the amount of material available to form stars. We also note a related trend between $L_{IR}$ and our derived $n(H_2)$ for the five galaxies within which we have derived spatial densities. The three galaxies with lower IR luminosities have lower ($n(H_2) \simeq 10^5$~cm$^{-3}$) derived spatial densities, while the two higher luminosity galaxies have higher ($n(H_2) = 10^{5.7}$~cm$^{-3}$) derived spatial densities. This is likely another representation of the $L_{IR}$-$M_{dense}$ correlation. | 7 | 10 | 0710.2115 |
0710 | 0710.0867_arXiv.txt | The observed cosmic acceleration presents the physics and cosmology communities with amazing opportunities to make exciting, probably even radical advances in these fields. This topic is highly data driven and many of our opportunities depend on us undertaking an ambitious observational program. Here I outline the case for such a program based on both the exciting science related to the cosmic acceleration and the impressive impact that a strong observational program would have. Along the way, I challenge a number of arguments that skeptics use to question the value of a strong observational commitment to this field. | This is truly remarkable time to be involved in cosmology research. There are many reasons for this, but none stand out quite as dramatically as the observed cosmic acceleration (attributed in current nomenclature to the ``dark energy''). In the words of the Dark Energy Task Force (DETF)~\cite{Albrecht:2006um} ``most experts believe that nothing short of a revolution in our understanding of fundamental physics will be required to achieve a full understanding of the cosmic acceleration''. As things stand, this revolution is being motivated both by remarkable new data sets and by exciting theoretical developments. Many ambitious researchers in cosmology and related fields have been galvanized by these extraordinary developments. There have been a number of excellent talks at this conference on the topic of cosmic acceleration which capture some of this climate. The DETF, charged with charting a way forward with dark energy observations, received 50 thoughtful and thorough whitepapers from leaders in the field (despite the lack of any specific commitment at that time to fund future dark energy experiments). Most of us look at these developments and are astonished at our good fortune to be part of what will surely be viewed as one of the great moments in the history of science. Indeed, since the discovery of dark energy every group that has deliberated on future directions for the field has recognized the exciting opportunities and challenges presented by the cosmic acceleration\footnote{See for example \cite{Turner:2003pe,Albrecht:2005np,Peacock:2006kj}. Since the 2006 DETF report a number of panels have recommended pursuit of specific ground and space based dark energy projects, and just since PASCOS 07 the US National Research Council's {\em Committee on NASA's Beyond Einstein Program} named a ``Joint Dark Energy Mission'' the top priority for that program \url{http://nationalacademies.org/morenews/20070907b}}. But a small number of negative voices continue to be heard alongside the building enthusiasm for further studies of the dark energy. To some degree this negativity may reflect the natural skepticism of scientists in the face of unbridled enthusiasm (such as the above). To the extent that that is the explanation, the skepticism is surely a good thing that will lead to a healthy debate and produce more rigorous research. On the other hand, I suspect some of the negativity is just the result of sloppy thinking and needs to be challenged and simply put to rest. Regardless of how one might attribute explanations and motivations, one purpose of this paper is to engage in this debate on a number of fronts. Here are some illustrations of some of the negative views I am talking about. At lunch on the first day of the PASCOS 07 conference there was a lively discussion about the merits of dark energy experiments. One colleague commented \begin{quote} {\em ``Studies of dark energy are unlikely to be interesting because we already have a theory of dark energy.''} \end{quote} moments later, another cosmologist declared \begin{quote} { \em ``Studies of dark energy are unlikely to be interesting because we have no theory of dark energy.''} \end{quote} The two were united in their conclusion, and apparently not too bothered by subtle differences in their reasoning One unavoidable feature of physics research is that at any given time the total amount of data is finite. That necessarily means there will be more than one theory that fits the data. In many fields, this universal fact is seen as a source of vitality. Curiosity about further resolving these degeneracies drives exciting new experiments and theoretical work. The momentous progress in physics over the last century can be seen in this light as the fundamental particles, the atomic theory of heat, quantum physics and general relativity emerged from ``under the radar'' of earlier data sets that were fit perfectly by more primitive theories. As was amply evident in the talks at the PASCOS 07 conference, we have good reason to look forward to similar advances as the LHC data starts coming in. But for some reason, the fact that a given future dark energy experiment will not remove all uncertainties about the nature of dark energy seems to generate considerable angst among some physicists and astronomers\footnote{See the question section of my PASCOS 07 talk for an example\cite{PASCOS07}} and is sometimes given as a reason to be discouraged from even doing more experiments. As I shall quantify below (and as was also shown by others at PASCOS 07), the proposed new experiments will have an impressive impact on our knowledge of dark energy. This is all one can ever ask of a new experiment. This paper focuses on two key areas. In the next section I review some of the thriving theoretical work that has been stimulated by the cosmic acceleration. The striking theoretical issues raised by the cosmic acceleration and the remarkable directions we have been driven in our initial attempts to understand it are the key reasons I find this topic so deeply interesting. I also believe this is why so many excellent researchers are taking risks and changing direction in their careers in order to get involved. The third section summarizes a number of results which demonstrate the tremendous impact future experiments can have on our understanding of dark energy, both on our understanding of the general properties of dark energy and in terms of constraining specific models of dark energy that are currently of interest. It is these results that demonstrate that an ``aggressive program of dark energy probes'' is indeed possible. This paper is based on my talk at the PASCOS 07 meeting at Imperial College. My slides and a video of the talk are available online\cite{PASCOS07}. This online material as well as my ``Origins of Dark Energy'' talk\cite{ODE} and related papers\cite{Albrecht:2007qy,Abrahamse:2007ip} are a good source for the technical material on which this paper is based. My goal here is to assemble some key arguments in a concise form. Readers seeking more details should refer to this other material. | \label{Sect:CAS} The case for aggressive pursuit of new data on dark energy is twofold. Firstly, I have outlined how the subject of dark energy has generated very exciting and often radical new theoretical ideas. These include dramatic proposals that change how we think about equilibrium and initial conditions in cosmology and even how we formulate fundamental theories. Second, impressive new experiments are within reach that could have a tremendous impact on our understanding of dark energy. The best experiments will constrain dark energy properties orders of magnitude better than the current or medium sized future experiments, and will have the ability to strongly discriminate among and even fully eliminate popular dark energy models based on subtle variations in the equation of state. I have outlined and challenged some of the skeptical perspectives I have heard regarding future dark energy studies. The discovery of the cosmic acceleration has caused a great upheaval in our thinking about fundamental physics and cosmology. I often sense that the skeptics are hoping this upheaval will end quickly and are grasping for arguments that will allow things to rapidly return to normal. I feel this outcome is very unlikely, and this is exactly why I find the topic so exciting. Nature has handed us an amazing opportunity. I hope that the physics and cosmology communities have the strength to face the challenge of the cosmic acceleration head on and give a response that we can be proud of when people write the history of this era. \begin{theacknowledgments} I would like to thank A. Abrahamse, M. Barnard, G. Bernstein, B. Bozek, L. Sorbo, M. Yashar and the DETF members who collaborated with me on some of the work reviewed here, and B. Bozek for helpful comments on the manuscript. Also, I thank the organizers, especially Arttu Rajantie, for a really excellent conference. This work was supported in part by DOE grant DE-FG03-91ER40674 and NSF grant AST-0632901. \end{theacknowledgments} | 7 | 10 | 0710.0867 |
0710 | 0710.5359_arXiv.txt | {% The study of rotation and activity in low-mass stars or brown dwarfs of spectral classes M and L has seen enormous progress during the last years. I summarize the results from different works that measured activity, rotation, and sometimes magnetic fields. The generation of magnetic activity seems to be unchanged at the threshold to completely convective stars, i.e. no change in the efficiency of the magnetic dynamos is observed. On the other hand, a sudden change in the strength of rotational braking appears at the threshold mass to full convection, and strong evidence exists for rotational braking weakening with lower mass. A probable explanation is that the field topology changes from dipolar to small scale structure as the objects become fully convective. } | Rotation and activity are intimately connected in sun-like stars. Rotation is believed to generate a magnetic field through a dynamo mechanism that scales with rotation. The stellar wind couples to the rotating magnetic field lines carrying away angular momentum so that the star is being braked. Hence stars of spectral type F--K rotate more rapidly and are more active when they are young, but they decelerate and become less active as they age; this is the so-called rotation-activity connection (Noyes et al., 1984; Pizzolato et al., 2003). The scaling of activity with rotation depends on the type of dynamo inside the star. The fact that the rotation-activity connection is similar in virtually all stars that harbor convective envelopes and at all ages (Pizzolato et al., 2003; Reiners, 2007) indicate that the dynamo mechanism is comparable in all these stars. Around spectral type M3.5, the internal structure of the stars changes; stars later than around M3.5 are believed to be fully convective. They cannot harbor an interface dynamo working at the tachocline, which is believed to be the most important dynamo mechanism in the hotter sun-like stars. If the interface dynamo was the only important mechanism driving a magnetic dynamo, one could expect a sharp break in magnetic field generation around spectral type M3.5. Such a break would imply a sudden change in observable stellar activity and in the braking of stellar rotation. A break in stellar activity is not observed in X-rays or H$\alpha$ in the surveys that crossed the M3.5 border (e.g., Delfosse et al., 1998; Mohanty \& Basri, 2003; West et al., 2004). Activity rather stays at comparable levels if normalized to bolometric luminosity, and quite surprisingly the fraction of active stars even raises up to $\sim 80\,\%$ in late M-type objects before it goes down again. This clearly shows that the interface dynamo is not the only dynamo operating in stellar interiors and that fully convective stars can have quite efficient dynamos as well. On the other hand, Delfosse et al., 1998, presented a plot that shows a different behavior in rotational braking in stars later than spectral type M3.5. In their Fig.\,3, they show that all M dwarfs earlier than M3.5 are slow rotators ($v\,\sin{i} \la 3$\,km\,s$^{-1}$) regardless of what disk population they belong to (young or old). M stars later than M3.5, however, show substantial rotation velocities of up to $v\,\sin{i} = 50$\,km\,s$^{-1}$ in the young disk population, and in the old disk population some late M dwarfs still have velocities around $v\,\sin{i} = 10$\,km\,s$^{-1}$. This can be interpreted as an indication for a sudden change in the timescales of rotational braking at the mass where stars become fully convective. Delfosse et al. conclude that spin-down timescales are on the order of a few Gyrs at spectral type M3--M4, and of the order of 10\,Gyr at spectral type M6. The investigation of rotation and activity in ultra-cool stars (M7 and later) was put on firm ground by Mohanty \& Basri, 2003. From high-resolution spectra they determined projected rotation velocities, and from H$\alpha$ emission they derived the level of activity. Reiners \& Basri, 2007, added more M stars to this sample. In addition, they measured magnetic fields in low-mass M dwarfs and showed that in M dwarfs the level of activity is still coupled to magnetic flux -- high magnetic flux levels lead to strong H$\alpha$ emission and stars without H$\alpha$ emission show no magnetic fields. | High resolution spectroscopy in M and L dwarfs becomes a technique that allows to investigate the physics and the evolution of low mass objects in great detail. No change in H$\alpha$ activity is observed at the threshold to complete convection. Activity can be followed to objects as late as mid-M and it might continue to even lower masses but below the current observational threshold. The decline in normalized activity among the ultra-cool dwarfs could be explained by the enhanced electrical resistivity at the low temperatures. Currently, no indications for a mass or temperature dependence of stellar or substellar dynamos can be concluded from the activity measurements. A sudden change in the behavior of rational braking at spectral class M3.5 was already found by Delfosse et al., 1998. In the young disk, stars earlier than M3.5 rotate slowly while later stars still show significant rotation. The lack of slowly rotating ultra-cool dwarfs and the rise of the minimum rotation rate with spectral class could be explained by rotational braking that is weaker with lower mass or lower temperature. The reason for a weaker braking at later spectral type could be a different magnetic topology. The geometry of the magnetic field is essential for the strength of magnetic braking as described in Krishnamurti et al. (1997) and Sills et al. (2000). These authors find that the description of magnetic braking ($\omega_\mathrm{crit}$) needs to be different in mid-M stars and in earlier objects, which they connect to different convective turnover times. Although the details of magnetic braking in ultra-cool dwarfs are not quite understood, a substantial amount of evidence exists that rotational braking is weaker with lower mass or lower temperature. The limiting factor of rotational braking may be the topology of the magnetic fields which in fully convective stars might be generated on smaller scales so that the topology is different from a dipolar configuration leading to the weaker rotational braking. | 7 | 10 | 0710.5359 |
0710 | 0710.3396_arXiv.txt | We modify a stellar structure code to estimate the effect upon the main sequence of the accretion of weakly interacting dark matter onto stars and its subsequent annihilation. The effect upon the stars depends upon whether the energy generation rate from dark matter annihilation is large enough to shut off the nuclear burning in the star. Main sequence WIMP burners look much like protostars moving on the Hayashi track, although they are in principle completely stable. We make some brief comments about where such stars could be found, how they might be observed and more detailed simulations which are currently in progress. Finally we comment on whether or not it is possible to link the paradoxically young OB stars found at the galactic centre with WIMP burners. | 7 | 10 | 0710.3396 |
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0710 | 0710.3169_arXiv.txt | We compute the electromagnetic radiative corrections to all leading annihilation processes which may occur in the Galactic dark matter halo, for dark matter in the framework of supersymmetric extensions of the Standard Model (MSSM and mSUGRA), and present the results of scans over the parameter space that is consistent with present observational bounds on the dark matter density of the Universe. Although these processes have previously been considered in some special cases by various authors, our new general analysis shows novel interesting results with large corrections that may be of importance, e.g., for searches at the soon to be launched GLAST gamma-ray space telescope. In particular, it is pointed out that regions of parameter space where there is a near degeneracy between the dark matter neutralino and the tau sleptons, radiative corrections may boost the gamma-ray yield by up to three or four orders of magnitude, even for neutralino masses considerably below the TeV scale, and will enhance the very characteristic signature of dark matter annihilations, namely a sharp step at the mass of the dark matter particle. Since this is a particularly interesting region for more constrained mSUGRA models of supersymmetry, we use an extensive scan over this parameter space to verify the significance of our findings. We also re-visit the direct annihilation of neutralinos into photons and point out that, for a considerable part of the parameter space, internal bremsstrahlung is more important for indirect dark matter searches than line signals. | During the last few years a strong consensus has emerged about the existence of a sizeable dark matter contribution to the total cosmological energy density. The identification of experimental signatures that eventually may determine the nature of the cosmological dark matter is thus becoming ever more important. The present estimates \cite{wmap} give the fraction of the critical density of cold dark matter particles as $\Omega_{CDM}h^2\sim 0.105 \pm 0.013$, where the Hubble parameter (scaled in units of 100 km/s Mpc$^{-1}$) is $h\sim 0.70\pm 0.02$. Also, on the scales of galaxies and smaller, a number of methods including measurements of rotation curves as well as gravitational lensing agree well with the predictions from N-body calculations of gravitational clustering in cold dark matter cosmologies (see e.g. \cite{vialactea}). The methods of detection of dark matter (for reviews, see \cite{reviews}) can be divided into {\em accelerator} production and detection of missing energy (especially at the LHC at CERN, which will start operating some time in 2008), {\em direct detection} (of dark matter particles impinging on a terrestrial detector, with recent impressive upper limits reported by \cite{direct}), or {\em indirect detection} of particles generated by the annihilation of dark matter particles in the Galactic halo or in the Sun/Earth. All these methods are indeed complementary -- it is probable that a signal from more than one type of experiment will be needed to fully identify the particle making up the dark matter. The field is just entering very interesting times, with the LHC soon starting and new detectors of liquid noble gases being developed for direct detection. For indirect detection, the satellite PAMELA \cite{pamela} was launched a year ago and will soon reveal its first sets of data for positron and antiproton yields in the cosmic rays \cite{antimatter}. AMANDA \cite{amanda} at the South Pole that has searched for detection of neutrinos from the centre of the Earth or the Sun \cite{neutrinos}, will soon give way to the much larger detector IceCUBE \cite{icecube}, and for gamma-rays coming from annihilations of dark matter particles in the halo \cite{gammas} the space satellite GLAST \cite{glast}, to be launched in 2008, will open up a new window to the high-energy universe, for energies from below a GeV to about 300 GeV. One problem with all these discovery methods is that the signal searched for may be quite weak, with much larger backgrounds in many cases. For indirect detection through gamma-rays, the situation may in principle be better, due to (i) the direct propagation from the region of production, without significant absorption or scattering; (ii) the dependence of the annihilation rate on the square of the dark matter density which may give "hot spots" near density concentrations as those predicted by N-body simulations; (iii) possible characteristic features like gamma-ray lines or steps, given by the fact that no more energy than $m_{\chi}$ per particle can be released in the annihilation of two non-relativistic dark matter particles (we denote the dark matter particle by $\chi$). As an example, it was recently shown \cite{idm} that in models of an extended Higgs sector, the line signal from the two-body final states $\gamma\gamma$ and $Z\gamma$ could give a spectacular signature in the gamma-ray spectrum between 40 and 80 GeV. On the other hand, in models of universal extra dimensions (UED) \cite{uedline} or in the theoretically perhaps most favoured, supersymmetric, models of dark matter the line feature is in general not very prominent, except in some particular regions of the large parameter space. However, it was early realised that there could be other important spectral features \cite{lbe89}, and recently it has been shown that internal bremsstrahlung (IB) from produced charged particles in the annihilations could yield a detectable "bump" near the highest energy for heavy gauginos or Higgsinos annihilating into $W$ boson pairs, such as expected in split supersymmetry models \cite{heavysusy}. In \cite{birkedal}, it was furthermore pointed out that IB often can be estimated by simple, universal formulas and often gives rise to a very prominent step in the spectrum at photon energies of $E_\gamma=m_\chi$ (such as in UED models \cite{Bergstrom:2004cy}). Encouraged by these partial results, we have performed a detailed analysis of the importance of IB in the minimal supersymmetric extension to the standard model (MSSM). We have therefore calculated the IB contributions for all two-particle charged final states from general neutralino annihilations. Besides confirming the mentioned partial results for the universal radiative corrections, in particular those relating to soft and collinear bremsstrahlung, we also point out interesting cases of model-dependent ``virtual'' brems\-strahlung (i.e. photons emitted from charged virtual particles), see Fig.~1. We confirm the suspicion expressed already in \cite{lbe89} that this type of emission may circumvent the chiral suppression, i.e., the annihilation rate being proportional to $m_f^2$ for annihilation into a fermion pair from an $S$-wave initial state, as is the case in lowest order for non-relativistic dark matter Majorana particles in the Galactic halo (see also \cite{baltz_bergstrom}). Since this enhancement mechanism is most prominent in cases where the neutralino is close to degenerate with charged sleptons, it is of special importance in the so-called stau coannihilation region in models of minimal supergravity (mSUGRA, as implemented in \cite{isajet}). We therefore run through an extensive scan over these models (based on \cite{baltz_peskin}) and find, indeed, remarkable cases of enhancement of the gamma-ray rate in the stau coannihilation region, near the maximal possible photon energy $E_\gamma=m_\chi$. Let us stress that the radiative corrections to the main annihilation channels, here computed systematically for the first time, may turn out to be of utmost importance when fitting gamma-ray data, e.g. from GLAST, to supersymmetric dark matter templates. Over much of the parameter space we have scanned, these corrections give a large factor of enhancement over the commonly adopted estimates, especially at the observationally most interesting, highest energies. More importantly, they add a feature, the very sharp step at the dark matter mass, that would distinguish this signal from all other astrophysical background (or foreground) processes. | As can be seen already from our benchmark points in Table~\ref{benchmark}, and in more detail from the scatter plots in Fig.~\ref{FSRcomp1}, the internal bremsstrahlung effects computed in this work can be very significant, changing sometimes by more than an order of magnitude the lowest-order prediction for the high-energy gamma-ray signal from neutralino dark matter annihilation. Although some of these enhancements have been found before \cite{lbe89,heavysusy,birkedal}, this is the first time the first-order radiative corrections have been computed systematically, for all relevant final states in supersymmetric dark matter models. The resulting enhancements of the expected fluxes are surprisingly large over significant regions in the parameter space of the MSSM, including the more constrained mSUGRA models. Despite the fact that some large corrections apply to absolute rates that are too small to be of practical interest, Fig.~4 shows that the quantity ${\cal S}$, which is directly proportional to the expected signal in gamma-ray detection experiments, also is significant for the internal bremsstrahlung contribution in large regions of parameter space. For $m_\chi < 300$ GeV, for example, values of ${\cal S}_{IB}$ greater than 0.1 are generic, and for masses below 100 GeV, values of 1 or higher are common, which in very many cases is higher than the corresponding values for the line signals $\gamma\gamma$ and $Z\gamma$. One should also bear in mind that the sensitivity of Air Cherenkov Telescopes increases significantly with energy; detectional prospects for a $m_\chi\sim1~$TeV neutralino with ${\cal S}\sim0.01$, e.g., correspond very roughly to those for a $m_\chi\sim100~$GeV neutralino with ${\cal S}\sim0.5$ (see, e.g., \cite{Morselli:2002nw}). In this light, the situation becomes very interesting even for TeV scale Higgsinos, where IB generically contributes more than 10 times as much as secondary photons. We note that (as anticipated in \cite{lbe89}) helicity suppression and also CP selection rules of certain final states may be circumvented by emitting a photon; this is for example the origin of the very substantial enhancements of the signal obtained in the stau annihilation region in mSUGRA models. In this situation, the probability of emitting gamma rays vanishes at zero photon energy but increases rapidly at high energy (see Fig.~\ref{fig:spec}b and \ref{fig:spec}c), which gives a photon ``bump'' at $E_\gamma\approx m_\chi$. We also note that in less constrained versions of the MSSM than considered here, we expect even more situations where large enhancements of the annihilation signal due to internal bremsstrahlung can be found. An example is heavy Wino dark matter \cite{heavysusy,Chattopadhyay:2006xb}, which becomes possible when relaxing the condition $M_1\approx \frac{1}{2}M_2$ (as realized, e.g., in anomaly mediated supersymmetry breaking scenarios \cite{amsmb}). Of course, the line signals, in particular $\gamma\gamma$, have the virtue of being at the highest possible energy, so in order to make a more accurate comparison between these and the IB signal computed here, one would have to model also the expected spectral shape of possible astrophysical gamma-ray backgrounds and the energy resolution of the detector. This is left for future work \cite{torstenetal}. We note, however, that in general also the new contributions have a characteristic signature, ({\em cf.} Fig.~2) which can hardly be mimicked by any known astrophysical gamma-ray source. In fact, in some cases these spectra could even be used by future experiments to distinguish between different dark matter candidates (note that, e.g., the distinction between Kaluza-Klein dark matter and a neutralino in the focus point region like our BM4 point would be possible already with the energy resolution of present Air Cherenkov Telescopes \cite{Bergstrom:2006hk}). To conclude, we have shown that the commonly neglected first-order radiative corrections to neutralino dark matter annihilation should definitely be taken into account when predicting rates for gamma-ray telescopes. In particular, the soon to be launched GLAST space telescope \cite{glast} will have an enhanced possibility over what has previously been assumed to detect radiation from supersymmetric dark matter annihilation. The routines needed to compute these new processes will be included in the next release of the \ds\ package \cite{ds,joakim}. \smallskip | 7 | 10 | 0710.3169 |
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