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0710 | 0710.0795_arXiv.txt | Two aspects of our recent N-body studies of star clusters are presented:\\ 1) What impact does mass segregation and selective mass loss have on integrated photometry?\\ 2) How well compare results from N-body simulations using NBODY4 and STARLAB/KIRA? | ``Mass segregation'' describes the effect that high-mass stars are preferentially found in the center of clusters, while the outskirts are preferentially occupied by lower-mass stars. (Primordial) mass segregation is studied observationally in a number of clusters. Some examples are: the Orion Nebula Cluster (\cite{1998ApJ...492..540H}), 6 LMC clusters (\cite{2002MNRAS.331..245D}). Dynamical mass segregation is also found from N-body simulations (e.g. \cite{1985ApJ...292..339I,1997MNRAS.286..709G}), caused by two-body encounters and energy equipartition. A number of authors also point at the preferential mass loss of low-mass stars when the cluster evolves in a tidal field (e.g. \cite{1975ApJ...201..773S,1997MNRAS.286..709G}), as mass segregation populates the cluster outskirts preferentially with low-mass stars where they are most easily stripped from the cluster potential. The resulting changes in the overall stellar mass function inside the cluster are quantified by \cite{2003MNRAS.340..227B}. In \cite{2006A&A...452..131L} we use the results from \cite{2003MNRAS.340..227B} to incorporate the changing mass function slope in a simplified manner into the GALEV code (see \cite{2003A&A...401.1063A} and references therein) to calculate its impact on the integrated photometry of star clusters. Our main findings are: \begin{itemize} \item at 0 -- 40 per cent of the cluster's lifetime: the cluster colours are comparable to standard models (i.e. without preferential loss of low-mass stars) \item at 40 -- 80 per cent of the cluster's lifetime: a cluster appears too blue/young (compared to standard models) due to the loss of lower main-sequence stars \item at 80 -- 100 per cent of the cluster's lifetime: a cluster appears too red/old (compared to standard models) due to the loss of main sequence turn-off stars \item when interpreting photometry of mass-segregated clusters, that have preferentially lost low-mass stars, with standard photometric models (without taking mass-segregation effects into account) the derived ages can be wrong by 0.3 -- 0.5 dex \end{itemize} | 7 | 10 | 0710.0795 |
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0710 | 0710.0089_arXiv.txt | {The Very Energetic Radiation Imaging Telescope Array System (VERITAS) is an array of four 12m diameter Imaging Atmospheric Cherenkov Technique (IACT) telescopes operated at the base of Mt. Hopkins in southern Arizona. The four-telescope experiment started operation in April, 2007. GeV and TeV gamma-ray observations of blazars can be used to probe the structure and composition of their jets, and to contribute to our understanding of how supermassive black holes accrete matter. In this contribution, we present first VERITAS blazar results obtained with three and four telescopes.} \begin{document} | The EGRET (Energetic Gamma Ray Experiment Telescope) detector on board the Compton Gamma-Ray Observatory discovered strong MeV $\gamma$-ray emission from 66 Active Galactic Nuclei (AGNs), mainly from Flat Spectrum Radio Quasars and Flat Spectrum Radio Sources \cite{1999yCat..21230079H}. As of writing these proceedings in May 2007, ground-based Cherenkov telescopes have discovered TeV $\gamma$-ray emission from 17 AGB \cite{2004NewAR..48..367K}. Sixteen of the 17 sources are blazars and one is the radio galaxy M 87 \cite{2006Sci...314.1424A,2003A&A...403L...1A}. The blazars are mainly high energy peaked BL Lac objects, with BL Lac itself (an intermediate peaked BL Lac) being the only exception \cite{Albe:07}. The redshifts of the TeV $\gamma$-ray sources range from $z\,=$ 0.031 for Mrk 421 \cite{1992Natur.358..477P} to $z\,=$ 0.188 for 1ES 0347-121 \cite{2006Natur.440.1018A}. In this contribution, we will give an overview of the blazar observations performed with the VERITAS experiment. VERITAS is an array of four 12~m diameter Cherenkov telescopes located at an altitude of 1268 m above sea level on Mt. Hopkins, Az (31$^{\circ}$ 40' 30.21" N, 110$^{\circ}$ 57' 07.77" W) \cite{2002APh....17..221W}. The experiment started operation with two telescopes in spring 2006, and with four telescopes in winter 2006. The telescope system achieves an angular resolution of 0.16$^{\circ}$ and a 250 GeV-1 TeV $\nu F_{\nu}$ sensitivity of 10$^{-12}$ ergs cm$^{-2}$ s$^{-1}$ for 10 hours of integration. For a detailed description of the status and performance of the telescope system the reader is referred to the contributions of Meier et al. \cite{Meie:07} and Celik et al. \cite{Celi:07} in this volume. The most important blazar detections are described in dedicated contributions, please see Fortin et al. \cite{Fort:07} for the 1ES 1218+304 results, Cogan et al. \cite{Coga:07} for the 1ES 0806+524 and 1ES 647+250 result, Fegan et al. for Mrk 421 and Mrk 501 results \cite{Fega:07}, and Colin et al. \cite{Coli:07} for the M 87 results. The Whipple 10 m Cherenkov telescope is used to monitor blazars on a regular basis. The results of the observations taken in 2006 are described by Steele et al. \cite{Stee:07}. \begin{table}[t] \begin{center} \begin{tabular}{p{1.4cm}p{1.4cm}p{3.1cm}} \hline Observatory & Wavelength & Contact\\ \hline { Owen V.} & Radio & A. Readhead\\ Metsahovi & Radio & A. Lahteenmaki\\ WEBT & Radio/IR/O & M. Villata\\ Abastumani & Opt. & O. Kurtanidze\\ Antipodal & Opt. & J. Buckley\\ Bell & Opt. & M. Carini\\ Boltwood & Opt. & P. Boltwood\\ Bordeaux & Opt. & P. Charlot\\ Tuorla & Opt. & A. Sillanpaa\\ WIYN 0.9m & Opt. & T. Montaruli\\ Swift & X-ray & H. Krimm\\ AGILE & $\gamma$-ray & M. Tavani\\ GLAST & $\gamma$-ray & J. McEnery\\ MAGIC & $\gamma$-ray & D. Mazin\\ H.E.S.S. & $\gamma$-ray & S. Wagner\\ IceCUBE & Neutrino & T. Montaruli\\ \hline \end{tabular} \caption{List of VERITAS multiwavelength collaborators. Only one contact person is given for each observatory.} \end{center} \end{table} | The VERITAS AGN program is fully underway. The program includes intensive multiwavelength observations of blazars in a flaring state, deep observations of blazars to determine their energy spectra with high accuracy, and the search for TeV $\gamma$-ray emission from a wide range of different types of blazars. The VERITAS collaboration is working together with a large number of observers to sample the spectral energy distribution of blazars along the entire electromagnetic energy spectrum, and to obtain complementary information through the detection of high-energy neutrinos. Agreements have been reached to assure a fruitful collaboration between the three Cherenkov telescope experiments VERITAS, MAGIC, and H.E.S.S.. The first observations have resulted in the highly significant detection of the blazars Mrk 421, Mrk 501, 1ES 1218+304, and M 87. In this contribution, we have described 3-telescope observations of a number of HBL, IBL, and FSRQs. The flux upper limits are between 2.2\% and 8.6\% of the flux from the Crab Nebula. We anticipate exciting results with the full VERITAS system of 4 telescopes. | 7 | 10 | 0710.0089 |
0710 | 0710.4157.txt | We report new L and T dwarfs found in a cross-match of the SDSS Data Release 1 and 2MASS. Our simultaneous search of the two databases effectively allows us to relax the criteria for object detection in either survey and to explore the combined databases to a greater completeness level. We find two new T dwarfs in addition to the 13 already known in the SDSS DR1 footprint. We also identify 22 new candidate and bona-fide L dwarfs, including a new young L2 dwarf and a peculiar L2 dwarf with unusually blue near-IR colors---potentially the result of mildly sub-solar metallicity. These discoveries underscore the utility of simultaneous database cross-correlation in searching for rare objects. Our cross-match completes the census of T dwarfs within the joint SDSS and 2MASS flux limits to the $ \approx$97\% level. Hence, we are able to accurately infer the space density of T dwarfs. We employ Monte Carlo tools to simulate the observed population of SDSS DR1 T dwarfs with 2MASS counterparts and find that the space density of T0--T8 dwarf systems is $0.0070_{-0.0030}^{+0.0032} $~pc$^{-3}$ (95\% confidence interval), i.e., about one per 140~pc$^3$. Compared to predictions for the T dwarf space density that depend on various assumptions for the sub-stellar mass function, this result is most consistent with models that assume a flat sub-stellar mass function $d N/d M \propto M^ {0.0}$. No $>$T8 dwarfs were discovered in the present cross-match, though less than one was expected in the limited area (2099~deg$^2$) of SDSS DR1. | Our knowledge of the properties of ultra-cool L and T dwarfs has increased dramatically over the past decade as a result of the completion of several large-area optical and near-IR imaging surveys, and the implementation of fast computerized access to survey databases. L and T dwarfs are readily identified in imaging surveys by their characteristic red optical minus near-infrared (near-IR) colors. There are now hundreds of L dwarfs and over 100 T dwarfs known\footnote{A database of known L and T dwarfs is maintained at http://DwarfArchives.org \citep{kirkpatrick03, gelino_etal04}.}, the vast majority of which have been found in the Two-Micron All-Sky Survey \citep[2MASS;][]{skrutskie_etal06} and in the Sloan Digital Sky Survey \citep[SDSS;][]{stoughton_etal02}. The large number of L and T dwarfs identified in these two uniform and well-characterized data sets allows detailed investigations of the population properties of sub-stellar objects, namely their mass and luminosity functions and their multiplicity. A detailed investigation focusing on a flux-limited sample of field L dwarfs has already been presented in \citet{cruz_etal07}. However, a similarly comprehensive empirical investigation of field T dwarfs has not been performed yet. The most detailed study of T dwarfs to date is the 2MASS T5--T8 dwarf survey of \citet{burgasser02}. \citeauthor{burgasser02}'s focus on the T5-- T8 sub-range was driven by their characteristic {\sl blue} near-IR colors ($J-K_S\sim0$~mag) that set them apart from the majority of main-sequence stars in 2MASS. T0--T4 dwarfs, on the other hand, have red to neutral near-IR colors (2.0~mag~$\gtrsim J-K_S \gtrsim0.5$~mag) and searches for them face a vast contamination by background low-mass stars. As a result, our understanding of the field T0--T4 population has lagged. Although a number of T0--T4 dwarfs have been identified in the optical in SDSS \citep[][and references therein]{geballe_etal02, knapp_etal04, chiu_etal06}, an adequate analysis of the population of early T dwarfs is still lacking. Accurate knowledge of the number density of early T dwarfs relative to those of late L and mid T dwarfs is important for studies aimed at constraining the time scale of dust sedimentation and cloud formation in sub-stellar photospheres at the L/T transition. Completing the census of known T dwarfs to allow such studies is the primary science motivation for the present study. With hundreds of L and T dwarfs now known, a small number of peculiar L and T dwarfs have also emerged from the larger sample. These unusual and rare objects are set apart from their counterparts either by having abnormal surface gravities \citep[e.g.,][]{kirkpatrick_etal06, burgasser_etal06, cruz_etal07} or lower metallicities \citep{burgasser_etal03b, burgasser04b}. The recognition of such variety among the known L and T dwarfs has revealed a necessity for dimensional expansion of the present L and T dwarf classification schemes to include the effects of surface gravity and metallicity \citep{kirkpatrick05}. However, the number of known peculiar objects is presently too small to enable their accurate characterization; a larger sample will be needed to adequately anchor an expanded classification scheme. The defining photometric characteristics of peculiar ultra-cool dwarfs that set them apart from the normal population, e.g., redder near-IR colors for young L dwarfs \citep {kirkpatrick_etal06}, or redder optical and bluer near-IR colors for metal-poor ultra-cool sub-dwarfs \citep{lepine_etal03b, burgasser_etal03b, cruz_etal07}, are only now being recognized. Targeted photometric searches for such objects in the existing databases may be more fruitful in the near future. As a by-product of the present study, we remark on the characteristics of two peculiar L dwarfs discovered in our search. Finally, the analysis of the late-T dwarf population of \citeauthor{burgasser02} (\citeyear{burgasser02}; see also \citealt{burgasser04, burgasser07}, \citealt{allen_etal05}) has shown that the number density of sub-stellar objects monotonically increases until the cool end of the present spectral type sequence (at T8; $T_{\rm eff}\approx750$~K), and is expected to continue increasing for even cooler objects. That is, brown dwarfs with spectral types $>$T8 are likely numerous, but have eluded detection in present large-area surveys because of being intrinsically faint. The photospheres of extremely cool brown dwarfs, with effective temperatures below 400~K, are expected to have undergone a chemical transformation that is similar to the one occurring at the transition between the L and T spectral types, with the dominant source of opacity in the near-IR becoming water clouds, as opposed to methane clouds \citep{burrows_etal03}. Even cooler ($\lesssim$200~K) brown dwarfs may have ammonia- dominated photospheres that are very similar to those of giant planets in the Solar System. At very low effective temperatures, the emergent spectral energy distribution (SED) may be such that these objects may require a new spectral type (``Y'') for classification. The discovery and characterization of such extremely cool brown dwarfs are among the primary science drivers for present and future deep large area surveys, e.g., with UKIRT \citep[The UKIRT Infrared Deep Sky Survey;][]{lawrence_etal06}, with the Panoramic Survey Telescope and Rapid Response System \citep[Pan-STARRS;][] {kaiser_etal02}, or with the Wide-field Infrared Survey Explorer \citep[WISE;][]{mainzer_etal06}. These large sensitive projects will undoubtedly dramatically expand our knowledge of sub-stellar objects at the bottom of the main sequence. Nevertheless, it is possible that a small population of such extremely cool objects may already be present in the current generation of sky surveys. Among the existing surveys, SDSS and 2MASS offer the best chance for finding brown dwarfs later than spectral type T8 because they cover the most volume. Given their anticipated faintness, very red optical colors, and potentially blue near-IR colors, $>$T8 dwarfs may be present only as low signal-to-noise ($S/N$) single-band detections in SDSS (at $z$) and 2MASS (at $J$). As such, they are more likely to have been overlooked or flagged as artifacts in either survey. A combined consideration of the optical and near-IR data from SDSS and 2MASS may improve the chance for their discovery. That is, a cross- correlation of the SDSS and 2MASS databases may allow us to not only probe deeper, but also cooler, than is possible in either survey alone. Such a cross-correlation is the underlying approach of the present work. The ability to cross-correlate large astronomical databases is one of the main technological goals of the National Virtual Observatory (NVO). In this paper we present results from a pilot project to test an implementation of this approach, focusing on the search of new brown dwarfs from a rapid cross- match of the 2MASS All-Sky Point Source Catalog (PSC) and SDSS Data Release 1 (DR1). The project was selected by the NVO as one of three demonstration research projects that would inform of the long-term hardware and software technology needs of the NVO. The brown dwarf project in particular was aimed at identifying the technologies that will be needed to cross-match source catalogs at scale. In the present paper we describe the implementation of our cross-matching technique (\S~\ref{sec_xmatch}), and report first results from the project, including identifications of two previously overlooked T dwarfs, a new peculiar L dwarf, and a young L dwarf in SDSS DR1 and 2MASS (\S~\ref{sec_results}). We demonstrate that our dual-database cross-correlation search is more sensitive to T dwarfs than previous searches performed on SDSS or 2MASS alone, and take advantage of the high degree of completeness to T dwarfs attained in our search to estimate the T dwarf space density in the solar neighborhood (\S~\ref{sec_density}). We discuss reasons for the omission of the newly identified T dwarfs in previous SDSS and 2MASS searches and draw lessons from our experience in cross-correlating large imaging databases in \S~\ref{sec_discussion}. Finally, we outline the improved prospects for finding brown dwarfs cooler than spectral type T8 in a future iteration of the SDSS/2MASS cross-match using the much expanded Fifth Data Release (DR5) of the SDSS imaging survey (\S~\ref{sec_conclusion}). | \label{sec_conclusion} Our pilot project to search for previously overlooked T dwarfs in 2MASS and SDSS DR1 demonstrates the feasibility and utility of large database cross-correlation in discovering rare interesting objects. Our simultaneous positional and color cross-match of the 2MASS and SDSS DR1 databases uncovered 2 more T dwarfs in addition to the 13 already known in the SDSS DR1 footprint. Despite the great scrutiny with which this area has already been explored for T dwarfs, both of the new T dwarfs had previously been overlooked, probably because of suspect photometry flags in SDSS. The discovery of the two new T dwarfs demonstrates the superior sensitivity to ultra-cool dwarfs that can be attained by simultaneously cross-correlating large optical and near-IR databases, compared to searches based on individual optical or near-IR databases alone. As a by-product of our search, which focused on objects with very red optical minus near-IR colors, we also report the discovery of two new peculiar L dwarfs: an L2 dwarf with unusually blue near-IR colors, potentially linked to mildly sub-solar metallicity, and another young L2 dwarf. We took advantage of the high degree of completeness attained through our approach to obtain a flux- limited estimate of the local T dwarf space density. We used Monte Carlo analysis to reproduce the observed T dwarf population in the overlap area of SDSS DR1 and 2MASS, and found that the local space density of T dwarfs is $0.0070_{-0.0030}^{+0.0032}$~pc$^{-3}$ (95\% confidence interval), i.e., about one per 140 pc$^3$. This empirical result is the first empirical estimate of the number density of T dwarfs over the entire range of T0--T8 spectral type range and extends earlier work by \citet {burgasser02} that focused on T5--T8 dwarfs. In the context of various predictions for the local sub- stellar population \citep{burgasser04, burgasser07, allen_etal05}, we find that our result is most consistent with model-dependent estimates that assume a flat sub-stellar mass function, $d N/d M \propto M^{0.0}$. Given the success of the 2MASS/SDSS-DR1 cross-match, we expect that the approach will be instrumental for the identification of brown dwarfs cooler than the coolest ones presently known, with spectral types $>$T8. While no such brown dwarfs were identified in the present cross-match covering the 2099~deg$^2$ area of SDSS DR1, we anticipate with a 86\% probability that at least one T9 dwarf will be detectable in a similar cross-comparison of the entire 8000~deg$^2$ SDSS DR5 footprint with 2MASS. | 7 | 10 | 0710.4157 |
0710 | 0710.2335_arXiv.txt | Gravitational waveforms from the inspiral and ring-down stages of the binary black hole coalescences can be modelled accurately by approximation/perturbation techniques in general relativity. Recent progress in numerical relativity has enabled us to model also the non-perturbative merger phase of the binary black-hole coalescence problem. This enables us to \emph{coherently} search for all three stages of the coalescence of non-spinning binary black holes using a single template bank. Taking our motivation from these results, we propose a family of template waveforms which can model the inspiral, merger, and ring-down stages of the coalescence of non-spinning binary black holes that follow quasi-circular inspiral. This two-dimensional template family is explicitly parametrized by the physical parameters of the binary. We show that the template family is not only \emph{effectual} in detecting the signals from black hole coalescences, but also \emph{faithful} in estimating the parameters of the binary. We compare the sensitivity of a search (in the context of different ground-based interferometers) using all three stages of the black hole coalescence with other template-based searches which look for individual stages separately. We find that the proposed search is significantly more sensitive than other template-based searches for a substantial mass-range, potentially bringing about remarkable improvement in the event-rate of ground-based interferometers. As part of this work, we also prescribe a general procedure to construct interpolated template banks using non-spinning black hole waveforms produced by numerical relativity. | \label{sec:intro} A network of ground based gravitational-wave (GW) detectors (LIGO, Virgo, GEO~600, TAMA) is currently collecting data, which a world-wide scientific collaboration is involved in analyzing. Among the most promising sources detectable by these observatories are coalescing compact binaries consisting of black holes (BHs) and/or neutron stars spiraling toward each other as they lose orbital energy and angular momentum through gravitational-wave emission. The gravitational-wave signal from coalescing binaries is conventionally split into three parts: inspiral, merger and ring down. In the first stage, the two compact objects, usually treated as point masses, move in quasi-circular orbits (eccentricity, if present initially, is quickly radiated away). This part of the waveform is described very well by the post-Newtonian (PN) approximation of general relativity. In this approximation the Einstein equations are solved in the near zone (which contains the source) using an expansion in terms of the (small) velocity of the point masses. In the far zone, the vacuum equations are solved assuming weak gravitational fields, and these two solutions are matched in the intermediate region~\cite{fock,Blanchet:1985sp,Wagoner:1976am}. The PN approximation breaks down as the two compact objects approach the ultra-relativistic regime and eventually merge with each other. Although various resummation methods, such as Pad\'e~\cite{DIS98} and effective-one-body (EOB) approaches~\cite{Buonanno:1998gg}, have been developed to extend the validity of the PN approximation, unambiguous waveforms in the merger stage must be calculated numerically in full general relativity. Recent breakthroughs in numerical relativity \cite{Pretorius:2005gq,Campanelli:2005dd,Baker05a} have allowed many groups \cite{Pretorius:2005gq,Campanelli:2005dd,Baker05a,Herrmann2006, Sperhake2006,Bruegmann:2006at,Thornburg-etal-2007a,Etienne:2007hr} to evolve BH binaries fully numerically for the last several orbits through the plunge to single BH formation. The field is now rapidly developing the capability to routinely evolve generic black-hole binary configurations in the comparable-mass regime, and to accurately extract the gravitational-wave signal. Important milestones include simulations of unequal-mass binaries and calculations of the gravitational recoil effect and the evolution of black-hole binaries with spin \cite{Baker2006b,Gonzalez06tr,Campanelli2006b,Campanelli2006c,Herrmann:2007ac, Koppitz-etal-2007aa,Gonzalez:2007hi,Campanelli2007,Campanelli:2006fy, Pollney:2007ss,Rezzolla-etal-2007}. Comparisons with post-Newtonian results are essential for data analysis efforts, and several groups have published results showing good agreement of various aspects of non-spinning simulations with post-Newtonian predictions (see e.g.~\cite{Baker:2006yw,Baker:2006ha,Buonanno:2006ui,Berti:2007fi, Pan:2007nw,Ajith:2007qp,Hannam:2007ik,Boyle:2007ft}), and first results for certain configurations with spin have also become available \cite{Brugmann:2007zj,Lousto:2007db}. In order to overcome phase inaccuracies in long evolutions, significant progress has been made by the Caltech-Cornell group using spectral codes \cite{Scheel-etal-2006:dual-frame,Pfeiffer:2007}, and by the Jena group using higher (sixth) order finite differencing \cite{Husa2007a}. Methods to reduce the eccentricity to around $10^{-3}$ (so far only for equal-mass binaries) have been presented by the Caltech-Cornell group \cite{Pfeiffer:2007}, and the Jena group \cite{Husa:2007ec} (using initial parameters from PN solutions that take into account radiation reaction). Current numerical waveforms can be generated for the last ($\lesssim 10$) orbits, and these waveforms can be joined continuously with analytic PN inspiral waveforms to obtain one full signal. This was done in \cite{Buonanno:2006ui,Pan:2007nw,Ajith:2007qp,Buonanno:2007pf}. Indeed, there are no fundamental obstructions to generating the whole waveform, including long inspiral over hundreds of orbits, by solving the full Einstein equations numerically. But, not only would this be computationally prohibitive with current methods, it is also unnecessary: the PN formalism is known to work very well in the weak-field regime (when the BHs are well-separated), and is a low-cost and perfectly adequate substitute to fully general relativistic solutions in that regime. The numerically generated part of the gravitational-wave signal from coalescing binaries also includes the final stage of the coalescence, when a single perturbed black hole is formed and it rapidly loses its deviations from a Kerr black hole via gravitational waves. This part of the signal can be decomposed as a superposition of exponentially damped modes, and is called quasi-normal mode `ring down', by analogy with the vibrations of a bell. The detectable part of the ring down is rather short and only a few modes (if not only the dominant one) are expected to be important/detectable by initial ground-based observatories. This will not be true, however, for the advanced detectors \cite{Berti:2007zu} and certainly it is not the case for LISA, the planned space-borne gravitational-wave observatory. Indeed, the majority of the signal-to-noise ratio (SNR) comes from the quasi-normal mode ringing of binary systems with a total mass above a few $10^6\, M_{\odot}$ \cite{Berti:2005ys}. For LISA, and also perhaps for the next generation of ground based detectors, it will be possible to detect several quasi-normal modes and test the `no hair' theorem, according to which all modes are functions of a BH's mass and spin~\cite{Dreyer:2003bv,Berti:2005qd, Berti:2005ys}. Joining analytically modeled inspiral with numerically generated merger and ring down allows us to produce the complete gravitational-wave signal from coalescing binaries, and to use it in the analysis of detector data. There are several benefits to using the whole signal in searches. The most obvious one is the increase in SNR in a fully coherent matched filtering search~\cite{FlanHugh98,BuonD00, DIS01,Ajith:2007qp}. Increase in SNR implies increase in the event rate and improvement in the parameter estimation. Including the inspiral, merger and ring down parts in a template waveform also means that the waveform has a more complex structure. This extra complexity will also bring about some improvement in the parameter estimation~\cite{ParamEstim} and possibly also a reduction in the false alarm rate in analysis of the data from the ground-based network of detectors. This is because it is in general harder for the noise to mimic a complex signal\footnote{At least we expect this to happen for those binaries for which both the inspiral and the merger contribute significantly to SNR.}. For LISA, the detection of inspiralling super-massive black holes is not a problem; the SNR is expected to be so large that we expect some signals to be visible by eye in LISA data. However, using the full signal for LISA data analysis is equally important because the full signal is essential in estimating parameters of the binary with the required accuracy. This is important not only from the astrophysical point of view, but also because we need to subtract loud signals from the data in order to detect/analyze other signals. Imperfect signal removal due to errors in the parameter estimation will result in large residuals and will adversely affect subsequent analyses. Improved parameter estimation will also enable GW observations (in conjunction with electromagnetic observations) to constrain important cosmological parameters, most importantly the equation of state of dark energy~\cite{Schutz86,Markovic:1993cr,ChernoffFinn:1993,HolzHugh05, Arun:2007hu,ParamEstim}. The numerical waveforms described above are still computationally expensive and cannot be used directly to densely cover the parameter space of the binary BHs that will be searched over by matched filtering techniques. A promising alternative is to use the post-Newtonian and numerical-relativity waveforms to construct an analytic model that sufficiently accurately mimics a true signal \cite{Ajith:2007qp, Buonanno:2007pf}. In~\cite{Ajith:2007qp} we have suggested a phenomenological family of waveforms which can match physical signals from non-spinning binaries in quasi-circular orbits with fitting factors above 99\%. In this paper we extend this formulation to propose a two-parameter family of template waveforms which are explicitly parametrized by the physical parameters of the binary. We show that this two-dimensional template family is not only \textit{`effectual'} in detecting the signals from binary BH coalescences, but also \textit{`faithful'} in estimating the parameters of the binary. This family of template waveforms can be used to densely cover the parameter space of the binary, thus avoiding the computational burden of generating numerical waveforms in each grid point in the parameter space. We compute the effectualness and faithfulness (see Section~\ref{sec:DA} for definitions) of the template family in the context of three different ground-based detectors: namely, Initial LIGO, Virgo and Advanced LIGO. We also compare the sensitivity of a search which coherently includes all three (inspiral, merger and ring down) stages of the BH coalescence with other template-based searches which look for each stage separately. Our `target signals' are constructed by matching the numerical-relativity waveforms to a particular family (\emph{TaylorT1} approximant~\cite{DIS01}) of post-Newtonian waveforms, but this choice is by no means necessary. Indeed, we expect that more robust ways of constructing post-Newtonian approximants, such as the effective one-body approach~\cite{Buonanno:1998gg} or Pad\'e resummation approach~\cite{DIS98}, will give better agreement with numerical-relativity (NR) waveforms. But the purpose of the current paper is to explicitly prescribe a general procedure to produce hybrid and phenomenological waveforms, and to construct interpolated template banks using parametrized waveforms. We show that, given the number of numerical wave cycles we employ, even a simple PN choice like TaylorT1 leads to very faithful and effectual templates, and significantly increases the possible range of gravitational-wave searches. The use of improved PN approximants will require a smaller number of NR cycles, thereby further reducing computational cost for template construction. There are also other approaches for comparing analytic and numerical waveforms and for constructing hybrid waveforms (see, for example \cite{Pan:2007nw}); it would be interesting to compare the results presented in this work with other approaches presented in the literature. The paper is structured as follows. In Section~\ref{sec:numrelintro} we summarize the methods of current numerical-relativity simulations, including a setup of the initial data that allows an unambiguous comparison with post-Newtonian results, and the wave extraction techniques. In Section~\ref{sec:DA} we briefly outline the waveform generation using the \emph{restricted} post-Newtonian approximation. There we briefly introduce the main data-analysis techniques and define notations that are used in the subsequent sections. In Section~\ref{sec:phenomtemplates} we construct a phenomenological template family parametrized only by the masses of the two individual black holes. First we combine restricted 3.5PN waveforms \cite{BDEI04} with results from NR simulations to construct `hybrid' waveforms for the quasi-circular inspiral of non-spinning binaries with possibly unequal masses. Then, we introduce a phenomenological family of templates constructed in the frequency domain. Initially the template family is parametrized by 10 phenomenological parameters. We then find a unique mapping of these 10 parameters to the two physical parameters: namely, the total mass $M$ and the symmetric mass ratio $\eta \equiv M_1 M_2 / M^2$, so that the template family is just two-dimensional. The resulting templates have remarkably high fitting factors with target waveforms. Here we also compute the faithfulness of the templates and the bias in the estimation of the parameter of the binary. A comparison of the sensitivity of the search using the proposed template family with other existing template-based searches is also presented. Finally, we summarize our main results in Section~\ref{sec:summary}. Some details of the calculations involved are described in Appendices~\ref{app:fitfactor} and~\ref{app:horDist}. We adopt geometrical units throughout this paper: $G=c=1$. \label{sec:DA} In this Section we will introduce notation that will be used later in the paper and describe briefly the main data-analysis techniques currently used in gravitational wave astronomy. \subsection{Restricted post-Newtonian waveforms} We use the restricted PN waveform at mass-quadrupole order, which has a phase equal to twice the orbital phase up to highest available order in the adiabatic approximation, and amplitude accurate up to leading order. The corresponding $\hc$ is given by \begin{equation} \hc=\frac{\eta M}{r}v^2(t) e^{2\rmi\phi}\left[(1+\cos\iota)^2 e^{-\rmi \varphi(t)} + (1-\cos\iota)^2 e^{\rmi \varphi(t)}\right] \label{eq:PNhOptOrient} \end{equation} where $M \equiv M_1+ M_2$ is the total mass, $\eta \equiv M_1 M_2/M^2$ is the symmetric mass ratio, $r$ is the observation radius, $\iota$ is the inclination angle; the quantity $v(t)$ is an expansion parameter, defined by $v=(M \dot{\varphi}/2)^{1/3}$ with $\varphi(t)$ equal to twice the adiabatic orbital phase. The {\it waveform} seen by the detector is given by \be s(t) = 4\,\eta\, \frac{M}{r} A \, v^2(t) \cos[\varphi(t)+\varphi_0], \ee where, for short-lived signals (i.e., with duration much shorter than the earth rotation time, as well as de-phasing time scale due to Doppler shifts induced by earth motion and rotation), $A$ and $\varphi_0$ are numerical constants depending on the relative position and orientation of the source relative to the detector, as well as the antenna pattern functions of the detector. In PN theory, the adiabatic phase $\varphi(t)$ is determined by the following ordinary differential equations (also called the \emph{phasing formula}): \begin{equation} \frac{{\rm d}\varphi}{\dt} = \frac{2v^3}{M},\ \ \ \ \frac{{\rm d}v}{\dt} = -\frac{{\cal F}(v)}{ME'(v)}. \label{eq:phasing1} \end{equation} In these expressions, $E'(v)={\rm d}E(v)/{\rm d}v$ where $E(v)$ is the binding energy (per unit mass) of the system, and ${\cal F}(v)$ is the GW luminosity. $E(v)$ and ${\cal F}(v)$ are computed as post-Newtonian expansions in terms of $v$~\cite{Blanchet:LivRev}. Currently, the binding energy function $E(v)$ has been calculated to $v^6$ (3PN) accuracy by a variety of methods~\cite{DJS2000, DJSequiv,BF00,BFeom,ABF01,BDE04,itoh1,itoh2}. The flux function ${\cal F}(v),$ on the other hand, has been calculated to $v^7$ (3.5PN) accuracy \cite{BFIJ02, BDEI04} up to now only by the multipolar-post-Minkowskian method and matching to a post-Newtonian source~\cite{Blanchet:LivRev}. The inspiralling phase is usually pushed up to the point where the adiabatic evolution of circular orbits breaks down due to the lack of further stable circular orbits. In the test-mass limit, the last (or innermost) stable circular orbit (ISCO) can be computed exactly (at 6$M$ in Schwarzschild coordinates). For comparable-mass binaries, on the other hand, the ISCO cannot always arise unambiguously from PN theories. In adiabatic models, the maximum-binding-energy condition (referred to as MECO, or the maximum binding energy circular orbit,~\cite{Blanchet:2002}) can be used in place of the ISCO. This condition is reached when the derivative of the orbital binding energy with respect to orbital frequency vanishes. As a consequence, in this paper, the waveforms are evolved in time up to MECO: $E'(v)=0$. It may be noted that the ISCO and MECO may not be physically meaningful beyond the test-mass limit, but they make convenient cutoff criteria. The appropriate region of validity of PN waveforms can only be determined by comparison with fully general relativistic results, such as the numerical simulations that we discussed earlier. Given $E(v)$ and ${\cal F}(v)$, one can construct different, but equivalent in terms of accuracy, approximations to the phasing by choosing to retain the involved functions or to re-expand them. Indeed, the different PN models which describe the GW signal from inspiralling binaries agree with each other in the early stages of inspiral; but start to deviate in the late inspiral. The classification and explicit form of various models is nicely summarized in~\cite{DIS01}. In this paper we use PN waveforms obtained by numerically solving Eqs. (\ref{eq:phasing1}), called the {\it TaylorT1} approximant, to construct the `hybrid waveforms' (see Section~\ref{sec:Matching}). \subsection{Introduction to matched filtering} Since we can model the signal reasonably well, it is natural to employ matched filtering (which is the optimal detection strategy for a signal of known shape in the stationary Gaussian noise) to search for the gravitational-wave signal. Suppose the detector's data $x(t)$ contains noise $n(t)$, and possible signal $s(t)$, i.e., $x(t) = n(t) + s(t)$. Assuming $n$ to be stationary Gaussian noise, it is convenient to work in the Fourier domain, because the statistical property of the noise is completely characterized by its power spectral density $S_n(f)$, which is given by (here we use a {\it single-sided} spectrum) \begin{equation} \langle\tilde{n}(f)\tilde{n}^*(f')\rangle = \frac1{2}S_n(f)\,\delta(f-f')\,, \end{equation} where $\tilde n(f)$ is the Fourier Transform of $n(t)$ \begin{equation} \tilde{n}(f) \equiv \int_{-\infty}^{\infty}n(t) e^{-2\pi \rmi ft}\;\dt\,, \end{equation} and $\langle\ldots\rangle$ denotes taking the expectation value. Based on the detector noise spectrum, we introduce a Hermitian inner product: \begin{equation} (g | h) \equiv 2 \int_0^{\infty} \frac{\tilde{g}^*(f)\tilde{h}(f) + \tilde{g}(f)\tilde{h}^*(f)}{S_n(f)}\;\df\,. \end{equation} For the data $x$ with known signal $s$, the optimal detection statistic is given by applying a template $h$ with the same shape as $s$, or $h = \alpha s$: \begin{equation} \rho_{\rm opt} \equiv ( x | h)\,. \end{equation} The detectability of the signal is then determined by the SNR of $\rho_{\rm opt}$, \begin{equation} \frac{S}{N} = \left. \frac{(s|h)}{\sqrt{\langle(h|n)(n|h)\rangle}}\right|_{h=\alpha s} = (s|s)^{1/2}. \end{equation} (Note that the SNR does not depend on the overall normalization of $h$.) In case the template $h$ is not exactly of the same shape as $s$, the SNR will be reduced to \begin{equation} \frac{S}{N} = (s|s)^{1/2}\mathcal{M}\,, \end{equation} where $\mathcal{M}\le 1$ is the {\it match} of the template to the signal, given by \begin{equation} \mathcal{M}[s,h]\equiv \frac{(s|h)}{\sqrt{(s|s)\,(h|h)}} \equiv (\hat s|\hat h)\,, \end{equation} and where a hat denotes a normalized waveform. For more details, we refer the reader to Ref.~\cite{Cutler:1994ys}. \subsection{Template banks, effectualness and faithfulness} We now consider the more realistic problem of attempting to detect a family of waveforms $s(\vectheta)$, parametrized by a vector of physical parameters $\vectheta \in \Theta$, using a family of templates $h(\veclambda)$ parametrized by a vector of parameters $\veclambda\in\Lambda$. We first introduce the concepts of {\it physical template bank} and {\it phenomenological template bank}. Roughly speaking, physical template banks are constructed from well-motivated physical models (e.g., approximation up to a certain order)~ \cite{Babak:2006ty}, while phenomenological banks are constructed in an ad-hoc manner to mimic the desired physical signals with high accuracy. For physical banks, the vectors $\vectheta$ and $\veclambda$ consists of the same set of {\it physical parameters}, while for phenomenological banks, the vector $\veclambda$ usually contains {\it phenomenological parameters,} which can be larger or smaller in number than the physical parameters. Two phenomenological template families~\cite{BCV, BCV2} are used currently in the search for BH binaries in LIGO data~\cite{Abbott:2007xi,Abbott:2005kq}. They each represent a different motivation for introducing phenomenological banks: (i) when we have uncertainty in the signal model, we can produce a template bank with larger detection efficiency by introducing extra (phenomenological) parameters (BCV1, \cite{BCV}) so that $\textrm{dim}(\Lambda) > \textrm{dim}(\Theta)$; (ii) when the true signal depends on too many parameters and is too difficult to search over, it is sometimes possible to come up with a model with fewer (phenomenological) parameters ($\textrm{dim}(\Lambda) < \textrm{dim}(\Theta)$) and still high fitting factors (BCV2, \cite{BCV2}). The detection efficiency of a template bank towards a specific signal $s(\vectheta)$ can be measured by the threshold SNR above which the detection probability exceeds a certain minimum (usually 50$\%$), while the false-alarm probability is kept below a certain maximum (usually 1$\%$ for one-year data). The threshold value depends (logarithmically, in the case of Gaussian noise) on the number of statistically independent templates, and (inverse-proportionally) on the {\it fitting factor} (FF)~\cite{Apostolatos:1995pj}: \begin{equation} {\rm FF}[h;\vectheta] \equiv \max_{\veclambda}\mathcal{M}[s(\vectheta),h(\veclambda)] \equiv \mathcal{M}[s(\vectheta),h(\veclambda_{\rm max})] \,. \end{equation} A bank with high FF is said to be \emph{effectual}~\cite{DIS98,DIJS03}. Typically, we require that the total mismatch between the template and true signal (including the effects of both the fitting factor and the discreteness of the template bank) to not exceed $3\%$. We shall see that this requirement is easily met by our template bank. \begin{figure*} \begin{center} \psfrag{Pmap}{$P$} \psfrag{PMap2D}{$P_{2\rm D}$} \psfrag{Pt}{$P(\Theta)$} \psfrag{Pint}{$P_{\rm int}(\Theta)$} \psfrag{BT}{$\vectheta$} \psfrag{BLM}{$\veclambda_{\rm max}$} \psfrag{BLMP}{$\veclambda_{\rm max^\prime}$} \psfrag{BLMInt}{$\veclambda_{\rm int}$} \psfrag{Th}{$\Theta$} \psfrag{(i)}{(i)} \psfrag{(ii)}{(ii)} \psfrag{(iii)}{(iii)} \includegraphics[width=17.5cm]{TemplSpace_1.eps} \caption{Construction of the phenomenological template bank: (i) mapping physical signals (solid curve) into a sub-manifold (dashed curve, with example templates marked by dots) of a larger-dimensional template bank (curved surface), (ii) obtaining a lower-dimensional phenomenological bank with the same number of parameters as physical parameters, through interpolation (solid curve on the curved surface, with example templates marked by triangles), and (iii) Estimating the bias of the lower-dimensional interpolated bank by mapping physical signals into the bank (with images of example signals marked by dots).} \label{fig:templSpaceCartoon} \end{center} \end{figure*} It is natural to associate every point $\vectheta$ in the physical space $\Theta$ with the best matched point $\veclambda_{\rm max}\in\Lambda$. This leads to a mapping $P:\Theta \mapsto \Lambda$ defined by \begin{equation} \label{def:P} P (\vectheta) = \veclambda_{\rm max} \,. \end{equation} This mapping will play a key role in the construction of our template bank. We will assume the mapping $P$ to be single-valued, i.e., given a target signal, the best-matched template is unique. We depict this mapping schematically in the left panel of Fig.~\ref{fig:templSpaceCartoon}. For a physical template bank with $\vectheta$ and $\veclambda$ the same set of parameters (which we use $\vectheta$ to denote), it is {\it most convenient} to identify the best-match parameter $\vectheta_{\rm max}$ as the estimation of the original parameter $\vectheta$. In general this will lead to a systematic bias \begin{equation} \Delta\vectheta = \vectheta_{\rm max}-\vectheta = P(\vectheta)-\vectheta\,. \end{equation} A bank with a small bias (as defined above) is said to be \emph{faithful}~\cite{DIS98,DIJS03}. However, if we assume no uncertainty in the true waveforms (thereby excluding the case of BCV1), then as long as $P$ is invertible, a non-faithful physical or phenomenological bank can always be {\it converted into} a faithful bank by the re-parametrization \begin{equation} \label{faithful} h_{\rm faithful}(\vectheta) \equiv h \circ P(\vectheta) \end{equation} where we have used the standard notation $h \circ P(\vectheta) :=h (P(\vectheta))$. In other words, each template $\veclambda$ in the image set of physical signals $P(\Theta)$ is labeled by physical parameters $\vectheta = P^{-1}(\veclambda)$. For this reason, we require $P$ to be invertible. It is quite conceivable that for physical banks, $P$ should be invertible, if the physical bank does not fail to describe the true waveforms too dramatically (and of course assuming the true waveform does contain independent information about the physical parameters $\vectheta$). In this way, {\it all reasonable physical banks can be made faithful}. By contrast, if for some phenomenological bank (e.g., BCV2 if we only take into account the intrinsic parameters of the bank), $P$ is a many-to-one map, with $P(\vectheta_1)=P(\vectheta_2)$ for some $\vectheta_1 \neq \vectheta_2$. Then for a physical signal with parameter $\vectheta_1$, the template bank $h_{\rm faithful}$ would achieve the same best match at both $\vectheta_1$ and $\vectheta_2$, making physical parameter determination non-unique. In this case, we can simply keep using the phenomenological bank $h(\veclambda)$; once a detection is made with $\veclambda_{\rm max}$, the a set of parameters $P^{-1} (\veclambda_{\rm max})$ would be the best knowledge we have about the physical parameters of the source. (In practice, statistical uncertainty also applies to $\veclambda_{\rm max}$.) \begin{figure*} \begin{center} \includegraphics[width=16cm]{MatchedWaves_AEIJena.eps} \caption{NR waveforms (thick/red), the `best-matched' 3.5PN waveforms (dashed/black), and the hybrid waveforms (thin/green) from three binary systems. The top panel corresponds to $\eta = 0.25$ NR waveform produced by the AEI-CCT group. The second, third and fourth panels, respectively, correspond to $\eta = 0.25, 0.22$ and $0.19$ NR waveforms from produced by the Jena group. In each case, the matching region is $-750 \leq t/M \leq -550$ and we plot the real part of the complex strain (the `+' polarization).}. \label{fig:timeDomWaveNRPNHyb} \end{center} \end{figure*} | \label{sec:summary} Making use of the recent results from numerical relativity we have proposed a phenomenological waveform family which can model the inspiral, merger and ring-down stages of the coalescence of non-spinning binary black holes in quasi-circular orbits. We first constructed a set of hybrid waveforms by matching the NR waveforms with analytical PN waveforms. Then, we constructed analytical phenomenological waveforms which approximated the hybrid waveforms. The family of phenomenological waveforms that we propose was found to have fitting factors larger than 0.99 with the hybrid waveforms. We have also shown how this phenomenological waveform family can be parametrized solely in terms of the physical parameters ($M$ and $\eta$) of the binary, so that the template bank is, in the end, two dimensional~\footnote{It may be noted that, the mapping from the phenomenological to physical parameters might not be unique in the case of spinning binaries, because of the degeneracies of different spin configurations.}. This two dimensional template family can be explicitly expressed in terms of the physical parameters of the binary. We have estimated the `closeness' of this two-dimensional template family with the family of hybrid waveforms in the detection band of three ground-based GW detectors, namely Initial LIGO, Virgo and Advanced LIGO. We have estimated the effectualness (larger overlaps with the target signals for the purpose of detection) and faithfulness (smaller biases in the estimation of the parameters of the target signals) of the template family. Having both types of overlap always greater than 0.99, the two dimensional template family is found to be both effectual and faithful in the detection band of these ground-based detectors. This phenomenological waveform family can be used to densely cover the parameter space, avoiding the computational cost of generating numerical waveforms at every grid point in the parameter space. We have compared the sensitivity of a search using this template family with other searches. For a substantial mass-range, the search using all three stages of the binary black hole coalescence was found to be significantly more sensitive than any other template-based searches considered in this paper. This might enable us to do a more sensitive search for intermediate-mass black holes using ground-based GW detectors. A number of practical issues need to be addressed before we can employ this template family in an actual search for GW signatures. The first issue will be how to construct a bank of templates sufficiently densely spaced in the parameter space so that the loss in the event rate because of the mismatch between the signal and template is restricted to an acceptable amount (say, 10\%). The explicit frequency domain parametrization of the proposed template family makes it easier to adopt the formalism proposed by Owen~\cite{Owen:1995tm} in laying down the templates using a metric in the parameter space. Work is ongoing to compare the metric formalism adopted to the proposed template family and other ways of laying out the templates, for example a `stochastic' template bank~\cite{StochTemplBank}. Also, this explicit parametrization makes it easier to employ additional signal-based vetoes, such as the `chi-square test'~\cite{Allen:2004gu}. This will also be explored in a forthcoming work. Since this template bank is also a faithful representation of the target signals considered, we expect that, for a certain mass-range, a search which coherently includes all three stages of the binary coalescence will bring about remarkable improvement in the estimation of parameters of the binary. This may be especially important for LISA data analysis in estimating the parameters of supermassive black hole binaries. This is also being explored in an ongoing work~\cite{ParamEstim}. It is worth pointing out that the family of target signals (the hybrid waveforms) that we have considered in this paper is not unique. One can construct alternate families of hybrid waveforms by matching PN waveforms computed using different approximations with NR waveforms. Also, owing to the differences in initial data and accuracy of numerical techniques, the NR waveforms from different simulations can also be slightly different. Thus, the coefficients listed in Tables~\ref{tab:polCoeffsAmpParams} and~\ref{tab:polCoeffsPhaseParams} have a unique meaning only related to this particular family of target waveforms. But we expect that the general parametrization that we propose in this paper will hold for the whole family of non-spinning black hole coalescence waveforms from quasi-circular inspiral. As we have mentioned in the Introduction, the purpose of this paper is to explicitly prescribe a general procedure to construct interpolated template banks using parametrized waveforms which mimics actual signals from binary black hole coalescence (as predicted by numerical relativity and analytical methods). Nevertheless, it may be noted that most of the PN waveforms constructed using different approximations are known to be very close to each other (see, for example,~\cite{DIS01}). Also, we expect that NR waveforms from different simulations will converge as the accuracy of numerical simulations improves (see, for example,~\cite{Baker:2007fb}). Thus, since different families of PN and NR waveforms, which are the `ingredients' for constructing our target signals, are very close to each other, we expect that the phenomenological waveform family proposed in this paper, in its present form, will be sufficiently close to other families of target signals for the purpose of detecting these signals. As a preliminary illustration of this, we have computed the fitting factors of the template waveforms with a different family of hybrid waveforms (constructed from longer and more accurate NR waveforms), and have shown that the overlaps are indeed very high. This will be explored in detail in a forthcoming work. Also, we remind the reader that this paper consider only the leading harmonic of the GW signal ($\ell=2,\ m=\pm 2$). We expect that the contribution from the higher harmonics become important for high mass ratios, which will be investigated in a forthcoming work. | 7 | 10 | 0710.2335 |
0710 | 0710.2103_arXiv.txt | We present a new version of the fully 3D photoionization and dust radiative transfer code, {\sc mocassin}, that uses a Monte Carlo approach for the transfer of radiation. The X-ray enabled {\sc mocassin} allows a fully geometry independent description of low-density gaseous environments strongly photoionized by a radiation field extending from radio to gamma rays. The code has been thoroughly benchmarked against other established codes routinely used in the literature, using simple plane parallel models designed to test performance under standard conditions. We show the results of our benchmarking exercise and discuss applicability and limitations of the new code, which should be of guidance for future astrophysical studies with {\sc mocassin}. | \label{s:intro} Photoionized environments characterize a wide range of astrophysical problems involving sources of X--radiation. With the advent of new technology used for instruments on board of (e.g.) {\it XMM-Newton} and {\it Chandra}, high resolution spectroscopy of such environments has become a reality. Paerels et al. (2000), for instance, observed the photoionized wind in Cygnus X--3 with {\it Chandra} High Energy Transmission Grating Spectrometer (HETGS) showing the discrete emission to be excited by recombination in a tenuous X--ray--photoionized medium which is not symmetric with the source of the wind. Other examples include the detection of several recombination emission lines (from Fe~{\sc xxvi} at 1.78~{\AA} to N~{\sc vi} at 29.08~{\AA}), by Jimenez-Garate et al. (2005) in a 50~ks observation of the bright X-ray binary Hercules X-1 with {\it Chandra} HETGS. We also note the {\it Chandra} ACIS observations of the deeply eclipsing cataclysmic variable DQ Herculis by Mukai et al. (2003), who were able to pin down the origin of the soft X-rays from this system as being due to scattering of the unseen central X-ray source, probably in an accretion disk wind. A number of 1D photoionization codes, including G.~Ferland's {\sc cloudy} (Ferland et al. 1998) and T.~Kallman's {\sc xstar} (Kallman \& McCray 1980, Kallman \& Bautista 2001) continue to represent powerful analytical tools for the analysis of astrophysical spectra from the X-ray to the infrared regime. These codes are designed for diffuse, optically thin media, which may also be irradiated by a non-thermal X-ray continuum, and have been, for instance, applied to the modelling of Narrow Line Regions (NLRs) of Active Galactic Nuclei (AGNs). Several other codes have been developed which specialize in emission and reflection spectra from optically thick hot photoionized media, irradiated by a non-thermal continuum extending to the hard X-ray region, such as X--ray irradiated accretion disks (e.g. Ross \& Fabian, 1993; Nayakshin et al. 2000; Dumont et al. 2000). To date, these and most other photoionization codes have numerically solved the equations of radiative transfer (RT) under the assumption of spherical symmetry or in plane parallel geometries, whereupon the problem is reduced to a 1D calculation. While very few real X-ray sources are spherically symmetric, this approximation has been driven by the available computing power and the complexity of the multi-dimensional case. Mauche et al. (2004) developed a Monte Carlo code to investigate the radiation transfer of Ly$\alpha$, He$\alpha$, and recombination continua photons of H- and He-like C, N, O and Ne produced in the atmosphere of a relativistic black hole accretion disk. This code, however, while accounting for Compton scattering and photoabsorption followed by recombination, does not calculate the ionization state of the plasma. To our knowledge, there are currently no general, self-consistent and publicly available X-ray photoionization and dust RT codes capable of working in 3D. Although the 1D codes mentioned above are powerful tools for the analysis of the pan-chromatic spectra of numerous astrophysical environments, their application is restricted to rather simplified geometries. The computational demand of realistic 3D simulations has recently come within the reach of low-cost clusters. Taking advantage of this, the first self-consistent, 3D photoionization and dust RT code was developed for the IR-UV regime using Monte Carlo techniques (Ercolano et al. 2003a, 2005). The code, {\sc mocassin} (MOnte CArlo SimulationS of Ionized Nebulae) was designed to build realistic models of photoionized environments of arbitrary geometry and density distributions, and can simultaneously treat the dust RT. The code can also treat illumination from multiple point- or arbitrarily extended sources. The fully parallel version of {\sc mocassin} (documented and publicly available) is well-tested for classical nebulae, according to standard photoionization benchmarks (P\'equignot et al., 2001), and has been successfully applied to the modeling of H~{\sc ii} regions (e.g. Ercolano, Bastian \& Stasi\'{n}ska, 2007) and planetary nebulae (e.g. Ercolano et al. 2003b,c, 2004; Gon\c{c}alves et al., 2006; Schwarz \& Monteiro, 2006; Wright, Ercolano \& Barlow, 2006). We now present a significantly improved version of the {\sc mocassin} code (version 3.00) which extends to the X-ray regime. In the tradition of previous releases of the {\sc mocassin} code, the X-ray version will also be made publicly available to the scientific community shortly after publication of this article\footnote{Previous versions are available upon request from BE.}. In Section~\ref{s:code} we summarize the physical processes and atomic data added/modified in the new implementation and discuss the applicability and limitations of the code in its current form. In Section~\ref{s:bench} we compare our code to established 1D codes for benchmark tests, comprising emission line spectra from model NLR and from three thin low-density slab models illuminated by a hard continuum. A brief summary is given in Section~\ref{s:summary}. | \label{s:summary} We have presented a new version of the fully 3D Monte Carlo photoionization code, {\sc mocassin}. The code was extended to allow the modelling of plasma irradiated by a hard continuum spanning from radio to gamma rays. The atomic data set of the code was also significantly updated and it is now synchronized with the latest release of the {\sc chianti} database. The applicability and limitations of the new code were discussed, and the results of a thorough benchmarking exercise presented. No major problems were found by the benchmark tests, although some minor differences have been found. We have highlighted a number of significant improvements in the atomic datasets available today compared to those available at the time the original benchmark tables were compiled. We provide here updated values and we emphasize the need of a new benchmarking exercise to be undertaken by the plasma modelling community. The good performance of the {\sc mocassin} code in all benchmark tests demonstrates that it is ready for application to real astrophysical problems. The public version of the X-ray enhanced {\sc mocassin} code is available on request from the author. | 7 | 10 | 0710.2103 |
0710 | 0710.0324_arXiv.txt | {Emission lines in quasars are believed to originate from a photoionized plasma. There are, however, some emission features which appear to be collisionally excited, such as the \FeII\ multiplet bands. Shortward of \Lya, there also are a few permitted lines of species from low to intermediate ionization.} {\ton\ ($\zq=1.928$) exhibits the steepest far-UV continuum decline known ($\Fnu \propto \nu^{-5.3}$) shortward of 1050\,\AA. This object also emits unusually strong low to intermediate excitation permitted lines shortward of the Lyman limit.}{Using archive spectra of \ton\ from HST, IUE and Palomar, we measure the fluxes of all the lines present in the spectra and compare their relative intensities with those observed in composite quasar spectra.} {Our analysis reveals unusual strengths with respect to \Lya\ of the following low to intermediate excitation permitted lines: \OII+\OIII\ (835\,\AA), \NIII+\OIII\ (686--703\,\AA) and \NIII+\NIV\ (765\,\AA). We compare the observed line spectrum with both photoionization and shock models.}{Photoionization cannot reproduce the strengths of these far-UV lines. Shocks with $\Vs \simeq 100$\,\kms\ turn out to be extremely efficient emitters of these lines and are favored as excitation mechanism. } | \label{sec:intro} In this work, we analyze the emission lines of an unusual quasar, \ton, which is alternatively named PG\,1017+280 or J1019+2745 with redshift $\zq=1.928$. It is severely deficient in ionizing photons, since its Spectral Energy Distribution (\sed) shows a remarkable steepening of the continuum in the rest-frame far-UV, shortward of 1100\,\AA\ (Binette \& Krongold 2007, hereafter BK07; Binette et\,al. 2007). If the far-UV is fitted by a powerlaw ($\Fnu \propto \nu^{+\alpha}$), the index\footnote{Among the 77 quasars whose far-UV indices could be measured by Telfer et\,al. 2002, there were only 3 objects with a continuum steeper than $\nu^{-3}$.} is as steep as $\nu^{-5.3}$. BK07 suggest that the extreme-UV flux might undergo a recovery shortward of 450\,\AA. While the near-UV emission-line spectrum appears to be `normal', the far-UV spectrum shows low to intermediate ionization species with unusual strengths. Using the UV \sed\ constructed by BK07 from archive data, we will quantify this statement and present photoionization and shock models for comparison. The aim is to understand how the extreme deficiency of ionizing photons in \ton\ might be impacting the emission line spectrum. The emission-line spectrum of quasar and Seyfert\,I galaxies is generally believed to originate from gas photoionized by a nuclear UV source. State of the art photoionization models of the Broad Emission Line Region (BELR) such as those developed by Baldwin et\,al. (1995) and dubbed `locally optimally emitting clouds' (LOC) models can successfully reproduce most of the emission lines observed in quasars. A grid of such models can be found in Korista et\,al. (1997, hereafter KO97) and more recently in Casebeer et\,al. (2006 and references therein). There are, however, a few exceptions to the success of pure photoionization. In particular, photoionization models require micro-turbulences in order to reproduce the shape and intensity of the \FeII\ UV-band (Baldwin et\,al. 2004). A possible alternative is that the region producing \FeII\ is collisionally ionized, as proposed by Grandi (1981, 1982), Joly (1987), V\'eron-Cetty et\,al. (2004, 2006) and Joly et\,al. (2007). In this work, we present evidence that photoionization might not be sustainable in the case of some of the far-UV permitted lines reported in this paper. | 7 | 10 | 0710.0324 |
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0710 | 0710.3866_arXiv.txt | We present here X-ray spectra of the HMXB SMC~X-1 obtained in an observation with the \xmm{} observatory beginning before eclipse and ending near the end of eclipse. With the Reflection Grating Spectrometers (RGS) on board \xmm{}, we observe emission lines from hydrogen-like and helium-like ions of nitrogen, oxygen, neon, magnesium, and silicon. Though the resolution of the RGS is sufficient to resolve the helium-like $n$=2$\to$1 emission into three line components, only one of these components, the intercombination line, is detected in our data. The lack of flux in the forbidden lines of the helium-like triplets is explained by pumping by ultraviolet photons from the B0 star and, from this, we set an upper limit on the distance of the emitting ions from the star. The lack of observable flux in the resonance lines of the helium-like triplets indicate a lack of enhancement due to resonance line scattering and, from this, we derive a new observational constraint on the distribution of the wind in SMC~X-1 in velocity and coordinate space. We find that the solid angle subtended by the volume containing the helium-like ions at the neutron star multiplied by the velocity dispersion of the helium-like ions must be less than 4$\pi$\,steradians\,\kms{}. This constraint will be satisfied if the helium-like ions are located primarily in clumps distributed throughout the wind or in a thin layer along the surface of the B0 star. | In isolated early-type stars winds are driven as ultraviolet photons from the stellar surface impart their outward momentum to the wind in resonance line transitions. In a high-mass X-ray binary (HMXB), the wind is ionized by X-rays from the compact object. If the X-radiation is intense enough, the resulting ions will not have transitions in the ultraviolet and this greatly affects the dynamics of the wind \citep[see, e.g., ][ and references therein]{blo94}. The SMC~X-1/Sk~160 system, which consists of a 0.71 second X-ray pulsar and a B0I star together in a 3.9-day orbit, is the most X-ray luminous known HMXB and, therefore, presumably, an extreme example of wind disruption by X-ray ionization. The behavior of the wind of an early-type star under the influence of ionization from an X-ray emitting companion has been the subject of many theoretical studies. For SMC~X-1, the most relevant and detailed such study is, arguably, the one by \citet{blo95} which included a numerical hydrodynamic simulation. The features that appeared in that simulation included a regular, UV-driven wind on the X-ray shadowed side of the star, a thermal wind on the X-ray illuminated side, transition regions between the two types of winds and dense, finger-like structures in the equatorial plane. However, even this simulation includes significant approximations: the gravity of the compact object is not included and X-ray photoionization and its dynamical effects are treated in a very approximate way. Because of the complexity of these systems and the difficulty of accounting for all of the relevant physics, observations are critical to characterizing the behavior of HMXB winds. Though previous X-ray observations of SMC~X-1 have had spectral resolving powers ($\equiv\lambda/\Delta\lambda=E/\Delta{}E$) of $\sim$50 or less, X-ray spectroscopy has already revealed much about the wind. X-ray spectroscopic observations of this resolving power can, through measurements of line fluxes, indicate the quantity and ionization level of X-ray emitting material. In observations of SMC X-1 with \asca{}, emission lines were not detected and upper limits were set on the quantity of material of intermediate ionization level in the wind of SMC~X-1, excluding the presence of finger-like structures as they appeared in the simulation of \citeauthor{blo95} \citep{woj00}. Further observations of SMC~X-1 with the Advanced Camera for Imaging Spectroscopy (ACIS) on-board the {\it Chandra X-ray Observatory}, however, have detected emission lines and, therefore, the presence of material of intermediate ionization in the wind of SMC~X-1 \citep{vrt01}. These results indicate that overdense regions do exist in the wind of SMC~X-1 and that while the simulations of \citeauthor{blo95} may not be accurate in detail, the general form of the structure predicted by \citeauthor{blo95} may, in fact, be present. To develop further constraints on models of the wind in SMC~X-1, it is desirable to have constraints in addition to the quantity and ionization level of the X-ray emitting plasma. High-resolution spectroscopic observations with resolving powers $\gtrsim$200 have the potential to provide useful constraints on the kinematics of the high-ionization wind in HMXBs through direct measurements of Doppler line shifts and broadening. In addition, the $n=2\to1$ triplets of He-like ions are resolved at high resolution and measurements of the fluxes of the individual lines of these triplets provide a constraint on the structure and kinematics of the X-ray emitting material. Of the three components of the helium-like triplet, only one --- the resonance line --- has a large oscillator strength and can be enhanced by resonant line scattering. Because resonant line scattering saturates and because this saturation depends on the distribution of the scatterers in physical and velocity space, constraints on the distribution of the wind in physical and velocity space can be derived from the relative fluxes of these three lines \citep{woj03}. The flux ratios of the helium-like triplets are also affected by ultraviolet radiation \citep*{blu72} and, therefore, the flux ratios measured from an HMXB constrain the distance from the photosphere of the high-mass star to the emission region. Because of the location of SMC X-1 outside of the plane of the Galaxy, the column density of interstellar material to it is small and, unlike the HMXBs in the Galaxy, it can be observed in the wavelength range 15--35\,\AA. The $n=2\to1$ transitions of the hydrogen-like and helium-like ions of oxygen and nitrogen as well as hydrogen-like carbon lie in this range. Therefore, unlike with the Galactic HMXBs, it is possible to derive observational constraints on the density and velocity distributions of the regions of moderate ionization in SMC~X-1 where these ions exist. Because of the high spectral resolution (FWHM of $\sim$65\,m\AA, corresponding to resolving powers in the range 230--540) and large effective area (total of $\sim$80\,cm$^2$) of the Reflection Grating Spectrometer on \xmm{} in this band, the RGS on \xmm{} is particularly well-suited to spectroscopic study there. We present here observations of SMC~X-1 with \xmm{} beginning before an eclipse and ending near the end of that eclipse. We present an analysis in which we focus on the data from the RGS and derive constraints on the space and velocity distribution of material in the wind. In \S\ref{sec:empir} we describe our observations and an empirical analysis of the line emission in which we measure line fluxes, shifts, and widths. In \S\ref{sec:phys} we derive constraints on the distribution of the wind in physical and velocity space from the line fluxes we measure. In \S\ref{sec:discuss}, we discuss the implications of our results for models of HMXB winds. | \label{sec:discuss} We have observed SMC~X-1 in eclipse with \xmm{} and, with the RGS, resolved the helium-like triplets of nitrogen, oxygen, neon, and magnesium. To our knowledge, this is the first time helium-like triplets from SMC~X-1 have been resolved and the first time the helium-like triplets of nitrogen, oxygen, or neon have been resolved for any HMXB. In all cases, only one component of the triplet, the intercombination line, is observably present. The lack of observable forbidden line fluxes in the triplet is easily explained by photoexcitation pumping of the forbidden line into the intercombination line by the ultraviolet radiation of the B0 star Sk~160 and allows upper limits to be set on distance of the emitting helium-like ions from Sk~160. The absence of observable fluxes in the resonance lines is consistent with what we expect from recombination in the photoionized wind. However, the fact that the resonance lines are not enhanced by resonant scattering implies a lower limit on the optical depth of the wind in the resonance line and, therefore, constraints on the structure and kinematics of the wind are also implied. We have modeled the helium-like line emission as recombination and scattering emission from a region that is partly obscured by the companion star and is described by a single solid angle ($\Omega$) and column density ($N$) and the velocities along lines of sight from the neutron star have a boxcar distribution with width $\Delta{}v$ for each of the ions. We have inferred infer the quantity $\Omega\Delta{}N$ from the sum of the flux of the intercombination and forbidden lines which are not affected by resonant line scattering. We infer the quantity $N/\Delta{}v$ from flux of the resonance line relative to the sum of the other two (the inverse of the $G$ ratio). From the constraints on these two quantities, we also obtain a constraint on the quantity $\Omega\Delta{}v$. Because we do not detect the resonance line, we obtain only lower limits on $N/\Delta{}v$ and only upper limits on $\Omega{}\Delta{}v$. As previously mentioned, in our analysis we have assumed that photons escape isotropically along lines from the location of their production or first scattering. While we do expect recombination emission to be isotropic, resonant line scattering is not isotropic. The angular distribution for resonant line scattering is a linear combination of an isotropic distribution and an dipole distribution \citep[see, e.g.,][Ch.\ 19]{cha60}. Furthermore, for large optical depths that are comparable to or greater than unity such as we infer, photons may undergo several scatterings before escaping. In this case, photons will escape preferentially along directions where the optical depth is least. In general, the calculation of photon escape direction is also complicated by the fact that photons may travel large distances across the plasma before escaping. While it is difficult to include this possibility in analytic calculations, it is straightforward using Monte Carlo techniques. However, if the region that emits the helium-like line emission has supersonic velocity differentials, as we know that the bulk of the wind does, the transfer of resonant line photons is essentially local (the Sobolev approximation, \citealt{cas70}) and can be approached analytically. For a supersonically expanding wind, the optical depth in a direction given by the coordinate $x$ is inversely proportional to $dv_x/dx$. If the flow moves along lines from a single point and $y$ is the distance from that point, then these derivatives, in directions parallel and perpendicular to the flow, respectively, are $dv/dy$ and $v/y$. While these quantities will not generally be equal, they will generally be of the same order over the bulk of the wind and, therefore, if we assume complete redistribution in frequency and direction for individual scatterings, then the radiation will escape isotropically \citep[see][]{cas70}. In fact, individual scattering events do not redistribute frequency and direction completely and so even with $v/s=dv/ds$, radiation does not escape isotropically. However the difference from the case of complete redistribution is no greater than 10\%--20\% (\citealt{car72}, see also \citealt{mih80}). In light of these facts, we proceed to explore the implications of our results. However, it must be kept in mind that the approximations we have made for the radiative line transfer may quantitatively affect our conclusions. The implication of our results on fundamental wind parameters, such as the mass-loss rate and the terminal wind velocity, is complex owing to factors such as the complex dependence of the ion fractions on density in photoionization balance. Therefore, interpreting these results in terms of fundamental wind parameters is difficult to do without computing line fluxes for complete wind models which is beyond the scope of this work. However, the upper limits on $(\Omega/4\pi)\Delta{}v$ that we obtain, approximately 1\,\kms, are quite small compared to 1000\,\kms{}, the order of the terminal wind velocities of massive stars and the wind velocity in a simulation of the wind of SMC~X-1 by \citet{blo95}. Therefore, the part of the wind that emits the helium-like lines cannot be homogeneously distributed around the entire wind. Instead, we believe that our results indicate either that the volume (or volumes) containing the helium-like ions of the various elements subtend a small solid angle at the neutron star or that the helium-like ions exist primarily in a region where the wind is at a small fraction of its terminal velocity or, perhaps, some combination of both. One possible explanation of our results is that the helium-like ions exist primarily in dense clumps throughout the wind. Dense clumps in the wind of an HMXB have previously been invoked in the case of Vela~X-1 by \citet{sak99} in order to explain the strong fluorescence lines observed from that system. Another possible explanation is that helium-like ions exist mainly near the surface of the star where the density is higher, the ionization less, and the velocity lower than in the outer parts of the wind \citep[see, e.g.,][]{lie01}. Furthermore, the part of a volume along the surface of the companion star with a thickness significantly less than one stellar radius that is visible during eclipse would subtend a solid angle significantly less than $4\pi$ from the neutron star. In Figure~\ref{fig:clump_face} we illustrate both of these possibilities. \begin{figure} \begin{center} \includegraphics[angle=-90,width=2.4in]{f6a.eps}\hspace{0.5in} \includegraphics[angle=-90,bb=-83 -27 295 263,width=2.4in]{f6b.eps} \end{center} \caption{In this figure we illustrate to possible distributions of the helium-like ions consistent with the helium-like line fluxes we measure. In the first panel we illustrate the possibility that the helium-like ions are in clumps distributed throughout the wind. In the second panel we illustrate the possibility that the helium-like ions are near the surface of the companion star.} \label{fig:clump_face} \end{figure} Definitive tests of these hypotheses would require calculations of the line emission from detailed wind models and that is beyond the scope of this work. However, we expect that our results, including the remarkably small values of $\Omega\Delta{}v$ are consistent with current expectations about the nature of winds in HMXBs. | 7 | 10 | 0710.3866 |
0710 | 0710.4488_arXiv.txt | We construct evolutionary models of the populations of active galactic nuclei (AGN) and supermassive black holes, in which the black hole mass function grows at the rate implied by the observed luminosity function, given assumptions about the radiative efficiency and the luminosity in Eddington units. We draw on a variety of recent X-ray and optical measurements to estimate the bolometric AGN luminosity function and compare to X-ray background data and the independent estimate of Hopkins et al.\ (2007) to assess remaining systematic uncertainties. The integrated AGN emissivity closely tracks the cosmic star formation history, suggesting that star formation and black hole growth are closely linked at all redshifts. We discuss observational uncertainties in the local black hole mass function, which remain substantial, with estimates of the integrated black hole mass density $\rho_{\bullet}$ spanning the range $3 - 5.5 \times 10^5\, {\rm M_{\odot}\, Mpc^{-3}}$. We find good agreement with estimates of the local mass function for a reference model where all active black holes have a fixed efficiency $\epsilon = 0.065$ and $L_{\rm bol}/L_{\rm Edd} \approx 0.4$ (shifting to $\epsilon = 0.09$, $L_{\rm bol}/L_{\rm Edd} \approx 0.9$ for the Hopkins et al.\ luminosity function). In our reference model, the duty cycle of $10^9 M_\odot$ black holes declines from 0.07 at $z=3$ to 0.004 at $z=1$ and $10^{-4}$ at $z=0$. The decline is shallower for less massive black holes, a signature of ``downsizing'' evolution in which more massive black holes build their mass earlier. The predicted duty cycles and AGN clustering bias in this model are in reasonable accord with observational estimates. If the typical Eddington ratio declines at $z<2$, then the ``downsizing'' of black hole growth is less pronounced. Models with reduced Eddington ratios at low redshift or black hole mass predict fewer low mass black holes ($M_{\bullet}\lesssim 10^8\, {\rm M_{\odot}}$) in the local universe, while models with black hole mergers predict more black holes at $M_{\bullet}>10^9\, {\rm M_{\odot}}$. Matching the integrated AGN emissivity to the local black hole mass density implies $\epsilon = 0.075 \times (\rho_{\bullet} / 4.5\times 10^5\, {\rm M_{\odot}\, Mpc^{-3}})^{-1}$ for our standard luminosity function estimate, or 25\% higher for Hopkins et al.'s estimate. It is difficult to reconcile current observations with a model in which most black holes have the high efficiencies $\epsilon \approx 0.16-0.20$ predicted by MHD simulations of disk accretion. We provide electronic tabulations of our bolometric luminosity function and our reference model predictions for black hole mass functions and duty cycles as a function of redshift. | \label{sec|intro} The long-standing hypothesis that quasars are powered by accretion onto supermassive black holes (Salpeter 1964; Lynden-Bell 1969; Rees 1984) is now strongly supported by many lines of evidence, including the apparent ubiquity of remnant black holes in the spheroids of local galaxies (Richstone et al. 1998). The strong correlations between the masses of central black holes and the luminosities, dynamical masses, and velocity dispersions of their host spheroids (e.g., Magorrian et al. 1998; Ferrarese \& Merritt 2000; McLure \& Dunlop 2002; Marconi \& Hunt 2003; H\"{a}ring \& Rix 2004; Ferrarese \& Ford 2005; Greene \& Ho 2006; Graham 2007; Hopkins et al. 2007b; Shankar \& Ferrarese 2008; Shankar et al. 2008) imply that the processes of black hole growth and bulge formation are intimately linked. Theoretical models typically tie black hole growth to episodes of rapid star formation, perhaps associated with galaxy mergers, and ascribe the black hole-bulge correlations to energy or momentum feedback from the black hole or to regulation of black hole growth by the bulge potential (e.g., Silk \& Rees 1998; Kauffmann \& Haehnelt 2000; Cavaliere \& Vittorini 2000; Granato et al. 2004, 2006; Murray, Quataert, \& Thompson 2004; Cattaneo et al. 2005; Miralda-Escud\`{e} \& Kollmeier 2005; Monaco \& Fontanot 2005; Croton et al. 2006; Hopkins et al. 2006a, Malbon et al. 2006). These correlations also make it possible to estimate the mass function of black holes in the local universe (e.g., Salucci et al. 1999; Yu \& Tremaine 2002; Marconi et al. 2004; Shankar et al. 2004, S04 hereafter). This local mass function provides important constraints on the co-evolution of the quasar and black hole populations. The most general and most well known of these constraints is the link between the integrated emissivity of the quasar population, the integrated mass density of remnant black holes, and the average radiative efficiency of black hole accretion (So{\l}tan 1982; Fabian \& Iwasawa 1999; Elvis et al. 2002). In the paper we construct self-consistent models of the quasar population using a method that can be considered a ``differential'' generalization of So{\l}tan's (1982) argument. Given assumed values of the radiative efficiency and the Eddington ratio $L/L_{\rm Edd}$, the observed luminosity function of quasars at a given redshift can be linked to the average growth rate of black holes of the corresponding mass, and these growth rates can be integrated forward in time to track the evolution of the black hole mass function. This modeling approach has been developed and applied in a variety of forms by numerous authors, drawing on steadily improving observational data (e.g. Cavaliere et al. 1973; Small \& Blandford 1992; Salucci et al. 1999; Cavaliere \& Vittorini 2000; Yu \& Tremaine 2002; Steed \& Weinberg 2003, hereafter SW03; Hosokawa 2004; Yu \& Lu 2004; Marconi et al. 2004; S04; Merloni 2004; Vittorini, Shankar \& Cavaliere 2005; Tamura et al. 2006; Hopkins et al. 2007). The consensus of recent studies is that evolutionary models with radiative efficiencies of roughly 10\% and mildly sub-Eddington accretion rates yield a reasonable match to the observational estimates of the remnant mass function. However, uncertainties in the bolometric luminosity function of active galactic nuclei (AGN, a term we will use to describe both quasars and less luminous systems powered by black hole accretion) and in the local black hole mass function remain an important source of uncertainty in these conclusions. We begin our investigation by constructing an estimate of the bolometric AGN luminosity function. Our estimate starts from the model of Ueda et al. (2003, hereafter U03), based on data from several X-ray surveys, but we adjust its parameters based on more recent optical and X-ray data that provide more complete coverage of luminosity and redshift. Comparison to the X-ray background and to the independent luminosity function estimate of Hopkins, Richards \& Hernquist (2007; HRH07 hereafter) gives an indication of the remaining uncertainties, associated mainly with bolometric corrections and the fraction of obscured sources. In agreement with Marconi et al. (2004) and Merloni (2004), we find similar trends in the evolution of AGN emissivity and the cosmic star formation rate, supporting models in which the growth of black holes occurs together with the formation of stars in their hosts (e.g., Sanders et al. 1988; Granato et al. 2004; Hopkins et al. 2006a). We also reassess current estimates of the local black hole mass function, finding that different choices and calibrations of the black hole-bulge correlation lead to significantly different results, with integrated mass densities that span nearly a factor of two. Our simplest models of the evolving black hole and AGN populations assume that all active black holes have a single radiative efficiency $\epsilon$ and a single accretion rate $\dot{m}$ in Eddington units. The work of Kollmeier et al. (2006) suggests that a constant value of $\dot{m}$ may be a reasonable approximation for luminous quasars, but the observational evidence on this point is mixed (e.g., McLure \& Dunlop 2004; Vestergaard 2004; Babic et al. 2007; Bundy et al. 2007; Netzer \& Trakhtenbrot 2007; Netzer et al. 2007; Rovilos \& Gorgantopoulos 2007; Shen et al. 2007b), and for lower luminosity, local AGN there is clearly a broad range of Eddington ratios (e.g., Heckmann et al. 2004). We also consider models in which $\dot{m}$ depends on redshift or black hole mass, and we consider a simple model of black hole mergers to assess their potential impact on the mass function. We will examine models with distributions of $\dot{m}$ values and a more realistic treatment of mergers in future work. Our modeling allows a consistency test of the basic scenario in which the observed luminosity of black hole accretion drives the growth of the underlying black hole population, and comparison to the local black hole mass function yields constraints on the average radiative efficiency and the typical accretion rate. For specified $\epsilon$ and $\dot{m}$, the model predicts the black hole mass function as a function of time, and combination with the observed luminosity function yields the duty cycle as a function of redshift and black hole mass. These predictions can be tested against observations of AGN host galaxy properties and AGN clustering, and they can be used as inputs for further modeling. While significant uncertainties remain, we find that a simple model of the black hole and AGN populations achieves a good match to a wide range of observational data. | \label{sec|DiscuConclu} We have constructed self-consistent models for the evolution of the supermassive black hole population and the AGN population, in which black holes grow at the rate implied by the observed luminosity function given assumed values of the radiative efficiency $\epsilon$ and the characteristic accretion rate $\dot{m}_0=\dot{M}_{\bullet}/\dot{M}_{\rm Edd}$ (see Eqs.~\ref{eq|conteq}, ~\ref{eq|mdotav}, ~\ref{eq|POsingle}). These models can be tested against the mass distribution of black holes in the local universe, and they make predictions for the duty cycles of black holes as a function of redshift and mass, which can be tested against observations of quasar hosts and quasar clustering if one assumes an approximately monotonic relation between the masses of black holes and the masses of their host galaxies and halos. Our method is similar to that used previously by Cavaliere et al. (1971), Small \& Blandford (1992), Yu \& Tremaine (2002), SW03, Marconi et al. (2004), and S04. However, we have drawn on more recent data on the AGN luminosity function and the local black hole mass function, and we have considered a broader spectrum of models. This approach can be considered a ``differential'' generalization of the So{\l}tan (1982) argument relating the integrated emissivity of the quasar population to the integrated mass density of the local black hole population. Our model for the bolometric AGN luminosity function starts from the model of U03, based on X-ray data from a variety of surveys, but we adjust its parameters as a function of redshift in light of more recent measurements and data at other wavelengths. In agreement with previous studies we find that the LF of optical AGNs is roughly consistent with that of X-ray AGN that have absorbing column densities $\log N_H/{\rm cm^{-2}}\le 22$ and that unobscured AGN dominate the bright end of the LF. We show that the latest constraints on the hard X-ray background ($E\sim 10-100$ keV) from \emph{INTEGRAL} and from the \emph{PDS} instrument on BeppoSax support a reduced normalization relative to extrapolations from other missions at lower energies. They therefore favor a lower contribution from very highly obscured AGN ($\log N_H/{\rm cm^{-2}}\gtrsim 24.5$) than some previous estimates (but see also Gilli et al. 2007). Our estimate of the bolometric AGN LF is independent in implementation but similar in spirit to that of HRH07. The most significant differences for purposes of this investigation are that HRH07 have a higher LF normalization at $z\sim 1-3$, principally because of their choice of bolometric correction, and they have a steeper bright-end slope at $z\sim 1-2.5$, where they adopt the slopes measured by Richards et al. (2006) from the SDSS and we use the slopes inferred by U03 from X-ray data. We regard the differences between our estimate and that of HRH07 as a reasonable indication of the remaining systematic uncertainties in the bolometric LF of AGNs. With our LF estimate, the bolometric emissivity of the AGN population tracks recent estimates of the cosmic star formation rate as a function of redshift. (Comparisons based on the space density of high luminosity quasars [e.g., Richards et al. 2006, Osmer 2004 and references therein] reach a different conclusion because at low redshifts the bright end of the AGN LF drops much more rapidly with time than the overall emissivity.) The integrated black hole mass density implied by this emissivity is $\sim 8\times 10^{-4}$ of the stellar mass at all redshifts, or about half of the estimated ratio of black hole mass to bulge stellar mass in local galaxies. This tracking favors scenarios in which black holes and the stellar mass of bulges grow in parallel, with about $50\%$ of the star formation linked to black hole growth at all redshifts. These findings are hard to reconcile with any models where black hole growth substantially precedes stellar mass buildup, or with recent claims that the ratio of black hole mass to stellar mass is much larger at high redshift than the local value. However, our finding refers to integrated densities, so it does not indicate the relative timing of black hole and bulge growth on an object-by-object basis. Observational estimates of the local black hole mass function still show substantial discrepancies among different authors, depending on the correlation used to derive it (e.g., $M_{\bullet}-\sigma$, $M_{\bullet}-L_{\rm bulge}$, $M_{\bullet}-M_{\rm star}$, $M_{\bullet}$-S\'{e}rsic index, or fundamental plane), the calibration of the correlation, and the intrinsic scatter of the correlation. Above $M_{\bullet}\sim 10^9\, {\rm M_{\odot}}$, estimates depend on extrapolation of the observed scaling relations, and below $M_{\bullet}\sim 10^{7.5}\, {\rm M_{\odot}}$ they are sensitive to the treatment of spiral bulges. The grey band in Figure~\ref{fig|LocalMFs} (and subsequent figures) encompasses most estimates, but the fundamental plane (Hopkins et al. 2007b) and S\'{e}rsic index (Graham et al. 2007) methods imply more sharply peaked mass functions. The integrated mass densities of all of these estimates are in the range $\rho_{\bullet}\sim 3-5.5\times 10^5\, {\rm M_{\odot}\,Mpc^{-3}}$. Our simplest models assume a single characteristic Eddington accretion rate $\dot{m}_0$, independent of redshift and black hole mass, and all of our models assume a single value of the radiative efficiency $\epsilon$. Matching the local black hole mass density requires $\epsilon=0.075\times(\rho_{\bullet}/4.5\times 10^5\, {\rm M_{\odot}\,Mpc^{-3}})^{-1}$ for our standard estimate of the AGN LF, or $\epsilon \sim 0.094(\rho_{\bullet}/4.5\times 10^5\, {\rm M_{\odot}\,Mpc^{-3}})^{-1}$ for the HRH07 luminosity function.\footnote{More precisely, it is $\epsilon/(1-\epsilon)$ that is proportional to $\rho_{\bullet}^{-1}$, but the difference from $\epsilon\propto \rho_{\bullet}^{-1}$ is tiny over the allowed range.} With $\epsilon$ thus fixed, the value of $\dot{m}_0$ determines the peak of the predicted local black hole mass function in the $M_{\bullet}\Phi(M_{\bullet})$ vs $M_{\bullet}$ plane. Note that with our definitions the Eddington luminosity ratio is $\lambda=L/L_{\rm Edd}\approx \dot{m}_0(\epsilon/0.1)$ (equation~\ref{eq|Lambda}). Matching the observed peak at $\log M_{\bullet}/{\rm M_{\odot}}\sim 8.5$ implies $\dot{m}_0\approx 0.6$ ($\lambda \approx 0.45$) for our standard LF estimate and $\dot{m}_0\approx 1$ ($\lambda \approx 0.95$) for the HRH07 LF. Lower values, $\dot{m}_0=0.1-0.3$, shift the peak location to untenably high masses, $\log M_{\bullet}/{\rm M_{\odot}}\sim 8.9-9.3$ or, for HRH07, $\log M_{\bullet}/{\rm M_{\odot}}\sim 9.1-9.6$. The single-$\dot{m}_0$ models achieve a reasonable match to our ``grey-band'' observational estimates of the width and asymmetry of the local mass function, though for our LF estimate the predicted mass function is too high at $M_{\bullet}>10^9\, {\rm M_{\odot}}$ and is therefore somewhat too broad. Single-$\dot{m}_0$ models cannot reproduce the more sharply peaked local mass functions estimated by Graham et al. (2007) or Hopkins et al. (2007b). Our reference model, which has $\epsilon=0.065$, $\dot{m}_0=0.60$, and our standard LF estimate, predicts a duty cycle for activity of $10^9\, {\rm M_{\odot}}$ black holes that declines steadily from 0.15 at $z=4$ to 0.07 ($z=3$), 0.035 ($z=2$), 0.004 ($z=1$), and $10^{-4}$ ($z=0$). The decline in duty cycle for lower mass black holes is much shallower. Massive black holes therefore build their mass relatively early while low mass black holes grow later, the phenomenon often referred to as ``downsizing''. Our results on mean radiative efficiency and duty cycle evolution are also in qualitative agreement with those found by Haiman et al. (2004). The predicted duty cycles seem in reasonable accord with observational estimates, though these estimates have considerable uncertainty and do not, as yet, probe mass and redshift dependence in much detail. The electronic tables described in \S~\ref{subsec|Tables} provide tabulations as a function of redshift of our AGN bolometric LF estimate, the HRH07 LF, and black hole mass functions and duty cycles for single-$\dot{m}_0$ models that are in good agreement with the observed $z=0$ mass functions given these LF inputs. We have examined models in which the Eddington ratio accretion rate $\dot{m}_0$ is reduced at low redshift or at low black hole mass. Declining redshift evolution of $\dot{m}_0$ damps ``downsizing,'' reducing the dependence of duty cycles on black hole mass and redshift. This model produces a typical duty cycle $P_0\sim 10^{-2.5}$ at $z=0$, about two orders of magnitude higher than in the reference model and consistent with the local duty cycle inferred from observations by Greene \& Ho (2007). In general terms, the observed luminosity-dependent density evolution of the AGN LF can be explained by preferential suppression of activity in high mass black holes at low redshift, by a decline in the typical accretion rate at low redshift, or by some combination thereof. The mass-dependent $\dot{m}_0$ model associates more of the AGN emissivity to high mass black holes, so it predicts a $z=0$ mass function that is more sharply peaked, in better agreement with the estimates of Hopkins et al. (2007b) and Graham et al. (2007) but worse agreement with other estimates. This model predicts a stronger mass dependence of duty cycles than our reference model because it maps low mass black holes, whose abundance is already suppressed, to less luminous, and hence more common, AGN. We have also considered a model in which each black hole has a $50\%$ probability per Hubble time of merging with another black hole of equal mass. Merger-driven growth in this model has little impact on the black hole mass function until $z<1$, when accretion-driven evolution has slowed. Low redshift mergers slightly depress the low mass end of the $z=0$ mass function and significantly enhance the high mass tail, worsening the agreement with observational estimates. Models incorporating theoretically predicted merger rates can allow more realistic calculations of the impact of mergers on the black hole population; this impact will probably be smaller than in the simplified model considered here. We have calculated the clustering bias of AGN as a function of luminosity and redshift for our reference model, assuming a monotonic relation between black hole mass and halo mass. The predictions are in reasonable accord with observational estimates. We will examine AGN clustering predictions in more detail in future work, with attention to what models can be excluded by the data and what can be learned by matching the full AGN correlation function in addition to an overall bias factor. MHD simulations (e.g., Gammie et al. 2004; Shapiro 2005) show that disk accretion onto Kerr black holes spins them up to an equilibrium spin rate $a\approx 0.95$ (where $a=1$ is the angular momentum parameter for a maximally rotating black hole). The radiative efficiency in these models is $\epsilon\approx 0.16-0.2$. These high efficiencies would lead to black hole mass densities a factor of two or more below our central estimate, and below our estimated lower bound. Furthermore, our results show that models with $\epsilon$ in the range $0.06-0.11$ can achieve a good match to the overall shape of the black hole mass function near the peak in $M_{\bullet}\Phi(M_{\bullet})$, not just the value of $\rho_{\bullet}$, given plausible choices of $\dot{m}_0$. Systematic uncertainties in the AGN LF do not appear large enough to accommodate $\epsilon\gtrsim 0.15$. Accommodating these high efficiencies would instead require a substantial downward revision of recent estimates of the local black hole mass function, reducing the integrated mass density to $\rho_{\bullet}\sim 2\times 10^5\, {\rm M_{\odot}\,Mpc^{-3}}$. Our results are consistent with a scenario like the one of King \& Pringle (2006) in which chaotic accretion spins down black holes because of counter-alignment with the accretion disk angular momentum, or with other mechanisms that reduce efficiencies below the MHD-simulation predictions. The assumption that all active black holes at a given mass and redshift have the same $\dot{m}_0$ is clearly an idealization, at least in the local universe where observations indicate a wide range of Eddington ratios (Heckman et al.\ 2005; Greene \& Ho 2007). Steed \& Weinberg (2003) and Yu \& Lu (2004) discussed continuity equation models evolved adopting a distribution of Eddington ratios. In particular, Yu \& Lu (2004) have derived the relation between the integrated number of AGNs shining at all times at a given luminosity $L$, the mean light curve of black holes, and the local black hole mass function. Following their equation (18), we find that a good match between the cumulative number of AGNs and of relic black holes can be achieved for $\epsilon\sim 0.07$ (required by the So{\l}tan argument) and a mean AGN light curve exponentially increasing with $\lambda=0.6$ and a negligible declining phase, similar to our reference model. Alternatively, we find that a good match can be obtained by assuming that black holes grow rapidly in a Super-Eddington phase with $\lambda\gtrsim 2$ and then have a long declining phase, qualitatively resembling our $\dot{m}(z)$ model discussed in \S~\ref{subsec|etamdot}. In future work we will investigate models that incorporate multiple $\dot{m}_0$-values and accretion modes, including the addition of modest log-normal scatter in $\dot{m}_0(M,z)$ (e.g., Kollmeier et al.\ 2005; Netzer et al.\ 2007; Shen et al.\ 2007b) and sharper revisions in which some black holes accrete at super-Eddington or highly sub-Eddington rates, perhaps with reduced radiative efficiencies (Narayan, Mahadevan, \& Quataert 1998 and references therein). We will also incorporate mergers at the rates predicted by theoretical models of cold dark matter subhalos and their associated black holes (Yoo et al.\ 2007). For appropriate parameter choices, we expect that many scenarios can be made consistent with the observed AGN LF and the local black hole mass function. However, direct measurements of Eddington ratio distributions and measurements of AGN clustering and host properties, all as a function of luminosity and redshift, should greatly narrow the field of viable models. Within the (often substantial) uncertainties of existing data, a simple model in which all black holes grow by accreting gas at mildly sub-Eddington rates with a radiative efficiency $\epsilon \approx 0.06-0.1$ is surprisingly successful at reproducing a wide range of observations. | 7 | 10 | 0710.4488 |
0710 | 0710.1863_arXiv.txt | We present the results of a search for molecular gas emission via the CO line in the far outer disk of the nearby spiral, NGC~6946. The positions targeted were chosen to lie on or near previously-identified outer disk HII regions. Molecular gas was clearly detected out to 1.3~R$_{25}$, with a further tentative detection at 1.4~R$_{25}$. The CO detections show excellent agreement with the HI velocities and imply beam-averaged column densities of $0.3-9\times 10^{20}$~cm$^{-2}$ and molecular gas masses of (2-70)$\times 10^{5}$~M$_{\sun}$ per 21$''$ beam (560pc). We find evidence for an abrupt decrease in the molecular fraction at the edge of the optical disk, similar to that seen previously in the azimuthally-averaged areal star formation rate. Our observations provide new constraints on the factors that determine the presence and detectability of molecular gas in the outskirts of galaxies, and suggest that neither the HI column, the metallicity or the local heating rate alone plays a dominant role. | It has become increasingly apparent in recent years that star formation is not just confined to the optically-bright parts of galaxies. Signposts of low-level star formation, such as HII regions and UV-bright clusters, have been discovered in the far outer regions of galactic disks, well beyond the classical R$_{25}$ radius (e.g. \cite{ferg98a,thilker05}), and in the tidal HI filaments of some interacting systems (e.g. \cite{ryanw04, Braine01}). The existence of star formation in these unusual environments, characterised by low metallicity, low HI column density and low interstellar radiation field, provides important constraints on star formation in galaxies, as well as the necessary conditions for sustaining a multi-phase interstellar medium (ISM). Much theoretical effort has been devoted to understanding the formation of a cold molecular phase in the ISM, a prerequisite for star formation. Thermal instabilities are generally believed to be the primary mechanism for converting cool atomic gas into cold molecular gas once a minimum pressure or local threshold surface density is exceeded (e.g. \cite{elme94,schaye04, blitz06}). Additional mechanisms for forming unstable cold clouds include shocks from spiral density waves and swing-amplifier and magneto-Jeans instabiities (e.g. \citet{dobbs06, kim02}). The threshold column densities predicted for molecular cloud formation are on average a few times smaller than those inferred for massive star formation in galaxies (e.g. \cite{skill87,mk01}). Direct detection of molecular clouds in extreme environments provides a means to test ideas about cloud formation and survival. Carbon monoxide (CO) is the most abundant heteronuclear molecule (i.e. with permitted rotational transitions) and is excited by collisions at densities typical of molecular clouds, making CO the standard tracer of molecular gas. Concerns have previously been raised \citep[e.g. ][]{Pfeniger94} as to whether the physical conditions in the outskirts of disks are sufficient to excite CO to detectable temperatures and whether the metallicities in these regions are sufficient to form enough CO for detection. With the discovery of CO emission out to 1.5~R$_{25}$ in the disk of NGC~4414 \citep{braine04} and in high column density tidal HI filaments in the M81 Group \citep{brou92,walter06}, these concerns have been proven somewhat unfounded. In this Letter, we present the results of a search for molecular gas in the far outer disk of the nearby spiral NGC~6946. While the study of molecular gas in NGC~4414 focused on high column HI regions \citep{braine04}, we have chosen to search for CO emission in the vicinity of outer disk star-forming regions. NGC~6946 was previously identified through H$\alpha$ observations to exhibit very extended massive star formation, out to $\sim 2$R$_{25}$ \citep{ferg98a}, and thus presents an excellent opportunity to attempt to measure and quantify molecular gas at extreme radii. We also report on the null detection of molecular gas at two positions in the far outer disk of another spiral with extended star formation, NGC~1058. | We have presented evidence for the existence of molecular gas in the far outskirts of NGC~6946 and noted a rather abrupt decrease in the molecular gas fraction, as well as the PAH emission, as one crosses R$_{25}$ along an individual spiral arm. An important question is whether this behavior could be explained by metallicity variations in the underlying gas disk. Indeed there is strong evidence for a radially increasing $\ratioo$ value \citep[e.g.][]{Sodroski95} as well as much higher values in low-metallicity systems \citep[e.g.][]{Rubio91}. The metallicity in the outer disk of NGC~6946 varies by a factor $\sim 5$ from log(O/H)$=-3.2$ to $-3.9$ \citep{ferg98b} while the I$_{\rm CO}$/N(HI) ratio decreases by a factor of $\sim30$. Focusing only on the positions which bridge R$_{25}$ (P12 to P2), the metallicity remains essentially constant while the I$_{\rm CO}$/N(HI) ratio decreases by a factor of $\sim6$. In the absence of a very strong radiation field, the $\ratioo$ factor should not vary more than linearly with metallicity and \citet{Wilson95} estimates that $\ratioo \propto [O/H]^{2/3}$ from a study of molecular clouds in Local Group galaxies. It is thus reasonable to expect that the $\ratioo$ ratio might be a factor 2 -- 3 higher in the most distant points than at P9, where the "standard" value is probably appropriate. Such a correction is not sufficient to explain the strong decrease in I$_{\rm CO}$/N(HI) at R$_{25}$ and we conclude that there is a genuine drop in the molecular gas content at the edge of the optical disk as defined by R$_{25}$. Although our pointings were selected to lie on or near outer disk HII regions, the H$\alpha$ luminosity contained within each beam shows significant variation and corresponds to local star formation rates (SFRs) ranging from $10^{-2}-10^{-4}$ M$_\odot$ yr$^{-1}$, using the \citet{kenn98} calibration (SFR$=$L(H$\alpha / 1.26 \times 10^{41}$ erg s$^{-1}$). The H$\alpha$ emission in some regions is predominantly diffuse (e.g. P12 and P7) and may more reflect ionization by photons produced elsewhere rather than {\it in situ} star formation \citep{ferg96}. The star formation efficiency (SFE), defined as SFR/M(H$_2$), ranges from $\sim 10-0.07$ Gyr$^{-1}$; eliminating the two outlier points (P12 and P2), the range reduces to $2 - 0.5$ Gyr$^{-1}$. The mere presence of a luminous HII region does not guarantee the detection of molecular gas. Position P1 in NGC~1058 is the second most luminous star-forming region targeted in our study, and is our second most sensitive integration, yet CO was not detected there. Intriguingly, there is a potential correlation between the detection of CO(2--1) emission and the presence of a luminous HII region in that the three non-detections of CO(2--1) are among the four lowest H$\alpha$ fluxes. With the exception of P4, the HI column densities at our observed positions within NGC~6946 are quite high, close to or greater than $10^{21}$cm$^{-2}$ on the scale of several hundred parsecs. The HI column density may be an important factor in determining whether molecular gas will be detected but a high HI column is clearly not a sufficient condition \citep[see also][]{Gardan07}. It may be purely coincidental that N(H$_2$)/N(HI) changes abruptly across R$_{25}$ in NGC~6946. While R$_{25}$ provides a means to characterise the optical extents of galaxies, it is not expected to relate to any underlying physical properties. In NGC~6946, neither the metallicity, the HI surface density nor the surface brightness of the stellar disk exhibit unusual features at this location. Interestingly, the azimuthally-averaged areal star formation rate in NGC~6946 does show a sharp decline near the edge of the optical disk ($\sim 0.8$~R$_{25}$), although widespread low-level star formation is observed much further out \citep{ferg98a,mk01}. The rough correspondence between the star formation and molecular profile breaks could provide support for threshold models \citep[e.g.][]{elme94,schaye04} although the physics that drives this behavior is still open to debate. Pressure based \citep{elme94,blitz06} and column density based \citep{schaye04} models are very difficult to distinguish once the HI surface density exceeds that of the stars. Larger-scale mapping of the outer disk of NGC~6946 and sensitive observations of the outskirts of additional galaxies are required in order to more thoroughly understand the molecular content and distribution at large galactocentric radii. | 7 | 10 | 0710.1863 |
0710 | 0710.1114_arXiv.txt | We present an \xmm\ detection of two low radio surface brightness SNRs, G85.4+0.7 and G85.9$-$0.6, discovered with the Canadian Galactic Plane Survey (CGPS). High-resolution \xmm\ images revealing the morphology of the diffuse emission, as well as discrete point sources, are presented and correlated with radio and \cha\ images. The new data also permit a spectroscopic analysis of the diffuse emission regions, and a spectroscopic and timing analysis of the point sources. Distances have been determined from \ion{H}{1} and CO data to be $3.5\pm 1.0$ kpc for SNR G85.4+0.7 and $4.8\pm 1.6$ kpc for SNR G85.9$-$0.6. The SNR G85.4+0.7 is found to have a temperature of $\sim12-13$ MK and a 0.5--2.5~keV luminosity of $\sim1-4\times 10^{33}D^2_{3.5}$ erg/s (where $D_{3.5}$ is the distance in units of 3.5 kpc), with an electron density $n_e$ of $\sim0.07-0.16(fD_{3.5})^{-1/2}$ cm$^{-3}$ (where $f$ is the volume filling factor), and a shock age of $\sim9-49(fD_{3.5})^{1/2}$ kyr. The SNR G85.9$-$0.6 is found to have a temperature of $\sim15-19$ MK and a 0.5--2.5~keV luminosity of $\sim1-4\times 10^{34}D^2_{4.8}$ erg/s (where $D_{4.8}$ is the distance in units of 4.8 kpc), with an electron density $n_e$ of $\sim0.04-0.10(fD_{4.8})^{-1/2}$ cm$^{-3}$ and a shock age of $\sim 12-42(fD_{4.8})^{1/2}$ kyr. Based on the data presented here, none of the point sources appears to be the neutron star associated with either SNR. | In 2001, two new supernova remnants (SNRs) with low radio surface brightness, G85.4+0.7 and G85.9$-$0.6, were discovered in Canadian Galactic Plane Survey (CGPS)\citep{tay03} data and confirmed in X-rays with \ro\ data \citep{kot01}. Both show distinct shells in the radio band with an extended region of X-ray emission in the centre. The radio surface brightness of G85.4+0.7 at 1 GHz is $\Sigma_{1 {\rm GHz}}\le 1\times 10^{-22}$ Watt ${\rm m}^{-2}$ ${\rm Hz}^{-1}$ ${\rm sr}^{-1}$ and the radio data also indicate that the SNR has a non-thermal shell with angular diameter $\approx0.4^{\circ}$ which is surrounded by a thermal shell with an angular diameter of $\approx0.6^{\circ}$ and is located within an H I bubble. The bubble also contains two B stars which may have been part of the same association as the SNR's progenitor star. G85.9$-$0.6 has a radio surface brightness $\Sigma_{1 {\rm GHz}}\le 2\times 10^{-22}$ Watt ${\rm m}^{-2}$ ${\rm Hz}^{-1}$ ${\rm sr}^{-1}$ and it has no discernible H I features, indicating that it is expanding into a low density medium, perhaps between the local and Perseus spiral arms. The most likely event which would produce an SNR in such a region would be a Type Ia supernova. X-ray observations are important to the study of SNRs, particularly those with low surface brightness, because they provide information about the morphology and emission processes of these objects, which are indicators of both the nature of the supernova explosion which formed them and the properties of the progenitor star. SNRs with low surface brightness are expected to be formed after the core collapse of a massive star in a Type Ib/c or Type II explosion, because the stellar wind would have blown away much of the interstellar medium (ISM) surrounding it, leaving a low ambient density into which the shock from the supernova expands. A Type Ia supernova can also result in a low surface brightness SNR if the surrounding ISM has a low density, such as it would if it were located between two spiral arms of the galaxy. Thermal X-ray emission from a SNR arises as the blast wave of the explosion travels through and shocks the ISM, and as a reverse shock travels back into and shocks the ejecta. The X-ray spectrum of the SNR gives information about the temperature, the density, and the luminosity of the shocked material, while imaging data provides information about the size and morphology of the region. The low angular and spectral resolutions as well as the small number of counts in the \ro\ data did not allow spectroscopy nor detailed imaging to be done, so \xmm\ observations, described in \S\ref{obs}, have been made in order to confirm the detection of the SNRs, and to perform imaging, spectroscopic and timing studies. Detailed X-ray imaging, described in \S\ref{imaging}, is used to map the diffuse emission and compare it to the location and size of the radio shells, and \cha\ data have been used to search for compact objects not resolved by XMM. Spectral parameters obtained in \S\ref{spec} lead to an estimate of such quantities as temperature, density, and luminosity of the SNRs. In \S\ref{ps} the point sources are catalogued and an attempt at identification is made by matching their positions to objects in other catalogues. Timing analysis is performed to identify any pulsar candidates. The distances to the SNRs are derived from \ion{H}{1} and CO data in \S\ref{dist}. The results are discussed in \S\ref{disc}. Preliminary results were presented in \cite{jac06} and \cite{saf06}. | \label{disc} \subsection{Determination of shock age and mass of emitting gas} To estimate $t$, the shock age, for the SNRs, the relation $t=\tau/n_{\rm e}$ is used, where $\tau$ is the upper limit of the ionization timescale from the VPSHOCK model. The electron density $n_{\rm e}$ is determined from the distance in cm $D_{\rm cm}$, angular size in radians $\alpha$, given by the diameter of the extraction region, and the relation $n_{\rm H}=n_{\rm e}/1.2$, where $n_{\rm H}$ is the volume density of hydrogen within. Given that ${\rm Norm}=\frac{\int n_{\rm e}n_{\rm{H}}dV}{10^{14}(4\pi D_{\rm cm}^2)}$, where Norm is the normalization of the VPSHOCK model, which is proportional to the emission measure, and the electron and hydrogen densities $n_{\rm e}$ and $n_{\rm{H}}$ are assumed to be uniform, a volume filling factor $f$ is employed so that $n_{\rm e}=(\frac{2.88\times 10^{15} ({\rm Norm})}{f \alpha^3D_{\rm cm}})^{1/2}$. From these calculations, it is determined that for G85.4+0.7, $n_e \approx (0.11\pm0.03)(fD_{3.5})^{-1/2}$ for the ESAS background (for the MOS instrument) or $(0.11^{+0.05}_{-0.04})(fD_{3.5})^{-1/2}$ for the observation background, and $t =(18^{+17}_{-9})(fD_{3.5})^{1/2}$ kyr for the ESAS background or $(23^{+26}_{-14})(fD_{3.5})^{1/2}$ kyr for the observation background. For G85.9$-$0.6, $n_e \approx (0.07^{+0.03}_{-0.02})(fD_{4.8})^{-1/2}$ for the ESAS background or $(0.07\pm0.03)(fD_{4.8})^{-1/2}$ for the observation background and $t = (30^{+12}_{-13})(fD_{4.8})^{1/2}$ kyr for the ESAS background or $(22^{+11}_{-10}) (fD_{4.8})^{1/2}$ kyr for the observation background. These are larger than the age estimates from the radio data. The mass of the emitting gas is calculated based on spherical emitting regions with the size given by the extraction radius (given in Table~\ref{dfit}), composed of 92\% hydrogen and 8\% helium, and the relation $n_{\rm H}=n_{\rm e}/1.2$ is used as for the above calculation of $t$. The mass of the emitting gas in G85.4+0.7 is $(2.8\pm1.5) D_{3.5}^{5/2} M_{\sun}$ for the ESAS background and $(2.8^{+1.8}_{-1.6})D_{3.5}^{5/2} M_{\sun}$ for the observation background and that for G85.9$-$0.6 is $(2.7^{+1.8}_{-1.6}) D_{4.8}^{5/2} M_{\sun}$ for the ESAS background and $(2.7\pm1.8)D_{4.8}^{5/2} M_{\sun}$ for the observation background. \subsection{Background subtraction} The background subtraction was problematic, particularly for the spectrum of G85.4+0.7, which is fainter and required a different model to be fit to it depending on which background region on the observation was used, leading to a very careful selection of the background extraction region, the process of which is described in \S\ref{bg}. Instrumental lines, which are fortunately different for PN and MOS, nevertheless increased the value of $\chi^2$ for the combined fits, even when the background was chosen very carefully, and needed to be explicitly fit when the ESAS background was used, as described in \S\ref{bg}. In addition, the soft spectra of the SNRs exhibit large amounts of noise in the high energy end of the spectra, allowing fits to be made only up to 2.5 keV. Future longer observations of these SNRs will help to resolve these difficulties, allow for better fits to the abundances, eliminate ambiguities in the spectral results, and perhaps allow for conclusive identifications of neutron star and pulsar candidates in addition to other point sources. \subsection{G85.4+0.7}\label{disc85.4} A comparison between the morphologies of the X-ray diffuse emission and the radio emission of G85.4+0.7 can be seen in Figure~\ref{85.4imagens} in which the X-ray point sources have been removed and the X-ray image has been smoothed to 1$^{\prime}$ to match the radio contours. The diffuse X-ray emission exhibits some structure, and it lies in the approximate centre of the radio shell. A spectrum was extracted from the central blob, the position and size of which is shown in the top panel of Figure~\ref{85.4reg}, and it could not be adequately fit with a non-thermal power law model, indicating that a pulsar wind nebula origin for the emission can be ruled out. The fact that the SNR does not exhibit any limb brightening and the X-ray emission is mostly concentrated in the centre, as projected in the plane of the sky and three dimensionally, can be interpreted as evidence that the X-ray emission is produced by the ejecta. The fact that there appears to be no emission from the swept up material can be explained if the remnant is evolutionarily young. With the ejecta interpretation, if it is assumed that the free electrons are evenly distributed in the extraction area, the resulting ages are 18 and 23 kyr for the two backgrounds. In Figure~\ref{85.4imagens} it can be seen that the outer radio shell has an approximately constant radius, whereas the inner shell has a smaller radius in the vertical centre, which expands as the latitude changes, and this indicates that the SNR is moving either toward us or away from us. The two shells appear to meet near the bottom of the image, so the radius of the SNR can be approximated as that of the outer shell, which is $\sim0.3^\circ$ on the image. The average velocity of the expanding SNR is then 910 or 730 km/s for the two ages. However, the morphology of the X-ray emission is not very smooth, indicating that the filling factor $f<<1$. Assuming $f=0.2$, the electron density $n_e$ would be 0.25 cm$^{-3}$ for either the ESAS and observation background, and the ejecta masses would be 1.3 $M_{\sun}$ for either background. The shock age would then be 8000 or 10300 years, and the average expansion velocity would be 2150 or 1640 km/s. For a $10^{51}$ erg supernova explosion, the ejecta mass would be 22 or 37 $M_{\sun}$, or 2.2 or 3.7 $M_{\sun}$ for a $10^{50}$ erg explosion, indicating that the lower energy explosion would result in an ejecta mass consistent with a young freely expanding SNR. The radio continuum emission indicates that there is some swept up material. Assuming that the density inside the stellar wind bubble was $\sim 0.01$ cm$^{-3}$ before the explosion, the SNR would have swept up $\sim 6.5$ $M_{\sun}$ of material, which is on the order of the ejecta mass, and means that the SNR is in the transition between free expansion and Sedov expansion. The slightly above-solar abundance of O in the diffuse spectra reinforces the hypothesis that the supernova most likely resulted from a core collapse, though the large error bars weaken the argument. The abundances of all elements other than O and Fe are below solar, but they would be enhanced if the X-rays were from the ejecta. However, it is possible that the spectral parameters for the abundances are affected by the poor quality of the spectra and by the uncertainties associated with the background subtraction. Since the most likely origin of G85.4+0.7 was a core collapse supernova (see \S\ref{dist}), it is possible that a neutron star or pulsar which is associated with this SNR can be found. Given its power law X-ray spectrum, proximity to the centre of the remnant, and similar $N_{\rm H}$ value to the diffuse emission, source 1 in Figure \ref{85.4reg} is possibly the associated neutron star, but sources 4 and 9 are within the radio shell and are also candidates, given their spectral parameters. However, the low fit quality of source 4 ($\chi^2\sim2$) indicates that it is not likely to be a neutron star, and source 9 is situated far from the centre of the diffuse emission so is less likely to be the associated neutron star. \cha\ sources 10, 11, 12, 13, 15, or 16 in Figure~\ref{cha85.4} can also be considered as neutron star candidates (provided the optical counterparts listed in Table~\ref{pscat85.4} are not the true counterparts. \subsection{G85.9$-$0.6}\label{disc85.9} Figure~\ref{85.9imagens} shows X-ray and radio images of G85.9$-$0.6, produced in a similar way to Figure~\ref{85.4imagens}. The diffuse X-ray emission appears to contain less structure than G85.4+0.7, and, as for G85.4+0.7, a spectrum extracted from the central blob, the position and size of which are shown in the top panel of Figure~\ref{85.9reg}, indicates that a pulsar wind nebula origin for the emission can be ruled out. Because Figure~\ref{85.4imagens} shows no limb brightening and the emission is mostly from the central blob, a similar argument to that used for G85.4+0.7 can be employed here to suggest an ejecta interpretation for the X-ray emission. Assuming that the free electrons are evenly distributed in the extraction area the resulting age is 30 or 22 kyr for the two backgrounds. Using the average distance between the centre of the X-ray emission and the shell, $\sim 0.2^\circ$, the radius of the shell is 15.3 pc The average expansion velocity would then be 500 or 680 km/s. The emission is concentrated in the inner $6^\prime$, which indicates an electron density of 0.20 cm$^{-3}$ and an ejecta mass of 1.0 $M_{\sun}$ which agrees with the predicted mass for a type Ia explosion, 1.4 $M_{\sun}$. This would result in an age of 10600 or 7800 years and average expansion velocity of 1350 or 1850 km/s. For a Type Ia supernova, the explosion energy is $\sim10^{51}$ erg and the ejecta mass is 1.4 $M_{\sun}$. Unlike for G85.4+0.7, the density in the interarm region is closer to 0.1 cm$^{-3}$, resulting in a swept up mass of nearly 50 $M_{\sun}$. This indicates that the SNR is in the Sedov expansion phase, which is described by $R=14 (E_0/n_0)^{1/5}t^{2/5}$, where $R$ is the radius in pc, $E_0$ is the explosion energy in units of $10^{51}$ erg, $n_0$ is the ambient density in cm$^{-3}$, and $t$ is the age in units of $10^4$ years. For an explosion energy of 10$^{51}$ erg, a radius of 15.3 pc, and an age of 10600 or 7800 years, the resulting ambient density is 0.84 or 0.55 cm$^{-3}$ which are consistent with the density of the interarm region, but the swept up mass would be 360 or 240 $M_{\sun}$, from which it should be possible to measure thermal X-ray emission, from the part of the shell that is included in the X-ray pointing. The current expansion velocity would be $dR/dt=0.4(R/t)$, which is 570 or 760 km/s. The interpretation of G85.9$-$0.6 as having been produced by a Type Ia supernova is reinforced by the Fe abundance, which is well above solar. As with G85.4+0.7, the ejecta interpretation is called into question by the remaining abundances, which should be above solar, but are instead below solar. This could again be a result of poor quality spectra. Given that the radio results described in \cite{kot01}, as well as the distance presented here, indicated that G85.9$-$0.6 was most likely produced by a Type Ia supernova, it was not expected to find a neutron star associated with this SNR. The above solar Fe abundance for this SNR is consistent with a Type Ia explosion. An identification of sources 6 and 7 in Figure~\ref{85.9reg} has not yet been made. They are bright X-ray emitting objects with no known optical or radio counterpart, making them good neutron star or radio-quiet AGN candidates, though if one of them were a neutron star, it would be unlikely that it is associated with the G85.9$-$0.6 SNR because they are both situated far outside the radio shell of the remnant, and furthermore, the value of $N_{\rm H}$ for source 6 does not match that of the SNR itself, and the fit quality of source 7 ($\chi^2\sim2$) indicates that an absorbed power law is not a good fit. Their identification with possible 2MASS counterparts (see Table~\ref{pscat85.9}) makes them more likely radio-quiet AGN than neutron stars, provided the 2MASS objects are the true counterparts. A future detailed deep X-ray or multiwavelength study of these objects should be undertaken to identify and further study them, even though they are probably not associated with the G85.9$-$0.6 SNR. Source 4 is on the edge of the radio shell of the SNR, but its identification as a neutron star is questionable because of its photon index and hardness ratio, as well as the fact that it is not expected that there is a neutron star associated with this remnant. \subsection{Mixed Morphology Interpretation} The centrally filled morphology of both SNRs and the thermal nature of their X-ray emission confined within the radio shells suggest that they belong to the class of mixed-morphology SNRs (also known as thermal composites; Rho \& Petre 1997). The origin of the thermal X-ray emission interior to the radio shells in these SNRs has been attributed to several mechanisms which include: a) cloudlet evaporation in the SNR interior (White \& Long 1991), b) thermal conduction smoothing out the temperature gradient across the SNR and enhancing the central density (Cox et al. 1999), c) a radiatively cooled rim with a hot interior (e.g. Harrus et al. 1997), and d) possible interaction with a nearby cloud (e.g. Safi-Harb et al. 2005). While modeling these SNRs in the light of the above mentioned models is beyond the scope of this paper and has to await better quality data, we can rule out the cloudlet evaporation model based on the low ambient densities inferred from our spectral fits (see \S7.1 and Table~1). Except for Fe and possibly O, the abundances inferred from our spectral fits are for the most elements consistent with or below solar values, as observed in most mixed-morphology SNRs. However, enhanced metal abundances have been observed in younger SNRs, e.g. 3C~397, estimated to be $\sim$5.3 kyr--old and proposed to be evolving into the mixed-morphology phase (Safi-Harb et al. 2005). The ages inferred for G85.4+0.7 ($\sim$8--10 kyr; see Table~1 and \S\ref{disc85.4}) and G85.9$-$0.6 ($\sim$8--11 kyrs; see Table~1 and \S\ref{disc85.9}) suggest a later evolutionary phase where only shock-heated ejecta from Fe (for G85.9$-$0.6) and possibly Oxygen are still observed. | 7 | 10 | 0710.1114 |
0710 | 0710.2737_arXiv.txt | We measure the relative evolution of the number of bright and faint (as faint as 0.05 $L^*$) red galaxies in a sample of 28 clusters, of which 16 are at $0.50\le z \le 1.27$, all observed through a pair of filters bracketing the $4000$ \AA \ break rest-frame. The abundance of red galaxies, relative to bright ones, is constant over all the studied redshift range, $0<z<1.3$, and rules out a differential evolution between bright and faint red galaxies as large as claimed in some past works. Faint red galaxies are largely assembled and in place at $z=1.3$ and their abundance does not depend on cluster mass, parametrized by velocity dispersion or X-ray luminosity. Our analysis, with respect to previous one, samples a wider redshift range, minimizes systematics and put a more attention to statistical issues, keeping at the same time a large number of clusters. | The evolution of faint red galaxies in clusters is a highly debated topic for two reasons: different observers have claimed controversial results, and clusters of galaxies are often claimed to be interesting laboratories where studying the effect of the environment. Red galaxies, in particular, have different assembly histories in halos of different masses, yet observationally the detection of a environmental dependence of their properties escapes a detection. For example, differences between cluster and field fundamental planes are small, if any (Pahre et al. 1998), so small that the Coma cluster fundamental plane (Jorgensen et al. 1996) is routinely used as zero-redshift reference for studying the evolution of field galaxies, and so small that previously claimed differences are probably due to having overlooked the difficulty of the statistical analysis (van Dokkum \& van der Marel 2007). Similarly, the colour of the red sequence seems not to depend on clustercentric distance (Pimbblet et al. 2002; Andreon 2003) or galaxy number density (Hogg et al. 2004, Cool et al. 2006). The red colour, by which red galaxies are defined and selected, induces a selection effect: at every redshift only galaxies whose stellar populations are red (i.e. old, modulo dust, of no interest here) enter the sample. It is not a surprise, then, to find old-selected populations to be old. A different question is whether galaxies that have an old stellar population were fully assembled at early or late times. Answering this question requires a measurement of the abundance of red galaxies as a function of look-back time. For clusters, there is a further complication: clusters have different richnesses, jeopardizing any look-back time trend if the richness dependence is not factored out. It is easy, furthermore, to qualitatively claim that the red sequence is built later (i.e. a lower redshift) in poor environments than it is in dense environments, but this might just be do to signal to noise issues, because in poorer environments the red population is a minority one, and its contrast with respect to other populations (e.g. background) noisier. A sound statistical assessment of the abundance of faint red galaxies is therefore compelling. \begin{table} \caption{The ACS $z\ge0.5$ cluster and control field samples} \begin{tabular}{l r c c l l} \hline Name & z & $N^{1}$ &\multispan{2}{\hfill Filters \hfill} \\ & & & blue & red \\ \hline Lynx W & 1.27 & 3 & F775W & F850LP \\ Lynx E & 1.26 & 3 & F775W & F850LP \\ RDCS J1252-2927 & 1.23 & 4 & F775W & F850LP \\ RDCS J0910+5422 & 1.11 & 1 & F775W & F850LP \\ GHO 1602+4329 & 0.92 & 1 & F606W & F814W \\ GHO 1602+4312 & 0.90 & 1 & F606W & F814W \\ 1WGA J1226.9+3332 & 0.89 & 6 & F606W & F814W \\ MACS J0744.8+3927 & 0.70 & 1 & F555W & F814W \\ MACS J2129.4-0741 & 0.59 & 1 & F555W & F814W \\ MACS J0717.5+3745 & 0.55 & 1 & F555W & F814W \\ MACS J1423.8+2404 & 0.54 & 1 & F555W & F814W \\ MACS J1149.5+2223 & 0.54 & 1 & F555W & F814W \\ MACS J0911.2+1746 & 0.50 & 1 & F555W & F814W \\ MACS J2214.9-1359 & 0.50 & 1 & F555W & F814W \\ MACS J0257.1-2325 & 0.50 & 1 & F555W & F814W \\ \\ CT344 & & 1 & F606W & F814W \\ B0455 & & 1 & F555W & F814W \\ GOODS+PAN & & $\sim30$ & F775W & F850LP \\ \hline \end{tabular} \hfill \break \footnotesize{$^1$ number of ACS field of view per filter. \hfill\break All clusters have coordinates and redshift listed in NED, except for MACS clusters, listed in Ebeling et al. (2007). MACS clusters have been also studied by Stott et al. (2007). \hfill\break} \end{table} Usually, the richness dependence of the abundance of faint red galaxies is removed by normalizing it to the number of bright red galaxies, i.e. by computing the faint-to-luminous ratio, or any related quantity, like the faint-end slope $\alpha$ of the luminosity function. The analysis of the faint-to-luminous ratio, performed by Stott et al. (2007), or its reciprocal, the luminous-to-faint ratio by De Lucia et al. (2007), both suggest an evolution of the relative abundance of faint red galaxies, in the sense that at high redshift there is a deficit of faint red galaxies per unit bright galaxy. On the other end, Andreon (2006a) suggests no deficit of red galaxies, using a very small cluster sample, and Andreon et al. (2006) discard a considerable build up of the red sequence on the basis of fossil evidence. Evidences presented in earlier works have been discussed in the mentioned papers and references therein. In this paper, we aim to understand if the colour--magnitude relation has been build up at early or late times, by studying many galaxy clusters at several look-back times. Throughout this paper we assume $\Omega_M=0.3$, $\Omega_\Lambda=0.7$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. All results of our stochastical computations are quoted in the form $x\pm\sigma$ where $x$ is the posterior mean and $\sigma$ is the posterior standard deviation. | The history of mass assembly of bright (massive) red galaxies in clusters is pretty well known: they were assembled at early times, as testified by the passive evolution of their characteristic magnitude (e.g. de Propris et al. 1998, 2007; Andreon et al. 2007), the constancy of their stellar mass function (e.g. Andreon 2006b) and of the halo occupation number, i.e. the number of galaxies per unit cluster mass (Lin et al. 2006, Andreon et al. 2008). We stress that all mentioned works favour the above scenario, but only one, (Andreon 2006b), excludes contender models, and we emphasize that most mentioned works have samples that are dominated, but not exclusively composed, by red galaxies. The history of mass assembly of faint red galaxies is far less clear. The present paper shows that a non-evolution of the faint end slope $\alpha$, or any related number such as the luminous-to-faint ratio, is fully compatible with the data. This implies that the history of mass assembly of faint red galaxies is strictly parallel to the one of their massive cousins, in order to keep the relative abundance constant. Therefore, the build up of the red sequence is largely complete by $z=1.3$ down to $0.05 L^*$, and, if a differential filling is envisaged, it should occur mostly at much larger redshift. Similarly, cluster mass, as parametrized by X-ray luminosity or velocity dispersion, seems not to play any role in shaping the relative abundance of faint galaxies, contrary to some previous claims. Our claims are based on one of the largest samples, spread over the wider redshift range studied thus far with a large cluster sample, with great attention to systematics. A recent ($z<1.3$) transformation of many blue galaxies in faint red galaxies would modify the faint-end slope of the LF, change the $F/L$ ratio and inflate the color scatter of the colour-magnitude relation, none of which have been observed. Yet, a redshift trend is expected because of the spiral morphology of some faint red galaxies in nearby clusters, but a larger sample of clusters (at $z \gg 0.5 $) is needed to measure its small amplitude. The present sample is however large enough to discard the claimed steep trends previously suggested in literature. | 7 | 10 | 0710.2737 |
0710 | 0710.0168_arXiv.txt | We present observations of the HCN and HCO$^+$ J=1-0 transitions in the center of the nearby spiral galaxy NGC 6946 made with the BIMA and CARMA interferometers. Using the BIMA SONG CO map, we investigate the change in the $I_{\rm HCN}/I_{\rm CO}$ and $I_{\rm HCO+}/I_{\rm CO}$ integrated intensity ratios as a function of radius in the central kiloparsec of the galaxy, and find that they are strongly concentrated at the center. We use the 2MASS $K_S$ band image to find the stellar surface density, and then construct a map of the hydrostatic midplane pressure. We apply a PDR model to the observed $I_{\rm HCN}/I_{\rm HCO+}$ integrated intensity ratio to calculate the number density of molecular hydrogen in the dense gas tracer emitting region, and find that it is roughly constant at $10^5$ cm$^{-3}$ across our map. We explore two hypotheses for the distribution of the dense gas. If the HCN and HCO$^+$ emission comes from self-gravitating density peaks inside of a less dense gas distribution, there is a linear proportionality between the internal velocity dispersion of the dense gas and the size of the density peak. Alternatively, the HCN and HCO$^+$ emission could come from dense gas homogeneously distributed throughout the center and bound by ambient pressure, similar to what is observed toward the center of the Milky Way. We find both of these hypotheses to be plausible. We fit the relationships between $I_{\rm HCN}$, $I_{\rm HCO+}$, and $I_{\rm CO}$. Correlations between the hydrostatic midplane pressure and $I_{\rm HCN}$ and $I_{\rm HCO+}$ are demonstrated, and power law fits are provided. We confirm the validity of a relation found by \citet{BR2006} between pressure and the molecular to atomic gas ratio in the high hydrostatic midplane pressure regime ($10^6$-$10^8$ cm$^{-3}$ K). | How is molecular gas organized in galaxies? In the disk of the Milky Way, molecular gas as traced by CO is concentrated in giant molecular clouds (GMCs) with a characteristic size of 50 pc \citep{B1987}. These clouds have masses around 10$^4$--10$^6$ $M_\odot$ and contain 80\% of the molecular mass \citep{B1993}. The hydrostatic pressure from the self-gravity of a GMC this size is $\sim10^5$ cm$^{-3}$ K \citep{B1993}, which is an order of magnitude larger than the mean pressure of the local ISM \citep{B1987a}; this leads to the conclusion that GMCs in the Milky Way disk are self-gravitating. GMCs in M31 \citep{R2007} and M33 \citep{RB2004,RKMW2007} have similar properties to those observed in the Milky Way. The densest parts of GMCs are traced by molecules with larger dipole moments than CO, such as HCN, HCO$^+$, and CS. These dense gas tracers need a molecular hydrogen density near $10^5$ cm$^{-3}$ to be excited collisionally, 100 times greater than CO. In the disk, emission from these molecules is limited to the star-forming cores of GMCs \citep{PJE1992,WE2003,HB1997b}. However, the molecular gas composition of the central kiloparsec of a spiral galaxy is very different from that of its disk. In the Milky Way, widespread CO emission in the inner kpc makes it difficult to differentiate between a distribution of discrete clouds or one in which the molecular gas is organized homogeneously \citep{DUCD1987}. The emission from HCN and CS within $\sim$250 pc of the Galactic center is distributed similarly to the CO \citep{JHPB1996,BSWH1987}, suggesting that high density regions are common, at least compared to the average ISM density of $\sim$1 cm$^{-3}$ in the solar neighborhood. In other galaxies, dense gas tracers are more concentrated than CO in the central kpc and have lower intensities elsewhere \citep{GS2004,HB1997a}, consistent with observations of the Milky Way \citep{HB1997b}. Furthermore, the centers of spiral galaxies tend to have high densities and pressures; bulges have hydrostatic midplane pressures 2-3 orders of magnitude higher than those of their disks \citep{SB1992}. \citet{BR2006} (hereafter BR06) found a simple relationship in disk galaxies between the hydrostatic midplane pressure and the ratio of molecular to atomic gas that holds in disks and bulges. In this paper, we use new HCN and HCO$^+$ observations of the central region of NGC 6946 to study molecular gas organization in this nearby spiral galaxy. Optically, this galaxy is classified as SAB(rs)cd \citep{D1963}, it has a moderate starburst in its nucleus \citep{TH1983}, and it is at a distance of 5.5 Mpc \citep{T1988}. At this distance, 1'' corresponds to 27 pc. Observations of NGC 6946 in the R, I, and K bands reveal a small bulge (radius 15'') within a larger bar (radius 63'') \citep{DTVR2003}. Recent studies revealed the rich details of its nuclear structure; \citet{SBEL2006} used high resolution images of CO to map the molecular gas spiral in the inner 10'', likely caused by an inner stellar bar 15'' in length (see also NIR images in \citealt{ECS1998} and \citealt{KDSB2003}). A further study used HCN emission in the central 2'' as a proxy for star formation in young clusters still embedded in their parent clouds, and combined these data with observations of HII regions to map the total amount of star formation \citep{SBED2007}. They concluded the nuclear star formation is being fed by an accumulation of dense gas driven by the inner stellar bar. Using HCN and HCO$^+$ observations in conjunction with CO and near-infrared data, we investigate the distribution of dense gas in the central kiloparsec of NGC 6946. After reducing our observations, we calculate HCN/CO and HCO$^+$/CO integrated intensity ratios, and we estimate the pressure using two different methods. We then describe two hypotheses for the distribution of the dense molecular gas, and explore the relationship between dense gas tracers and hydrostatic midplane pressure. | Using the BIMA and CARMA interferometers, we produced high-resolution maps of the dense gas tracers HCN and HCO$^+$ in NGC 6946. The maps are rich with detail, and we have 13 independent data points across the inner $\sim$750 pc of the galactic center in HCN and CO, and 7 in HCO$^+$. The integrated intensities of HCN and HCO$^+$ peak near the center of the galaxy, and fall off from there. Using BIMA SONG CO 1-0 and 2MASS $K_S$ band images, we calculated an approximation to the hydrostatic midplane pressure in the center of NGC 6946. We used a PDR model to calculate the number density in the dense gas regions traced by HCN and HCO$^+$, and found that the number density in these regions does not vary as a function of radius or the surrounding hydrostatic midplane pressure. We explored two hypotheses for the distribution of the dense gas in the central kpc of NGC 6946. If the dense gas is concentrated in clumps, we showed that self-gravity implies that $v_{\rm int}\propto r$. However, we also explored another plausible dense gas distribution; the dense gas could instead be homogeneously distributed throughout the galactic center if thermal pressure provides enough support for hydrostatic equilibrium. We then demonstrated clear correlations between the HCN and HCO$^+$ integrated intensities and the midplane pressure. Finally, we confirmed the validity of the relation between hydrostatic midplane pressure and the molecular to atomic gas ratio found in BR06 in the high pressure regime. | 7 | 10 | 0710.0168 |
0710 | 0710.0397_arXiv.txt | Our understanding of Pulsar Wind Nebulae (PWNe), has greatly improved in the last years thanks to unprecedented high resolution images taken from the HUBBLE, CHANDRA and XMM satellites. The discovery of complex but similar inner features, with the presence of unexpected axisymmetric rings and jets, has prompted a new investigation into the dynamics of the interaction of the pulsar winds with the surrounding SNR, which, thanks to the improvement in the computational resources, has let to a better understanding of the properties of these objects. On the other hand the discovery of non-thermal emission from bow shock PWNe, and of systems with a complex interaction between pulsar and SNR, has led to the development of more reliable evolutionary models. I will review the standard theory of PWNe, their evolution, and the current status in the modeling of their emission properties, in particular I will show that our evolutionary models are able to describe the observations, and that the X-ray emission can now be reproduced with sufficient accuracy, to the point that we can use these nebulae to investigate fundamental issues as the properties of relativistic outflows and particle acceleration. | Pulsar Wind Nebulae (PWNe) are bubbles of relativistic particles and magnetic field created when the ultra-relativistic wind from a pulsar interacts with the ambient medium, either SNR or ISM. The best example of a PWN, often considered the prototype of this entire class of objects, is the Crab Nebula. The first theoretical model of PWNe was presented by \citet{ree74}, developed in more details by \citet{ken84a,ken84b} (KC84 hereafter), and is based on a relativistic MHD description. The ultra-relativistic pulsar wind is confined inside the SNR, and slowed down to non relativistic speeds in a strong termination shock (TS). At the shock the toroidal magnetic field of the wind is compressed, the plasma is heated and particles are accelerated to high energies. A bubble of high energy particles and magnetic field is produced where the post-shock flow expands at a non relativistic speeds toward the edge of the nebula. Despite its simplicity the MHD model can explain many of the observed properties of PWNe. Acceleration at the TS accounts for the continuous, non-thermal, very broad-band spectrum, extending from Radio to X-rays, with spectral index in the range 0-1.2, steepening with increasing frequencies, and modeled as synchrotron emission \citep{ver93,ban99,wei00,wil01,mor04}. The under-luminous region, centered on the location of the pulsar, is interpreted as the ultra-relativistic unshocked wind. Polarization measures \citep{wil72,vel85,sch79,hic90,mic91} show that emission is highly polarized and the nebular magnetic field is mostly toroidal, as one would expect from the compression of the pulsar wind. This model also predicts that PWNe should appear bigger at smaller frequencies: high energy X-rays emitting particles have a short lifetime for synchrotron losses, and they are are present only in the vicinity of the TS; in contrast the synchrotron lifetime for Radio particles is longer than the age of the nebula, so they fill the entire volume. This increase in size at smaller frequencies is observed in the Crab Nebula \citep{ver93,bie97,ban98}. If one considers the pressure anisotropy due to the compressed nebular toroidal magnetic field \citep{beg92,van03}, it is also possible to recover the elongated axisymmetric shape of many PWNe ({\it i.e.} Crab Nebula, 3C58). By comparing observations with the predictions of the simple spherically symmetric MHD model it is possible to constraint some of the properties of the pulsar wind, at least at the distance of the TS. To explain the dynamics of the plasma, as well as the emission properties of the nebula, the Lorentz factor of the wind is estimated to be $\sim 10^6$, and the ratio between Poynting flux and kinetic energy $\sigma\sim 0.003$. This shows that nebular properties can be used to derive informations on the conditions of the pulsar wind at large distances. | \label{sec:concl} In the last few years, the combination of high resolution observations, and numerical simulations, has improved our understanding of the evolution and internal dynamics of PWNe. We can reproduce the observed jet-torus structure and we can relate the formation of the jet in the post shock flow to the wind magnetization. Simulated maps can reproduce many of the observed features, including the details of spectral properties. Results suggest that, the best agreement is achieved in the case of a striped wind, even if MHD simulations are not able to distinguish between dissipation of the current sheet in the wind or at the TS. Results also suggest that it is possible to use X-rays imaging to constrain the pulsar wind properties; already the rings and tori observed in many PWNe have been used to determine the spin axis of the pulsar \citep{rom05}. Interestingly in Crab the inner ring does not appear boosted, while the wisps (which are interpreted as its optical counterpart) are. There are still however unsolved questions, and possible future developments for research in this field. All present simulations are axisymmetric, and none is able to address the problem of the stability of the toroidal field, and cannot reproduce the observed variability in the jet. It is not clear if small scale disordered field is present in the inner region (may be a residual of the dissipation in the TS of the striped wind). A combination of simulations and polarimetry might help to answer this question. \begin{theacknowledgments} This work was supported partly by NASA through Hubble Fellowship grant HST-HF-01193.01-A, awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS 5-26555. \end{theacknowledgments} | 7 | 10 | 0710.0397 |
0710 | 0710.3394_arXiv.txt | {}{We analyze systematics in the asteroseismological mass determination methods in pulsating \pg\ stars.}{We compare the seismic masses resulting from the comparison of the observed mean period spacings with the usually adopted asymptotic period spacings, $\Delta \Pi_{\ell}^{\rm a}$, and the average of the computed period spacings, $\overline{\Delta \Pi_{\ell}}$. Computations are based on full PG1159 evolutionary models with stellar masses ranging from $0.530$ to $0.741 M_{\odot}$ that take into account the complete evolution of progenitor stars.} {We conclude that asteroseismology is a precise and powerful technique that determines the masses to a high internal accuracy, but it depends on the adopted mass determination method. In particular, we find that in the case of pulsating \pg\ stars characterized by short pulsation periods, like \ppg\ and \pp, the employment of the asymptotic period spacings overestimates the stellar mass by about 0.06 \msun\ as compared with inferences from the average of the period spacings. In this case, the discrepancy between asteroseismological and spectroscopical masses is markedly reduced when use is made of the mean period spacing $\overline{\Delta \Pi_{\ell}}$ instead of the asymptotic period spacing $\Delta \Pi_{\ell}^{\rm a}$.}{} | \label{intro} Pulsating PG1159 stars (or GW Virginis) are evolved hot stars that can pose constraints to the stellar evolution theory of post-Asymptotic Giant Branch (AGB). These variable stars belong to the population of hydrogen--deficient objects characterized by surface layers rich in helium, carbon and oxygen (Werner \& Herwig 2006) that are considered the evolutionary link between post-AGB stars and most of the hydrogen-deficient white dwarfs. The origin of most \pg\ stars is traced back to the occurrence of post--AGB thermal pulses: a born-again episode induced either by a very late thermal pulse (VLTP) experienced by a hot hydrogen-rich white dwarf during its early cooling phase --- see Herwig et al. (1999), Bl\"ocker (2001), Lawlor \& MacDonald (2003), Althaus et al. (2005), Miller Bertolami et al. (2006), or a late thermal pulse (LTP) during which hydrogen deficiency is a result of a dredge--up episode. (see Bl\"ocker 2001). During the VLTP, the convection zone driven by the helium flash reaches the hydrogen--rich envelope of the star, with the result that most of the hydrogen content is burnt. About a third of spectroscopic \pg\ stars exhibit multiperiodic luminosity variations with periods in the range $300-3000$ s, attributable to global nonradial $g$-modes pulsation (e.g. Quirion et al. 2007). The presence of a pulsational pattern in many \pg\ stars has allowed researchers to infer structural parameters --- particularly the stellar mass --- and the evolutionary status of individual pulsators --- e.g. Kawaler \& Bradley (1994), Kawaler et al. (1995), O'Brien et al. (1998), Vauclair et al. (2002) and more recently C\'orsico \& Althaus (2006). Stellar masses of PG1159 stars can independently be assessed by comparing the values of log $ g $ and \lteff, as inferred from detailed non--LTE model atmospheres (Werner et al. 1991), with tracks coming from stellar evolution modeling, i. e. the spectroscopic mass (Dreizler \& Heber 1998, Werner \& Herwig 2006). These two different approaches enable us to compare the derived stellar masses. Recently, considerable observational and theoretical effort has been devoted to the study of some pulsating \pg\ stars. Particularly noteworthy is the work of Fu et al. (2007) who have detected a total of 23 frequencies in \pp\ and Costa et al. (2007) who have enlarged to 198 the total number of pulsation modes in \pgs, making it the star with the largest number of modes detected besides the Sun. Parallel to these observational breakthroughs, substantial progress in the theoretical modeling of \pg\ stars has been possible (Herwig et al. 1999, Althaus et al. 2005, Lawlor \& Mac Donald 2006). In this sense, the new generation of \pg\ evolutionary models recently developed by Miller Bertolami \& Althaus (2006) (hereinafter MA06) has proved to be valuable at deriving structural parameters of pulsating \pg\ on the basis of individual period fits --- see C\'orsico et al. 2007a and C\'orsico et al. 2007b, respectively, for an application to the hot pulsating \rxj\ and the coolest member of the class, \pp. These evolutionary models are derived from the complete evolutionary history of progenitor stars with different stellar masses and an elaborate treatment of the mixing and extramixing processes during the core helium burning and born again phases. The success of these models at explaining both the spread in surface chemical composition observed in \pg\ stars and the location of the GW Vir instability strip in the $\log T_{\rm eff}- \log g$ plane (C\'orsico et al. 2006) renders reliability to the inferences drawn from individual pulsating \pg. As shown in MA06 the employment of detailed \pg\ evolutionary models yields spectroscopical masses that are systematically lower --- by about 0.05 \msun\ --- than those derived from hydrogen--rich post--AGB tracks (Werner \& Herwig 2006). Most importantly, the resulting asteroseismological masses (as inferred from the period spacings) are usually 10\% higher than the new spectroscopical masses, except for the hot pulsating PG1159 star \rxj, the spectroscopical mass of which is more than 20\% higher than the asteroseismological one (C\'orsico et al. 2007a). The mass discrepancy is a clear indication of the uncertainties weighting upon the mass determination methods, even though the spectroscopic uncertainties are of that order. In an attempt to understand the persisting discrepancy between the asteroseismological and spectroscopical masses, Miller Bertolami \& Althaus (2007) have recently shown that previous evolution is not the dominant factor in shaping hydrogen--deficient post--VLTP tracks. They conclude that the MA06 \pg\ tracks are robust enough as to be used for spectroscopical mass determinations of \pg--type stars, unless opacities in the intershell region are strongly subestimated. Their results make clear that the systematic discrepancy between asteroseismological and spectroscopical masses should not be attributed to uncertainties in post--AGB tracks; rather, they call for the need of an analysis of possible systematics in the asteroseismological mass determination methods. This is precisely the core feature of the present work. Specifically, we will concentrate on the usually adopted asymptotic period spacing approach (Kawaler et al. 1995; O'Brien et al. 1998; Vauclair et al. 2002; Fu et al. 2007) used in most mass determinations of individual pulsating \pg\ stars. The advantage of this approach lies in the fact that the mean period spacing of PG1159 pulsators depends primarily on the stellar mass (Kawaler \& Bradley 1994; C\'orsico \& Althaus 2006). However, the derivation of the stellar mass using the asymptotic predictions may not be entirely reliable because they are strictly valid for chemically homogeneous stellar models, while PG1159 stars are expected to be chemically stratified with strong chemical gradients built up during the progenitor star life. We will show that this approach overestimates the seismic mass for those pulsating \pg\ stars on the white dwarf cooling track. We will also show that the discrepancy between asteroseismological and spectroscopic masses is markedly alleviated if the average of the computed period spacings, instead of the asymptotic ones, is used. In the next Section, we summarize the seismological tools to infer the stellar mass from the observed mean period spacings. We also describe the evolutionary sequences employed. In Sect. 3 and \ref{period-spacing} we present our results and compare them with other mass determinations methods. We close the paper in Sect. \ref{conclusions} by summarizing our findings. | \label{conclusions} This paper explores the systematic discrepancy between spectroscopical and asteroseismological masses of pulsating \pg\ stars. Our motivation is the result of Miller Bertolami \& Althaus (2007) that such discrepancy should not be attributed to uncertainties in post--AGB tracks, but possibly to systematics in the asteroseismological mass determination methods. Recently, Quirion has pointed to one of us (M3B) that a possible opacity change resulting from the spread of He/C/O abundances in PG1159 stars could be a source of uncertainty in the location of the tracks. We addressed this issue by calculating sequences in which helium and carbon are changed in the whole envelope above the helium burning shell. We find that changing helium into carbon by an amount of 0.4 by mass shifts the track by only 0.02 dex in effective temperature (being bluer if carbon is higher). This translates into a shift of only 0.005 and 0.015 \msun\ for the spectroscopic mass near the 0.51 and 0.6 \msun\ tracks, respectively. Thus, the precise values of the He/C/O abundances do not seem to introduce appreciable changes in the masses derived by MA06. Specifically, we have concentrated on the seismic masses that result from a comparison of the observed period spacings with the usually adopted asymptotic period spacings ($\Delta \Pi_{\ell}^{\rm a}$) used in most mass determination of individual pulsating \pg\, and the better suited average of the computed period spacings ($\overline{\Delta \Pi_{\ell}})$. On the basis of full \pg\ evolutionary models that consider the evolutionary history of progenitor stars (MA06),and the ensuing internal chemical profile, we have shown that the derivation of the stellar mass using the asymptotic period spacing is not appropriate in the case of \pg\ stars. In particular, we demonstrate that for those pulsating \pg\ stars characterized by short pulsation periods, i. e., the pulsating \pg\ stars on the hot white dwarf regime (DOVs), the asymptotic $\Delta \Pi_{\ell}^{\rm a}$ differs appreciably (by more than 1 s) from the mean $\overline{\Delta \Pi_{\ell}}$. Only in the case of variables with long periods (PNNVs), like the high--luminosity, log--gravity pulsating \pg\ stars, do the $g-$ mode period spacings given by asymptotic $\Delta \Pi_{\ell}^{\rm a}$ resemble those predicted by mean $\overline{\Delta \Pi_{\ell}}$. This is expected because the asymptotic conditions are approached in the limit of very high radial order $k$. For quantitative inferences, we have computed the seismic mass resulting from the employment of the asymptotic and the average of the computed period spacing for those pulsating \pg\ which have a sufficiently large number of detected modes to infer an observed value of the mean period spacing. Our selected stars are listed in Table 1, together with the stellar mass inferences. The employment of the asymptotic theory, in principle formally valid for chemically homogeneous stellar models at high radial index k, overestimates the seismic mass by about 0.06 \msun\ in the case of very short period pulsating \pg\ stars like \ppg\ and \pp. Because \pg\ stars are expected to be chemically stratified, estimations of the stellar mass from mean $\overline{\Delta \Pi_{\ell}}$ are more realistic than those inferred by means of asymptotic $\Delta \Pi_{\ell}^{\rm a}$. Indeed, stellar mass derived from the mean $\overline{\Delta \Pi_{\ell}}$ are in good agreement with the mass values obtained from detailed period fittings. The discrepancy between asteroseismological and spectroscopical masses is markedly alleviated by the employment of the average of the computed period spacing instead of the asymptotic period spacings. In closing, a Fortran program to derive, from our evolutionary sequences, averages of the period spacing for arbitrary period intervals is available at our web site http://www.fcaglp.unlp.edu.ar/evolgroup. | 7 | 10 | 0710.3394 |
0710 | 0710.4057_arXiv.txt | {The shell-type supernova remnant RX J1713.7--3946 was observed during three years with the H.E.S.S. Cherenkov telescope system. The first observation campaign in 2003 yielded the first-ever resolved TeV gamma-ray image. Follow-up observations in 2004 and 2005 revealed the very-high-energy gamma-ray morphology with unprecedented precision and enabled spatially resolved spectral studies. Combining the data of three years, we obtain significantly increased statistics and energy coverage of the gamma-ray spectrum as compared to earlier H.E.S.S. results. We present the analysis of the data of different years separately for comparison and demonstrate that the telescope system operates stably over the course of three years. When combining the data sets, a gamma-ray image is obtained with a superb angular resolution of 0.06 degrees. The combined spectrum extends over three orders of magnitude, with significant gamma-ray emission approaching 100 TeV. For realistic scenarios of very-high-energy gamma-ray production, the measured gamma-ray energies imply efficient particle acceleration of primary particles, electrons or protons, to energies exceeding 100 TeV in the shell of RX J1713.7--3946.} | 7 | 10 | 0710.4057 |
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0710 | 0710.1444_arXiv.txt | One of the promising methods to search for life on extra-solar planets (exoplanets) is to detect its signature in the chemical disequilibrium of exoplanet atmospheres. Spectra at the modest resolutions needed to search for methane, oxygen, carbon dioxide, or water will demand large collecting areas and large diameters to capture and isolate the light from planets in the habitable zones around the stars. Single telescopes with coronagraphs to isolate the light from the planet will have to be 8\,m or more in diameter to generate sample sizes with a reasonable probability of finding at least one life-bearing planet; interferometers of smaller telescopes can overcome some of the limitations but will still need large similarly large collecting areas. Even larger telescopes will be needed to detect atmospheric signatures in transiting planets. In all cases, the sample sizes increase as the third power of telescope diameter. Direct observation using coronagraphs or interferometers are most sensitive to planets around stars with masses similar to that of the Sun, whereas transit observations favor low-mass stars near the nuclear burning limit. If the technical difficulties of constructing very large space telescopes can be overcome, they will be able to observe planets near hundreds to thousands of stars with adequate resolution and sensitivity to look for the signatures of life. | The detection of more than 200 planets outside the Solar System is a powerful incentive to search for extra-terrestrial life. Although extra-terrestrial life could take on many guises, economy of hypothesis (and a practical approach) implies that we should first search for signs of life similar to those seen on Earth. An advantage of using the Earth as a proxy for analysis is that it bounds the problem and provides concrete examples of signatures subject to passive detection, i.e. not requiring signals broadcast by sentient beings. In the Earth's present atmosphere, the chemical components have been altered by life. Evidence for life on Earth could be detected from afar in the spectral signatures of these molecules: oxygen, carbon dioxide, methane, and water vapor \citep{sag93}. The rise of photoplankton and plants on Earth created an atmosphere with a large reservoir of oxygen today that requires steady production by photosynthesis to maintain its present level. If we could study the same spectral signatures in the atmospheres of exoplanets, we could search for signs of life similar to some of the earliest and most robust forms on our own planet. But the Earth's atmosphere has been altered by life several times over the last 4\,Gyr, and there are potentially many signatures that could indicate the presence of life on exoplanets \cite{tra02, sea02, kalt07}. Chemical signatures of life on other planets would revolutionize our thinking about Earth's uniqueness and provide tantalizing evidence that we are not alone in the universe. Observing exoplanets directly is difficult owing to their proximity to the much brighter stars that keep them warm. Although technically challenging, this problem is well understood, and there are a variety of strategies that can reduce the brightness of the starlight without diminishing the light from the planet for direct detection \citep{guy06,cas06} or use the star itself as a background source to probe the atmosphere when the exoplanet transits the face of the star \citep{cha02, ehr06}. The first technique must overcome diffraction in the pupil of the telescope, a well understood phenomenon, and it is possible to characterize the detection problem in general terms to understand the kinds of instruments that will be needed to study exoplanets and search for signs of life. The second technique depends only on the photometric accuracy of an observation and is easy to calculate for any star. Any planet supporting life as on Earth must satisfy two broad criteria: (1) it must have surface temperatures in the range 273 to 373\,K, where water is in the liquid phase, and (2) it must have an atmosphere. The first criterion is met if the planet is in the {\it habitable zone} (HZ) around the star, a range of orbital distances where the equilibrium temperature for a rotating body is between the freezing and boiling points of water. The second criterion is met if the planet is rocky and can retain an atmosphere; current estimates specify a mass between 0.5 to 10 Earth masses. Smaller planets will not retain their atmospheres, and larger planets accrete gas and become gas giants. These criteria are necessary but not sufficient to create Earth-like life. Although they are probably far too restrictive to encompass all the possibilities for other life forms in the universe---or even on Earth itself---they are the only ones amenable to remote observation with technology that we can foresee at present and thus provide a good basis for a targeted search. There have been many calculations aimed at refining our ideas of the habitable zone \citep[e.g.][]{kast93}, suitable samples of stars to search \citep{tur03, tur04, tur06}, and the impact of specific telescopes on such a search \citep{alg07}, and there is some disagreement about the likelihood of success depending on the different assumptions used. Most authors concentrate on photometric detection alone, ignoring the means to search for life (e.g. Agol 2007). But the utility of photometric searches alone to identify exoplanets for subsequent study could be obviated by the difficulty of measuring the orbits accurately enough to recover them at a later time \cite{bro05, bro07}. For apertures sizes under discussion for the Terrestrial Planet Finder (TPF) mission, habitable zones of most potential target stars are blocked by the coronagraph. The planets of interest will spend the majority of time behind the central obscuration of the imaging instrument. To ensure efficient recovery of a newly discovered planet at future observing epochs, it will be necessary to estimate the orbit to high accuracy from a small number of astrometric measurements. This requirement implies a lower limit to the aperture size based on operational requirements (Brown et al. 2007). Achieving adequate astrometric precision may demand apertures larger than any so far discussed for TPF (Brown 2007, personal communication). If true, discovery and immediate spectroscopy of candidate sources near the stars may be the most efficient means of identifying those most interesting for follow up observations to determine if life is, indeed, present, making it essential to understand the spectroscopic requirements at the outset. The purpose of this article is to derive the main scaling parameters for the study of life-bearing exoplanets in known samples of stars to understand the size of the telescopes needed for a robust search. We adopt simple but optimistic assumptions to bound the problem and find the minimum size for survey telescopes. Using only the lowest order approximations and assuming ``best case'' observing conditions allows robust conclusions about the scale of facilities needed to tackle the search for life. The main premise is that direct photometric detection of exoplanets in a band where the exo-planet atmosphere is free of chemical signatures cannot be the endpoint of any mission to search for life-bearing planets; spectra of the atmospheres will be the major advance of observing the planets directly. Moreover, the rapid increase in the number of exoplanets discovered to date suggests that finding terrestrial planets will be easiest with indirect methods, such as observing radial velocity, photometric or astrometric variations in the host stars, and the real thrust of direct observations will be to search for signs of life. | 7 | 10 | 0710.1444 |
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0710 | 0710.1734_arXiv.txt | When two stars collide and merge they form a new star that can stand out against the background population in a starcluster as a blue straggler. In so called collision runaways many stars can merge and may form a very massive star that eventually forms an intermediate mass blackhole. We have performed detailed evolution calculations of merger remnants from collisions between main sequence stars, both for lower mass stars and higher mass stars. These stars can be significantly brighter than ordinary stars of the same mass due to their increased helium abundance. Simplified treatments ignoring this effect give incorrect predictions for the collision product lifetime and evolution in the Hertzsprung-Russell diagram. | In star clusters stars can experience close encounters with other stars, which in some cases can lead to the collision and merger of two or more stars. This is a possible formation mechanism for blue straggler stars. Blue stragglers are stars that appear on the extension of the main sequence in the colour-magnitude diagram (CMD) of star clusters (\citet{article:piotto_freq_bss} and figure \ref{fig:hrd_m67}). In the past, various mechanisms for their formation have been proposed (see \emph{e.~g.~} \citet{article:livio_blue_stragglers} for a list), but the most commonly accepted explanation is that they are stars that have gained mass long after they were formed, either through mass transfer in a close binary or by merging two stars. Such a merger can in turn be the result of normal binary evolution or of a collision with another star. These mechanisms are thought to operate and produce blue stragglers in different regions in a cluster \citep{article:davies_bss_formatio}. We have studied mergers that result from collisions. In some cases a sequence of collisions in the centre of a cluster can lead to a runaway where many stars merge together \citep{article:portegies_zwart_intermediate_mass_blackhole}. Such runaways may lead to the formation of very massive stars and ultimately intermediate mass blackholes. We have studied the evolution of the outcome of a single collision. One of the goals of the MODEST collaboration \citep{article:modest1} is to study the evolution of stellar mergers by combining stellar evolution, stellar hydrodynamics and stellar dynamics in one software framework. This work is the first result from combining stellar hydrodynamics and stellar evolution in one framework. | Stellar collisions are a means of making stars that appear in unusual locations in the colour magnitude diagram, such as blue stragglers. Remnants from high mass stars that merge after the end of core hydrogen burning can become much brighter while crossing the Hertzsprung-gap and stay there much longer. Approximating mergers with ordinary stellar models or fully mixed models can give significantly different results compared to proper detailed models. | 7 | 10 | 0710.1734 |
0710 | 0710.1028_arXiv.txt | We present new spectroscopic observations of the stellar cluster population of region~B in the prototype starburst galaxy M82 obtained with the Gillett Gemini-North 8.1-metre telescope. By coupling the spectroscopy with \emph{UBVI} photometry acquired with the Advanced Camera for Surveys (ACS) on the Hubble Space Telescope (HST), we derive ages, extinctions and radial velocities for seven young massive clusters (YMCs) in region B. We find the clusters to have ages between 70 and 200\,Myr and velocities in the range 230 to 350\kms, while extinctions $A_V$ vary between $\sim$1--2.5 mag. We also find evidence of differential extinction across the faces of some clusters which hinders the photometric determination of ages and extinctions in these cases. The cluster radial velocities indicate that the clusters are located at different depths within the disk, and are on regular disk orbits. Our results overall contradict the findings of previous studies, where region~B was thought to be a bound region populated by intermediate-age clusters that formed in an independent, offset starburst episode that commenced 600\,Myr--1\,Gyr ago. Our findings instead suggest that region B is optically bright because of low extinction patches, and this allows us to view the cluster population of the inner M82 disk, which probably formed as a result of the last encounter with M81. This study forms part of a series of papers aimed at studying the cluster population of M82 using deep optical spectroscopy and multi-band photometry. | The extensively studied galaxy M82 is a local example of a nuclear starburst galaxy. \citet[hereafter OM78]{om78} first catalogued the complex star-forming regions seen in ground-based images of the M82 disk, and introduced the nomenclature A--H. Region~B is the brightest region in the disk and is located 350--1050\,pc north-east of the nucleus. OM78 and \citet{marcum96} find that the integrated spectrum of this region is indicative of a fossil starburst region as it shows the characteristic `E$+$A' post-starburst spectrum, suggestive of a truncated burst of star formation occurring 100--1000\,Myr ago. Moreover, they find that the intrinsic brightness of region~B is such that the burst must have been of comparable intensity to the present starburst in the nucleus. More recently, de Grijs, O'Connell, \& Gallagher (2001) studied the M82-B cluster population using photometry obtained with the Wide Field and Planetary Camera 2 (WFPC2) and the Near-Infrared Camera and Multi-Object Spectrometer (NICMOS) onboard the Hubble Space Telescope (HST). They identified 113 clusters and, by estimating ages and extinctions from $BVI$ photometry, they found that a concentrated episode of cluster formation occurred 400--1000\,Myr ago with a peak at 600\,Myr. Subsequent to this analysis, de Grijs, Bastian, \&~Lamers~(2003a) re-derived ages and extinctions by combining $BVI$ and $JH$ photometry; they find that the peak of cluster formation occurred at 1.10\,Gyr with an age spread of 500--1500\,Myr. de Grijs, Bastian, \&~Lamers~(2003b) used these new ages to derive the cluster luminosity function (CLF) for a fiducial age of 1.0\,Gyr and find that it has a Gaussian shape, similar to globular clusters (GCs), rather than the standard power-law CLF found for populations of young massive clusters (YMCs) \citep[e. g. ][]{larsen04,gieles06a}. This result is surprising but supports theoretical models advocating that an initial power-law distribution of cluster masses will be transformed into a Gaussian distribution because low mass clusters will be preferentially disrupted \citep[e.g.][]{fall01, vesperini03, gieles06b}. \citet{RdG03-2} therefore suggest that because M82-B is of intermediate age, it provides the `missing evolutionary link' between the power-law CLFs found for YMCs and the Gaussian CLFs of globular clusters. However, for an environment as turbulent as that of M82, one would expect a timescale of several Gyr for preferential disruption to produce this effect \citep[e. g. ][where such a timescale is calculated for NGC 1316 and NGC 3610 respectively]{goud1316,goud3610}. To date, spectroscopy of the region~B cluster population has been extremely limited; the only published spectra are those of \citet{stis} for the two brightest members. The spectra were obtained with the Space Telescope Imaging Spectrograph (STIS) and, although they are of low quality, the derived ages of 350$\pm$100\,Myr are much lower than the photometrically-based ages of $\sim$0.7--6~Gyr \citep{stis, RdG03-1} for the same clusters. This discrepancy hints at the possibility that region~B may be younger and that the ages derived from the photometry are too high. It is important to determine accurate ages for the M82-B cluster population to verify its unusual CLF, and for understanding the cluster formation history of M82 and its relationship to encounters with its close neighbour M81. We therefore acquired new spectroscopic and photometric data for the M82-B cluster population. By using both techniques we are able to obtain information for a large number of clusters and also overcome the degeneracy between age and extinction that presents a hurdle in the analysis of photometric data \citep{gelys07a, gelys07b}. In this paper, we present optical spectroscopy for seven clusters obtained with the Gillett Gemini \mbox{8.1-m} telescope, and in a companion paper \citep{phot}, we present new HST imaging, including $U$-band photometry of the clusters. | We have presented Gemini GMOS spectroscopy and ACS \emph{UBVI} photometry for seven clusters in region~B of the starbust galaxy M82. Our aims were to obtain accurate ages, extinctions and radial velocities for the clusters. We have used both photometric and spectroscopic methods to determine cluster ages. Our main age determination method is fitting the Balmer spectral lines with all available models and finding the best fit (i.e. the lowest $\chi^2$ on the overall line profile fit), in order to eliminate the known Balmer line strength degeneracy. We find our method to agree with photometric ages derived using the 3DEF method (based on $UBVI$ colours). This demonstrates that the photometrically-derived ages are accurate but they are not as precise as the ones obtained from spectroscopy and confirms that the only truly independently reliable method for age-dating clusters in extragalactic environments is spectroscopy. The inclusion of $U$-band photometry may enhance the accuracy of photometric measurements, but unless they can be cross-checked by spectroscopic means, the potential age degeneracy may not be broken. We find cluster ages in the range of 80 to 200\,Myr, a distribution that is consistent with the timescale for the last encounter between M82 and M81 as proposed by \citet{yun99}, namely 220\,Myr. These findings in combination with the larger photometric sample of \citet{phot}, disagree with the `fossil starburst' scenario proposed by \citet{RdG01}, where region B in M82 was identified with the remnant of an off-centre starburst that commenced about 600\,Myr ago, a time scale based on their photometric determinations of star cluster ages. We suggest that the increased cluster formation rate in region~B is representative of the era of increased star formation across the galaxy disk triggered by the last encounter with M81. The extinction along our line of sight in this extended region appears to vary greatly, between 1.1 and 2.5 magnitudes. As M82 is seen virtually edge-on, we find that \mbox{M82-B} presents a view into various depths of the body of the galaxy through an arrangement of `windows' in the dust distribution \citep[as hypothesised for cluster M82-F by][]{smith01}. We also find differential extinction across the face of cluster \#91 (and also possibly \#108), which sets a serious obstacle in the photometric determination of extinction and age in these cases. In fact, the effect of dust in this environment is in some cases so grave that photometry cannot be used at all as an age/extinction diagnostic. We have also used the available spectroscopy to derive cluster kinematics. This allows us to reinforce our extinction-based findings and show that the seven clusters reside at different depths within the disk of M82. The large scatter of cluster velocities about the gently rising component of the rotation curve indicates that the clusters do not move in a co-ordinated fashion and that region B cannot be bound. Previous studies of clusters F and its neighbour L on the western side of the disk show that they both have an age of $60\pm20$ Myr \citep{GS99,bastian07}. This age fits in well with the region B cluster age distribution presented in this paper and \citet{phot}, and supports our suggestion that the increase in the cluster formation rate was not local to region B but part of a galaxy-wide burst. In addition, the very high masses of clusters F \citep[$\sim10^6$\,\Msun, ][]{smith01,bastian07} and L \citep[$4\times10^6$~\Msun, ][]{mccrady07} imply that many lower mass clusters were also likely to have formed with F and L. In summary, we find region~B to be optically bright owing to the presence of low internal extinction patches, thus offering a deep view into the M82 disk at radii between $\sim$ 400--1200\,pc. The range of cluster ages and other properties in this region then are typical of the evolution of the main body of M82 and reflect the large increase in star formation that occurred about 220\,Myr ago when M82 last passed close to M81. M82 region B stands out because it is representative of the mid-disk zone, and is relatively clear of dust, rather than being a special substructure. This model also receives support from the larger sample of photometric ages in Smith et al. (2007), and will be further discussed in the Konstantopoulos et al. (in preparation) study of another three dozen M82 star clusters distributed across M82. \\ | 7 | 10 | 0710.1028 |
0710 | 0710.5520_arXiv.txt | Virial mass is used as an estimator for the mass of a dark matter halo. However, the commonly used constant overdensity criterion does not reflect the dynamical structure of haloes. Here we analyze dark matter cosmological simulations in order to obtain properties of haloes of different masses focusing on the size of the region with zero mean radial velocity. Dark matter inside this region is stationary, and thus the mass of this region is a much better approximation for the virial mass. We call this mass the static mass to distinguish from the commonly used constant overdensity mass. We also study the relation of this static mass with the traditional virial mass, and we find that the matter inside galaxy-size haloes ($M\approx 10^{12}{\rm M}_{\sun}$) is underestimated by the virial mass by nearly a factor of two. At $z\approx 0$ the virial mass is close to the static mass for cluster-size haloes ($M\approx 10^{14}{\rm M}_{\sun}$). The same pattern -- large haloes having $M_{\rm vir} > M_{\rm static}$ -- exists at all redshifts, but the transition mass $M_0 = M_{\rm vir} = M_{\rm static}$ decreases dramatically with increasing redshift: $M_0(z) \approx 3\times 10^{15}h^{-1}{\rm M}_{\sun} (1+z)^{-8.9}$. When rescaled to the same $M_0$ haloes clearly demonstrate a self-similar behaviour, which in a statistical sense gives a relation between the static and virial mass. To our surprise we find that the abundance of haloes with a given static mass, i.e. the static mass function, is very accurately fitted by the Press \& Schechter approximation at $z=0$, but this approximation breaks at higher redshifts $z\simeq 1$. Instead, the virial mass function is well fitted as usual by the Sheth \& Tormen approximation even at $z\lesssim 2$. We find an explanation why the static radius can be 2--3 times larger as compared with the constant overdensity estimate. The traditional estimate is based on the top-hat model, which assumes a constant density and no rms velocities for the matter before it collapses into a halo. Those assumptions fail for small haloes, which find themselves in environment, where density is falling off well outside the virial radius and random velocities grow due to other haloes. Applying the non-stationary Jeans equation we find that the role of the pressure gradients is significantly larger for small haloes. At some moment it gets too large and stops the accretion. | The currently accepted paradigm of hierarchical clustering (\citealt{Whi78}, \citealt{Blu84}) provides a picture of assembly of dark matter haloes in which more massive haloes are formed through merging and accretion of smaller ones. This picture is supported for example by recent observations of ongoing mergers in clusters (e.g. \citealt{vDok99}, \citealt{Rin07}). The theoretical framework of the cold dark matter (CDM) has been proved successful in the description of the structure formation in the Universe from tiny fluctuations in the primordial density field (see \citealt{Pri03} for a review). This model has received support from many different observations, in particular of the gravitational lensing effect (\citealt{Smi01}, \citealt{Guz02}, \citealt{Kne03}, \citealt{Hoe04}, \citealt{She04}, \citealt{Man06}), CMB \citep{Spe07}, the abundance of clusters (\citealt{Pie01}, \citealt{Gla07}), and satellite dynamics (\citealt{Zar94}, \citealt{Pra03}). Cosmological simulations with ever increasing resolution play important role by making accurate predictions of different properties of dark matter haloes. Results of those simulations are used by other methods. For example, semi-analytical models of galaxy formation have been either incorporated into N-body simulations or use statistics such as halo mass function and merging trees, which were calibrated and tested using the simulations (e.g. \citealt{Som99,Cro06}). The simulations reveal important information about the internal structure of dark matter haloes (e.g. \citealt{Nav97}, \citealt{Bul01b}, \citealt{Tay01}) which is reflecting their underlying dynamics. Many results of simulations implicitly use some definition of what is the size and mass of a collapsed dark matter halo. The problem is that there is no well-defined boundary of a halo: density field is smooth around the halo. The common prescription for this boundary (and hence the mass belonging to the halo) is defined through the spherical collapse model (\citealt{Gun72}, \citealt{Gun77}). The size of the halo at redshift $z$ is given by the half of the radius of the spherical shell at turnaround which is collapsing at that redshift. This is the \textit{virial radius}. Thus, it is very common to measure the mass of haloes in cosmological simulations taking the particles inside a sphere of fixed spherical overdensity or to take those particles which are connected by a inter-particle separation below a given value (the friends-of-friends algorithm). We note that there is little justification for using the top-hat collapse model. Haloes do not collapse from perfect spherical homogeneous distribution. The environment of haloes is typically very non-spherical with most of accretion happening from few elongated filaments. The random velocities of the accreted matter also cannot be neglected. As the fluctuations collapse, the dark matter, which is being accreted, increases it rms velocities. This effective pressure should affect the accretion rate. The only motivation for using the top-hat model comes from simulations. Indeed, early simulations indicated that the radius of overdensity 178 is close to the virial radius \citep{Col96}: "the radius $r_{178}$ approximately demarcates the inner regions of haloes at $r\lesssim r_{178}$ which are in approximate dynamical equilibrium from the outer regions at $r\gtrsim r_{178}$ which are still infalling". Thus, the radius of overdensity 200 ($\approx 178$) became the virial radius. The reason why this was a good approximation is simply coincidental: the early simulations were mostly done for cluster-size haloes and, indeed, for those masses the virial radius is close to the radius of overdensity 200. The early models were flat models without the cosmological constant. Models with the cosmological constant have produced significant confusion in the community. The top-hat model must be modified to incorporate the changes due to the different rate of expansion and due to the different rate of growth of perturbations (see \citealt{Pri97}). That path produced the so called virial radius, which for the standard cosmological model gives the radius of overdensity relative to matter of about 340 \citep{Bul01}. Still, a large group of cosmologists uses the old overdensity 200 relative to the critical density even for the models with the cosmological constant. In this paper, we cast some light on this subject by searching for a physical extent of dark matter haloes, which is related to the physical processes that occur around collapsed structures. \citet{Pra06} provided first results, which indicated that spherically averaged, mean radial velocity profiles show an inner region in which there is no net infall or outflow. The size of this region in virial units is mass-dependent: for galactic-size haloes it may even reach three times the virial radius. We use this result to study the properties of the mass inside this region and in particular, to determine the redshift evolution of such mass. Recent effort (\citealt{Wec02}, \citealt{Zha03}) has been devoted to the analysis of the evolution of virial mass inside dark matter haloes, sometimes referred to as the mass accretion history. However, the picture presented in these works has been recently put into question. The analysis of the mass inside a fixed physical radius has revealed that galaxy-size haloes experience unphysical growth of their virial mass \citep{Die07}. The mean background density decreases as the Universe expands, making the virial radius to increase even in the case of no accretion or mergers. This is clearly an artifact of the definition. Other works realized this issue, for example, the evolution of the spin parameter when the matter inside a fixed radius is taken into account differs from that when using the evolving $R_{\rm vir}$ instead \citep{Don07}. On the other hand, there is another effect apart from accretion and merging, which has not been taken into account: haloes may also grow via relaxation of the surrounding regions near them. This is an important effect as it is a reflection of the dynamical processes which are turning non-virialized mass in the outskirts of a halo into the mass associated to it. Moreover, the long-term evolution of dark matter haloes show an interesting feature for $\Lambda$CDM cosmologies: their mass turns out to converge to an asymptotic value which depends on the definition of mass, as pointed out by \citet{Bus05}. Thus, it is desirable that the mass of a halo is measured using a virialization-based criteria (see \citealt{Mac03} for an interesting approach), instead of using boundaries of a given overdensity. The interest of the measurement of the physical mass associated to dark matter haloes is not only theoretical. Indeed, it may have a great impact on the number of collapsed objects in a given range of mass, i.e. the mass function \citep{Whi01}. Besides, it is of great importance for the process of formation and evolution of galaxies, and, hence, it is relevant for the results obtained from semi-analytical modelling of galaxy formation (e.g. \citealt{Cro06}), to predict the main properties of observed galaxies. Thus, it turns out to be mandatory to bring attention on the mass belonging to a halo, if accurate predictions of the physics behind galaxies from cosmological simulations are to be drawn. This paper is organized as follows: in Section~\ref{sec:simul} we describe the set of cosmological simulations used in our analysis and the properties of the halo samples. In Section~\ref{sec:static} we present the definition of static mass and how it is related to objects of different sizes. A brief analysis of equilibrium in this context is presented in Section~\ref{sec:jeans}. The scaling properties of the static to virial mass relation with redshift are shown in Section~\ref{sec:scaling}, together with a simple model derived from this scaling relation. We obtain the static mass function and compare it with analytical models in Section~\ref{sec:mf}. In Section~\ref{sec:evolution} we present the evolution of the static mass tracking the halo progenitors. We discuss our results in Section~\ref{sec:discussion} and present our conclusions in Section~\ref{sec:conclusion}. | \label{sec:conclusion} In this paper we use a set of high resolution $N$--body cosmological simulations for the analysis of the mass inside the region of dark matter haloes with no net infall or outflow velocities, i.e. the static region. Haloes show a typical radial velocity pattern which depends on their mass: low-mass haloes tend to display a region with outflow in their surroundings, while cluster-size haloes show a prominent infall velocity pattern. Galaxy-size haloes show an intermediate pattern with a sharper transition between the static region and the Hubble flow. We find that the former virial radius tends to underestimate the size of the region with zero mean radial velocity for haloes with masses $10^{10}h^{-1}{\rm M}_{\sun}<M_{\rm vir}\lesssim 10^{14}h^{-1}{\rm M}_{\sun}$ at $z=0$. The mass inside this static region can be about a factor of two larger than the former virial mass, when the threshold defining the static radius is $5$ per cent of the virial velocity $V_{\rm vir} =\sqrt(GM_{\rm vir}/R_{\rm vir})$. Lower values of the threshold ($\lesssim 1$ per cent) may lower this factor. However, we observe a clear trend in the $M_{\rm static}/M_{\rm vir}$ vs $M_{\rm vir}$ relation which remains unchanged with this threshold: for low-mass haloes, this ratio is an increasing function of $M_{\rm vir}$ due to reduction in the outflow. On the other hand, high-mass haloes show a larger infall for larger $M_{\rm vir}$, making this ratio to decrease. The maximum occurs at $M_{\rm vir}\simeq 5\times 10^{12}h^{-1}{\rm M}_{\sun}$, where the size of the static region is the largest in units of $R_{\rm vir}$. The mass function of objects whose mass is defined in the way presented here resembles that of Press \& Schechter in the range of mass we have studied, but only at $z=0$. At higher redshifts the static mass function deviates significantly from it. This disagreement could be accounted for by the scatter in $M_0(z)$, although it seems more likely that the resemblance with Press \& Schechter at $z=0$ is due to coincidence (in both shape and amplitude to within $\sim 20\%$) of the ratio $f_{PS}/f_{ST}$ with $M_{\rm static}/M_{\rm vir}$. Whatever the case may be, this resemblance is theoretically unexpected because the Sheth \& Tormen function, which fits very well the mass function for virial masses, is derived from a more realistic approach (ellipsoidal collapse). In any case it seems clear that when the virial mass is used, the number density of dark matter haloes of a given mass at $z=0$ might be underestimated in the mass range $10^{10}h^{-1}{\rm M}_{\sun}<M_{\rm vir}\lesssim 10^{14}h^{-1}{\rm M}_{\sun}$. The redshift evolution of the $M_{\rm static}$--$M_{\rm vir}$ relation turns out to be very weakly redshift-dependent from $z=2$ to $z=0$ when appropriate variables are used. This dependence on the redshift is encoded as an evolving mass scale $M_0$ which indicates the approximate mass at which $M_{\rm static}=M_{\rm vir}$ in the declining part of the relation. By doing this, it is straightforward to derive from this relation a differential equation for the evolution of the static mass, provided that the evolution of $M_{\rm vir}$ and of $M_0$ are already known. As we have shown in this paper, the evolution of the mass scale has an acceptable fit to a power law: $ M_0\simeq 2.8\times10^{15} a^{8.9} h^{-1}{\rm M}_{\sun}$ . On the other hand, the evolution of the virial mass $M_{\rm vir}$ of the major progenitor is well described by an exponential law in $z$, as claimed by \citet{Wec02}. Once we have fixed the evolution of $M_{\rm vir}$ and of $M_0$, we solve this differential equation and we find a simple model for the evolution of the static mass of the major progenitor, i.e. $M_{\rm static}(z) =M_{\rm static}(z=0)(1+z)^{-\beta}e^{-\alpha z}$, which turns out to be a very good fit to our data. In any case, there is a different behaviour for the evolution of the static mass in different mass bins. The static mass of low mass haloes is nearly constant from $z=0.5$ to $z=0$. Galaxy-size haloes keep growing at present, but the reason here is not the decreasing background density of the expanding Universe. Instead, the relaxation of the matter in the surroundings of these haloes is incorporating mass at even a higher rate than the unphysical growth of the virial mass. The same stands for clusters, although the static and virial mass are not very different at present. The mean radial velocity profile (averaged over a halo mass bin) evolves in the way found by \citet{Bus05} for cluster-size haloes: while the halo is forming, it accretes mass from its surroundings via infall, which is being reduced with time. This is followed by a period of outflow of unbound matter, and a final configuration thereafter consisting in a static region whose size is independent of time. Although the simulation used in that paper, was evolved until $a=100$, and our halo sample was instead tracked up to redshift $z=0$, we find a situation that is consistent with this picture. Moreover, we find that the timescale of this sequence is different for each mass bin: low-mass haloes at present are in the outflow phase, but cluster-size haloes are still in their infall phase. Interestingly, galactic-size haloes are experiencing a transition (infall to outflow) epoch at $z=0$. A.J.C., F.P. thank the Spanish MEC under grant PNAYA 2005-07789 for their support. A. K. acknowledges support from NASA and NSF grants to NMSU. A.J.C. acknowledges the financial support of the MEC through Spanish grant FPU AP2005-1826. We thank Juan Betancort-Rijo for valuable discussion and useful comments in this work, and Douglas Rudd for suggesting some points which improved the text. We do appreciate the useful feedback from the anonymous referee. A.J.C. also thanks the Helmholtz Institute for Supercomputational Physics for their outstanding Summerschool 2006 in Supercomputational Cosmology. Computer simulations were done at the LRZ Munich, NIC Julich, and NASA Ames. | 7 | 10 | 0710.5520 |
0710 | 0710.2246_arXiv.txt | The strong lensing modelling of gravitational ``rings'' formed around massive galaxies is sensitive to the amplitude of the external shear and convergence produced by nearby mass condensations. In current wide field surveys, it is now possible to find out a large number of rings, typically 10 gravitational rings per square degree. We propose here, to systematically study gravitational rings around galaxy clusters to probe the cluster mass profile beyond the cluster strong lensing regions. For cluster of galaxies with multiple arc systems, we show that rings found at various distances from the cluster centre can improve the modelling by constraining the slope of the cluster mass profile. We outline the principle of the method with simple numerical simulations and we apply it to 3 rings discovered recently in Abell~1689. In particular, the lens modelling of the 3 rings confirms that the cluster is bimodal, and favours a slope of the mass profile steeper than isothermal at a cluster radius $\sim 300 \kpc$. These results are compared with previous lens modelling of Abell~1689 including weak lensing analysis. Because of the difficulty arising from the complex mass distribution in Abell~1689, we argue that the ring method will be better implemented on simpler and relaxed clusters. | In recent years, the modelling of cluster mass distribution has been progressively improved by (a) coupling the strong lensing (SL) analysis in cluster cores with weak lensing (WL) measurements at large radius (e.g Gavazzi 2005\nocite{gavazzi05}; Limousin et al. 2007b\nocite{limousin1689}; Cacciato et al. 2006 \nocite{cacciato06}); (b) SZ measurements (e.g. Zaroubi et al. 2000 \nocite{zaroubi01}; Dore et al. 2001\nocite{dore01}; Seren 2007\nocite{Seren07}), (c) joint modelling of the cluster X-ray gas distribution (e.g. Mahdavi et al. 2007\nocite{Mahdavi07}), (d) dynamical analysis of the velocity distribution of the stars near the potential centre (e.g. Miralda-Escud\'{e} 1995\nocite{miraldababul95}; Sand et al. 2004 \nocite{Sand04}; Gavazzi 2005\nocite{gavazzi05} and Koopmans et al. 2006\nocite{Koopmans06}). Despite these important improvements, it is not fully proved that cluster Dark Matter (DM) distribution closely follows the ``universal'' profile (Navarro, Frenk \& White 1996\nocite{nfw}, hereafter NFW) predicted by N-body numerical simulations. For this profile, the DM space density is cuspy at the centre ($r \propto r^{-1}$ for $r \la r_s$) and behaves as $r^{-3}$ outward. Measuring with great accuracy these two main characteristics is challenging and no consensus has arisen yet with the present day lensing observations. Recent analysis has shown that S\'{e}rsic profile (Merritt et al. 2005\nocite{Merritt05}) is also fitting the universal mass profile of numerical simulations. Importantly, the mass profile has only a weak dependence on the halo mass or cosmology, allowing to stack different measurements together to improve their significance. At large radius, analyses of the hot gas distribution from the X-ray observations (Pointecouteau 2005\nocite{Pointecouteau05}; Schmidt 2007\nocite{Schmidt07}), as well as weak lensing (Kneib et al. 2003)\nocite{kneib03} seem to favour an NFW-like profile for galaxy clusters. For elliptical galaxies, Wilson et al. (2001\nocite{Wilson01}) claimed that they were consistent with isothermal profile out to about $1 \Mpc$. Also, for massive elliptical galaxies with gravitational rings, Gavazzi et al. (2007) \nocite{Gavazzi07} found that the WL slope of the mass density could be $\propto r^{-2}$ out to $300 \kpc$. In the core of mass concentrations, the actual existence of a DM cusp predicted by simple numerical simulations is more unclear. For spiral galaxy halos, the existence of a singular density profile is still a debate because rotation curves are better explained by isothermal profiles with a core radius (Salucci 2003\nocite{Salucci03}). Projection effects of non circular star orbits in triaxial halos has been invoked, but only in a few cases, to explain the linear increase of the velocity at the centres of galaxies (Hayashi et al. 2006 \nocite{Hayashi06}). As a consequence, various mechanical processes, such as gas cooling, supernova and AGN feedback, binary super massive black holes, dynamical friction of the gas outflow during AGN activities, were investigated to explain the formation of a DM core radius (Peirani et al. 2006\nocite{Peirani06} and references therein). For galaxy clusters, it is often argued that isothermal ellipsoid mass distributions with a flat core could better match the gravitational arcs geometry than NFW profiles (Sand et al. 2004\nocite{Sand04}; Gavazzi 2005\nocite{gavazzi05}; Gavazzi et al. 2003\nocite{Gavazzi03}). X-ray observations generally do not help much for this issue because, due to the limited spatial resolution of X-ray telescopes, one can just place upper limits on the radius of the smallest DM core, around $30-50\kpc$ (Chen et al. 2007\nocite{Chen07}). In summary, the DM density profile is still an open question, and high quality data are necessary to settle these questions. The deviation of light by masses is well described by gravitational lensing effect deduced from general relativity theory. The exquisite {\it Hubble Space Telescope} images (particularly from the now defunct ACS camera) are providing the necessary lensing constraints to model observed gravitational lensing systems and may directly probe the existence of a ``universal'' density mass profile. However, one has to struggle to fully take into account the many observational parameters entering a lens modelling. Firstly, it is difficult to assess the stellar mass contribution because the stellar mass-to-light ratio ($M/L$) is generally badly determined, as well as the number of sub-halos and their galaxy occupation numbers (Wright et al. 2002\nocite{Wright02}). Secondly, lens modelling only probes the projected mass distribution of lenses along the line of sight which introduces degeneracies in the 3D density profile if the mass profile is only determined on a small range of radius. When testing parametric models of mass distributions for a given DM condensation, these current difficulties are only partly alleviated if we can probe the projected mass at many different radius. As an example, to disentangle between a flat core or a cusp with strong lensing it is not enough to detect and to analyse gravitational images very close to the centre, the so-called demagnified central images (Gavazzi et al. 2003\nocite{Gavazzi03}) or inner radial arcs (Mellier et al. 1993\nocite{Mellier93}; Comerford et al. 2006\nocite{comerford06}). Even in such ideal cases, some information on the mass distribution beyond the Einstein radius is also critically needed. Beyond giant arc radii, one can use the information in the distortion of singly highly magnified arclets, in an intermediate shear regime ($2 < \mu < 3$), also called flexion regime (Bacon et al. 2006\nocite{Bacon06}; Massey et al. 2007\nocite{Massey07}). However, a generic difficulty similar to the one encountered in the weak lensing needs to be overcome. We do not know the shape of background sources and the flexion method must be used in a statistical way. Only with the most deepest space-based observations, it becomes possible to reach surface number density of background galaxies large enough to conduct such analysis on a single cluster (see recent work on Abell~1689 by Leonard et~al.(2007)\nocite{Leonard07}, where they reach a density of background sources equal to $\sim 200$ sources/arcmin$^2$). From the above discussion, we understand the difficulty to conduct an accurate measurement of the slope of the 3D density profile at large radial distances. Hence considering new probes of cluster mass profile is important. In this paper, we propose a method to investigate the slope of cluster mass profiles. Gravitational image systems, i.e. ``rings'' formed around galaxy cluster members, are used to analyse the slope of the cluster's density profile. Nowadays, such rings can be systematically searched with dedicated software (e.g. Gavazzi et al. 2007, in preparation; Cabanac et al. 2005). Here, we outline the method with three rings detected around Abell~1689. The main goal is to provide some constraints on the cluster potential at the location of the rings (i.e. at $\sim$ 100$\arcsec$ from the centre of the brightest cluster galaxy). The coordinate system in this work is centred on the brightest cluster galaxy: $\alpha_{J2000}=13:11:29.52, \delta_{J200}=-01:20:27.59$. The paper is organised as follows. First we rapidly summarise the properties of gravitational rings observed in the field around elliptical lenses. In section~\ref{simulation}, we illustrate the method by using simulated cluster profiles, lensing galaxies, and resulting images. These simulations show that we can put constraints on the local slope of the projected mass distribution. In section~\ref{reallife}, the method is applied to Abell~1689. In this case we show that the three rings confirm that the cluster is dominated by a bimodal mass distribution and that the local slopes of both clumps are not much steeper than isothermal. Then, the results are discussed relatively to previous lensing models of Abell~1689, including a weak lensing analysis in the field of the rings. Finally, we conclude that this method should be better used on very relaxed clusters (single halo) with regular geometry to better probe the slope of the mass profile at various distances from the cluster centre. Throughout this paper we assume a cosmological model with $\Omega_{\rm{m}}$=0.3, $\Omega_\Lambda$=0.7, $H_0=70$km s$^{-1} \Mpc^{-1}$. At the redshift of the cluster Abell~1689 ($z=0.185$) $1 \arcsec$ is equivalent to 3.089\kpc. | This work revisits the importance of external shear perturbations for the modelling of galaxy lenses (see Dye et al. 2007\nocite{Dye07}). But instead of focusing our attention on the improvement of the ring modelling by introducing an external shear, we probe the gravitational potential slope in the outskirts of groups and clusters of galaxies thanks to the high sensitivity of the ring modelling both to the local external shear and convergence of the cluster. We test the principle of the method with simple simulations using power law models to describe the cluster mass clumps. From this simulation we found that within the astrometric accuracy of ACS images a ring can be used to detect a logarithmic slope of the external mass density distribution up to about $n=2.8$. For an NFW profile with a typical concentration parameter $C_{\rm{vir}} \sim 5$ (Comerford \& Natarajan 2007\nocite{Comerford07}), this should allow to probe the cluster potential up to a large fraction of its virial radius. It is now possible to detect rings in the outskirts of clusters with new ground based surveys like CFHT-LS, but the determination of the direction and amplitude of the external shear depends crucially on the high resolution HST images. On the other hand, the knowledge of the lens stellar velocity dispersion and of the redshift of the ring itself would increase the robustness to the ring modelling (Koopmans et al. 2006). Using the 3 rings found in Abell~1689 we have shown that it is possible to estimate the mass density slope. Hence, we have shown with the information given by rings alone that the potential of Abell~1689 is bimodal. With this limited number of rings it is however not possible to better trace other DM substructures within each clump. Therefore a global modelling of the 31 arcs systems and rings remains a good challenge. We find a mass density slope a bit larger than $n=2$ for each DM clump at about $100\arcsec$ from their centre, but the complexity of the cluster potential really weakens the result. Since the most immediate appealing application of the method is to probe the slope departure from a $r^{-2}$ mass density profile, we strongly suggest similar analysis on less complex clusters. The ideal lenses to probe the reality of a universal NFW profile would be a cluster or a fossil group with a bright cD galaxy having a single dominant halo, and which displays at least a distant ring, a multiple arc with known redshift, and a radial arc. Such a configuration would probe the potential at various distances from the centre. The discovery of such ``golden lenses'' requests a large survey of massive clusters. Most often clusters with a great number of arcs have multi-polar potentials (i.e. longer caustics) so that are more complex to analyse. However, rings are influenced by all their nearby environment (including foregrounds). In this sense, a statistical study of the environmental effect on the rings should be conducted in the future. Many rings can be discovered now in the field of deep wide field surveys. It will be possible to improve the method and to use it more systematically. As an example, the CFHT-LS ``wide'' survey will cover 170 degree , and we expect to detect about 10 rings per square degree with an average redshift $z_{\rm{lens}} = 0.65$ (Cabanac et al. 2007\nocite{Cabanac07}). The clusters and groups in front of or at the same redshift as the lensing galaxies have approximately the same sky density (Oguri 2006\nocite{Oguri06}). Thus we can expect that several rings per square degree will be influenced by the external shear of a nearby mass condensation in the field. Such cases are already observed in the SL2S survey (Cabanac et al. 2007\nocite{Cabanac07}). For Abell~1689, we have found three rings within the small HST field ($r \sim 100\arcsec$), which may be consider as surprisingly high number. In fact, the large number of ellipticals in this cluster and the extra convergence that it adds to the lensing effect most probably compensate for the small field of view imaged with the ACS camera\footnote{After the submission of this paper, King (2007\nocite{king07}) has shown with cluster simulation that the ring cross section is increased by a factor about 3 near critical lines. }. We also found two possible rings closer to the critical region of Abell~1689 ({\it i.e.} at radii smaller than 50$\arcsec$). However, we are not considering these rings here, because these systems are in a crowded region and can not be described by a simple contribution of a galaxy and a cluster scale component. Any attempt of modelling these latter rings must involve all the multiple arc systems of Abell~1689. In wide field surveys, perturbing clusters may be found at a lower redshift than the ring-producing elliptical lenses (e.g. Smail et al. 2007\nocite{smailcosmiceye}). In such cases, a proper analysis requires a multi-plan ring modelling. In conclusion, rings seem to be a promising tool to constrain the mass distribution slope at large radii from the centres of groups and clusters provided dedicated observations are carried out on a few golden lenses. They can provide information on many structural parameters of halos of clusters and galaxies. In complement to the modelling of the multiple arcs in the cluster core, they can confirm or otherwise dispute the existence of the universal DM halo profile predicted by numerical simulation, as well as to study sub-halos and their cut-off radii. For these studies, it is crucial to measure the velocity dispersion of the lens galaxies which produce the gravitational ring images. It clearly appears from the study of Abell~1689 that this method should be implemented first on a sample of very relaxed clusters or (fossil) groups in order to analyse a single DM potential with the simplest possible geometry. Only when large sample of rings is available, it will become possible to start a systematic analysis of the external shear perturbation on ring shapes in correlation with more complex nearby (eventually foreground) environment. | 7 | 10 | 0710.2246 |
0710 | 0710.5466_arXiv.txt | {A variety of formation scenarios have been proposed to explain the diversity of properties observed in bulges. Studying their intrinsic shape can help to constraint the dominant mechanisms at the epochs of their assembly.} {The structural parameters of a magnitude-limited sample of 148 unbarred S0--Sb galaxies were derived in order to study the correlations between bulges and disks, as well as the probability distribution function of the intrinsic equatorial ellipticity of bulges.} {We present a new fitting algorithm (GASP2D) to perform two-dimensional photometric decomposition of the galaxy surface-brightness distribution. This was assumed to be the sum of the contribution of a bulge and disk component characterized by elliptical and concentric isophotes with constant (but possibly different) ellipticity and position angles. Bulge and disk parameters of the sample galaxies were derived from the $J-$band images, which were available in the Two Micron All Sky Survey. The probability distribution function of the equatorial ellipticity of the bulges was derived from the distribution of the observed ellipticities of bulges and misalignments between bulges and disks.} {Strong correlations between the bulge and disk parameters were found. About $80\%$ of bulges in unbarred lenticular and early-to-intermediate spiral galaxies are not oblate but triaxial ellipsoids. Their mean axial ratio in the equatorial plane is $\langle B/A \rangle = 0.85$. Their probability distribution function is not significantly dependent on morphology, light concentration or luminosity. The possible presence of nuclear bars does not influence our results.} {The interplay between bulge and disk parameters favors scenarios in which bulges have assembled from mergers and/or have grown over long times through disk secular evolution. However, all these mechanisms have to be tested against the derived distribution of bulge intrinsic ellipticities.} | \label{sec:introduction} The relative prominence of galactic bulges with respect to their disks is important in the definition of galaxy types. Therefore, understanding the formation of bulges is key to understanding the origin of the Hubble sequence. Bulges are diverse and heterogeneous \citep[see the reviews by][]{kormendy93,wyse97,kormendy04}. The big bulges of lenticulars and early-type spirals are similar to low-luminosity elliptical galaxies. Their surface-brightness radial profiles generally follow a De Vaucouleurs law \citep[][hereafter MH01]{andredakis95,carollo98,mollenhoff01}. The majority of these bulges appear rounder than their associated disks \citep{kent85}. Their kinematical properties are well described by dynamical models of rotationally flattened oblate spheroids with little or no anisotropy \citep{kormendy82,davies83,cappellari06}. They have photometrical and kinematical properties, which satisfy the fundamental plane (FP) correlation \citep{bender92,bender93,burstein97,aguerri05a}. On the contrary, the small bulges of late-type spiral galaxies seems to be reminiscent of disks. Their surface-brightness radial profiles have an almost exponential falloff \citep{andredakis94,dejong96,macarthur03}. In some cases they have apparent flattenings that are similar or even larger than their associated disks \citep{fathi03} and rotate as fast as disks \citep{kormendy93,kormendy02}. Late-type bulges deviate from the FP \citep{carollo99}. Different formation mechanisms (or at least a variety of dominant mechanisms at the epochs of star formation and mass assembly) were proposed to explain the variety of properties observed in bulges. Some of these formation processes are rapid. They include early formation from the dissipative collapse of protogalactic gas clouds \citep{eggen62,sandage90,gilmore98,merlinchiosi06} or later assembly from mergers between pre-existing disks \citep{kauffmann96,baugh96,cole00}. In both scenarios the disk forms after the bulge as a consequence of either a long star-formation time compared to the collapse time or a re-accretion around the newly formed bulge. Bulges can also grow over long timescales through the disk secular evolution driven by bars and/or environmental effects. Bars are present in more than half of disk galaxies in the local universe \citep{eskridge00,menendez-delmestre07} and out to $z\sim1$ \citep{elmegreen04,jogee04}. They are efficient mechanisms for driving gas inward to the galactic center and feed the galactic supermassive black hole \citep[see][and references therein]{corsini03a}. In addition, bar dissolution due to the growth of a central mass \citep{pfenninger90}, scattering of disk stars at vertical resonances \citep{combes90}, and coherent bending of the bar perpendicular to the disk plane \citep{raha91,debattista04,athanassoula05,martinez-valpuesta06} are efficient mechanisms in building central bulge-like structures, the so-called boxy/peanut bulges. Moreover, the growth of the bulge out of disk material may also be externally triggered by satellite accretion during minor merging events \citep{searlezinn78,aguerri01,eliche-moral06} and gas infall \citep{thakarryden98}. Traditionally, the study of the relations between the structural parameters of the galaxies have been used to understand the bulge formation processes, e.g., the correlation between the bulge effective radius and the scale length of the disk in many galaxy samples has always been interpreted as an indication that bulges were formed by secular evolution of their disks \citep[see][]{macarthur03}. However, one piece lost in this study is the three-dimensional shape of the bulges. By studying this, one might be able to provide the relative importance of rapid and slow processes in assembling the dense central components of disk galaxies. A statistical study can provide a crucial piece of information for testing the results of numerical simulations of bulge formation for different galaxy type along the morphological sequence. In this paper, we analyze a sample of unbarred early-type disk galaxies to derive the intrinsic ellipticity of their bulges in the galactic plane. The twisting of bulge isophotes \citep{lindblad56,zaritsky86} and misalignment between the major axes of the bulge and disk \citep{bertola91} are not possible if the bulge and disk are both oblate. Therefore, they were interpreted as a signature of bulge triaxiality. This idea is supported by the presence of non-circular gas motions \citep[e.g.,][]{gerhardvietri86,bertola89,gerhard89,berman01} and a velocity gradient along the galaxy minor axis \citep{corsini03b,coccato04,coccato05}. We improve the previous works in several aspects. First, we use near-infrared images to map the distribution of the mass-carrying evolved stars and avoid contamination of dust and bright young stars. Second, we retrieve the structural parameters of the bulge and disk by applying a new algorithm for two-dimensional photometric decomposition of the observed surface-brightness distribution. Finally, we obtain the probability distribution function (PDF) of the intrinsic equatorial ellipticity of bulges by using a new mathematical treatment of the equations describing their three-dimensional shape. The paper is organized as follow. The selection criteria of our sample galaxies and the analysis of their near-infrared images are described in Sect. \ref{sec:sample}. Our new photometric decomposition method for deriving the structural parameters of the bulge and disk by analyzing the two-dimensional surface brightness distribution of galaxies is presented in Sect. \ref{sec:decomposition}. The correlations between the structural parameters of the sample galaxies are discussed in Sect. \ref{sec:correlations}. The PDF of intrinsic equatorial ellipticity of the studied bulges is derived in Sect. \ref{sec:ellipticity}. Our conclusions and a summary of the results are given in Sect. \ref{sec:discussion}. | \label{sec:discussion} The structural parameters of the bulge and disk of a magnitude-limited sample of 148 unbarred S0--Sb galaxies were investigated to constrain the dominant mechanism at the epoch of bulge assembly. \begin{itemize} \item We presented a new fitting algorithm (GASP2D) to perform two-dimensional photometric decomposition of galaxy images. The surface-brightness distribution of the galaxy was assumed to be the sum of the contribution of a S\'ersic bulge and an exponential disk. The two components were characterized by elliptical and concentric isophotes with constant (but possibly different) ellipticity and position angles. GASP2D is optimized to deal with large image samples, and it adopts a robust Levenberg-Marquard fitting algorithm in order to obtain reliable estimates of the galaxy structural parameters. \item The bulge and disk parameters of the sample galaxies were derived from the $J-$band images, which were available in the Two Micron All Sky Survey. \item The bulges of the sample galaxies follow the same FP, FJ, and PP relationships found for elliptical galaxies. No statistically significant difference is observed when only bulges of lenticular and early-to-intermediate spiral galaxies were considered. This supports the idea that bulges and ellipticals formed in the same way. \item Tight correlations between the parameters of bulges and disks were found. In fact, the disk scale lengths increase with both the central velocity dispersion and bulge effective radius. Therefore, larger disks reside in galaxies with more massive and larger bulges. This was interpreted as an indication of the formation of bulges via secular evolution of their host disks. \item Our measurements of the exponential scale length of the bulge and disk, as well as of bulge shape parameter, were also fully consistent with numerical simulations of the effects of mergers on the mass distribution of the bulge and disk in galaxies formed in hierarchical clustering scenarios. \item These results indicate that the above relations are not enough to clearly distinguish between bulges formed by early dissipative collapse, merging or secular evolution. All these mechanisms could be tested against the intrinsic shape of bulges. Therefore, the PDF of the intrinsic equatorial ellipticity of the bulges was derived from the distribution of the observed ellipticities of bulges and their misalignments with disks. \item About $80\%$ of bulges in unbarred lenticular and early-to-intermediate spiral galaxies are not oblate but triaxial ellipsoids. Their mean axial ratio in the equatorial plane is $\langle B/A \rangle = 0.85$. This is consistent with previous findings by Bertola et al. (1991) and Fathi \& Peletier (2003). There is no significant dependence of the PDF on the morphology, light concentration, and luminosity of bulges. The derived PDF is independent of the possible presence of nuclear bars. \end{itemize} | 7 | 10 | 0710.5466 |
0710 | 0710.0844_arXiv.txt | We present a complete set of diagnostic tools aimed at reproducing synthetic non-thermal (synchrotron and/or Inverse Compton, IC) emissivity, integrated flux energy, polarization and spectral index simulated maps in comparison to observations. The time dependent relativistic magnetohydrodynamic (RMHD) equations are solved with a shock capturing code together with the evolution of the maximum particles energy. Applications to Pulsar Wind Nebulae (PWNe) are shown. | 7 | 10 | 0710.0844 |
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0710 | 0710.2888_arXiv.txt | Two-temperature, two-mass quasi-equilibrium plasmas may occur in electron-ion plasmas, nuclear-matter, as well as in electron-hole condensed-matter systems. Dense two-temperature hydrogen plasmas straddle the difficult partially-degenerate regime of electron densities and temperatures which are important in astrophysics, in inertial-confinement fusion research, and other areas of warm dense matter physics. Results from Kohn-Sham calculations and QMC are used to benchmark the procedures used in classical molecular-dynamics simulations, HNC and CHNC methods to derive electron-electron and electron-proton pair-distribution functions. Then, nonequilibrium molecular dynamics for two-temperature, two-mass plasmas are used to obtain the pair distribution. Using these results, the correct HNC and CHNC procedures for the evaluation of pair-distribution functions in two-temperature two-mass two-component charged fluids are established. Results for a mass ratio of 1:5, typical of electron-hole fluids, as well as for compressed hydrogen are presented. | The study of hot strongly-coupled dense charged fluids is a difficult task, especially near the regime of or molecular and atomic species\cite{ceperleyH}, or excitons in electron-hole plasmas. The system is better understood for fully ionized systems, such as hydrogen, in the form of free electrons and protons, and fully-ionized electron-hole condensates. In fact, considerable headway has been made using methods based on density-functional theory (DFT), even for plasmas with multiple states of ionization. DFT methods have been used with molecular-dynamics based approaches\cite{kwon,dejarlais,mazevet}, as well as within multi-component integral-equation approaches \cite{pdw95,ilciacco}. Equilibrium properties of plasmas, as well as their linear transport properties, have been successfully studied in these papers, and excellent agreement between the molecular-dynamics based DFT and integral-equation based DFT has been found \cite{lvm}. On the other hand, laser-produced plasmas are initially formed as two-temperature plasmas, where the electrons have absorbed the laser energy and have self-equilibrated to some ``electron temperature'' $T_e$, while the ions remain cool, at some temperature $T_i$, with $T_i<T_e$. The opposite situation arises in shock-wave generated plasmas, where the ions absorb the shock energy and $T_i>T_e$. Such two-temperature plasmas also occur in astrophysical settings, affecting the time of termination of synthesis of light-nuclei to occur at different stages of cooling of the electrons\cite{astrop}, and influencing the Coulomb nuclear-tunneling rates\cite{dewitt}. The possibility of such well-defined two-temperature plasmas is largely a result of the extreme mass ratio $m_i/m_e\ge 1836$ between ions and electrons. Similar, but less well defined situations can arise in electron-hole plasmas, where the masses are of the same order of magnitude (e.g, the electron and hole masses in GaAs are 0.067$m_e$ and 0.34$m_e$ respectively, with an electron-hole mass ratio of $\sim 5$). GaAs is a direct bandgap material, and electron-hole plasmas are more easily studied in indirect-gap systems like Si where the density-of states mass ratio is $\sim 3$. The simulation of such systems at two temperatures, using quantum Monte-Carlo methods is at present unavailable, even in regimes of densities and temperatures where bound states (or exciton formation in electron-hole systems) do not exist. Thus it is natural to look for analytical methods based on integral-equation techniques which are computationally simple and physically insightful. However, although $T_e,\, T_i$ define the temperatures of each subsystem and the pair-distribution functions (PDFs) $g_{ee}$ and $g_{ii}$, the `temperature' $T_{ei}$ entering into the cross-correlations $g_{ei}$, as well as the the effects of electron spin, exchange etc., relevant to two-temperature systems need to be clarified. In this context we use $T_{ee}=T_e$, $T_{ii}=T_i$, and $T_{ei}$ to refer to the electron-, ion-, and electron-ion temperatures as they enter independently into the Ornstein-Zernike (OZ) and hypernetted chain (HNC) equations. In fact, some authors\cite{seuf} have proposed to modify the well-established OZ equations in dealing with two-temperature (2T) two-mass (2M) systems. The objective of this paper is to study such 2T-2M plasmas using results from molecular-dynamics (MD)\cite{mc}, HNC\cite{hncref}, classical-map HNC (CHNC)\cite{prl1,prb3d}, quantum Monte Carlo (QMC)\cite{mc} and Kohn-Sham (KS)\cite{ilciacco} methods to establish the proper implementation of quantum effects and 2T, 2M situations in simulation studies. One of our main interests would be uniform hydrogenic plasmas free of bound states, in the regime of warm-dense matter. | The simplest classical rendering of quantum plasmas, based on the use of diffraction corrected potentials (Eq.~\ref{zerothset}) was used with HNC calculations and MD simulations to resolve the ambiguities and difficulties in handling the two-temperature, two-mass system. We conclude that the modifications to the OZ equations proposed by Seuferling et al.\cite{seuf}., are not needed. The classical mapping of quantum systems to the HNC equations, as used in the CHNC was confirmed by comparisons with Kohn-Sham DFT calculations as well as with available PIMC results for compressed hydrogen plasmas at finite temperatures. We conclude that the HNC and CHNC, together with the standard OZ equations provide excellent, accurate and simple analytical tools for the investigation of many-particle quasi-equilibrium systems for which direct quantum simulations continue to remain too prohibitive or unfeasible. | 7 | 10 | 0710.2888 |
0710 | 0710.3076_arXiv.txt | Ca II triplet spectroscopy has been used to derive stellar metallicities for individual stars in four LMC fields situated at galactocentric distances of 3\arcdeg, 5\arcdeg, 6\arcdeg\@ and 8\arcdeg\@ to the north of the Bar. Observed metallicity distributions show a well defined peak, with a tail toward low metallicities. The mean metallicity remains constant until 6\arcdeg\@ ([Fe/H]$\sim$-0.5 dex), while for the outermost field, at 8\arcdeg, the mean metallicity is substantially lower than in the rest of the disk ([Fe/H]$\sim$-0.8 dex). The combination of spectroscopy with deep CCD photometry has allowed us to break the RGB age--metallicity degeneracy and compute the ages for the objects observed spectroscopically. The obtained age--metallicity relationships for our four fields are statistically indistinguishable. We conclude that the lower mean metallicity in the outermost field is a consequence of it having a lower fraction of intermediate-age stars, which are more metal-rich than the older stars. The disk age--metallicity relationship is similar to that for clusters. However, the lack of objects with ages between 3 and 10 Gyr is not observed in the field population. Finally, we used data from the literature to derive consistently the age--metallicity relationship of the bar. Simple chemical evolution models have been used to reproduce the observed age--metallicity relationships with the purpose of investigating which mechanism has participated in the evolution of the disk and bar. We find that while the disk age--metallicity relationship is well reproduced by close-box models or models with a small degree of outflow, that of the bar is only reproduced by models with combination of infall and outflow. | Despite decades of work, there are still significant gaps in our knowledge of the LMC's star formation and chemical enrichment histories \citep{ol96}. This is motivated in part by the vastness of its stellar populations and to our limitations in observing sizable samples of stars. The age distribution of star clusters is relatively well known \citep[e.g.][]{geisler97}: there is an age interval between 10 and 3 Gyr with almost no clusters. This so-called age-gap may also correspond to an abundance gap \citep{ol91}, since the old clusters are metal-poor while the young ones are relatively metal-rich. The star formation history (SFH) of field stars is much less precisely known. Studies using {\itshape HST} data in small fields, suggest that the SFH of the LMC disk has been more or less continuous, with some increase in the star formation rate (SFR) in the last few Gyr \citep[e.g.][]{Smecker-Hane02,castro01,holtzman99}. This contrasts with previous results (based on much shallower data), which found a relatively young age (a few Gyr) for the dominant LMC population \citep{hardy84,bertelli92,westerlund95,vallenari96}. The SFH of the LMC can now be studied using sufficiently deep ground-based data (e.g.\ reaching the oldest main-sequence turn-off with good photometric precision) in large areas and in different positions of the galaxy \citep[e.g.][hereafter Paper I]{gshpz04,gshpz05}. Less is known about the LMC chemical enrichment history. The age--metallicity relation (AMR) normally used for the LMC is defined by star clusters \citep[e.g.][]{ol91,geisler97,bica98,dirsch00}. All works based on clusters obtain similar results. The mean metallicity jumps from [Fe/H] $\sim -1.5$ for the oldest clusters to [Fe/H] $\sim -0.5$ for the youngest ones. There were no studies on the field population AMR until the last decade. \citet{dopita97}, from a study of $\alpha$-elements in planetary nebulae, obtained a result qualitatively similar to that found in the clusters although only ten objects were used. Later, \citet{bica98}, using Washington photometry, and \citet{c00} and \citet{dirsch00}, using Str\"ongrem photometry, obtained the metallicity of RGB stars in different positions of the LMC. \citet{c00} also observed 20 stars with the infrared CaII triplet (CaT). All of them found that the age-gap observed in the clusters is not found in the field population. More recently, \citet{c05} have obtained stellar metallicities for almost 400 stars in the bar of the LMC, also using CaT lines. The metallicity for each star has been combined with its position in the color--magnitude diagram (CMD) to estimate its age. The obtained AMR shows a similar behavior to that of the clusters for the oldest population. However, while for the clusters the metallicity has increased over the last 2 Gyr, this has not happened in the bar. Finally \citet{gratton04} measured the metallicity of about 100 bar RR Lyrae, obtaining an average metallicity of [Fe/H] $=-1.48$. Another point that it is necessary to investigate is the presence, or not, of an abundance gradient in the LMC. From observations of clusters, \citet{ol91} and \citet{santos99} found no evidence for the presence of a radial metallicity gradient in the LMC. The only evidence for a radial metallicity gradient in the LMC cluster system was reported by \citet{kon93} based on six outer LMC clusters ($\geq$ 8 kpc). \citet{hill95} were the first to report evidences of a gradient in the field population from high-resolution spectroscopy, using a sample of nine stars in the bar and in the disk. They found that the bar is on average 0.3 dex more metal-rich than the disk population, but the disk stars studied are located within a radius of 2\arcdeg. Subsequent work by \citet{cioni03}, detected that the C/M ratio between number of asymptotic giant branch stars of spectral types C and M increased when moving away from the bar and within a radius of 6$\fdg$7. As the C/M ratio is anticorrelated with metallicity, this increment implies a decrease in metallicity. Finally, an outward radial gradient of decreasing metallicity was also found by \citet{alves04} from infrared CMD using the 2MASS survey. We have obtained deep photometry with the Mosaic II CCD Imager on the CTIO 4m telescope in four disk fields at different distances from the center of the LMC (RA=5$^h$23$^m$34$\fs$5, $\delta$=-69$\arcdeg$45'22``) with a quality similar to that obtained by {\itshape HST} in more crowded areas (see Paper I). The position of these fields, together with a description of their CMDs, are presented in Paper I, and detailed SFHs will be published in forthcoming papers. In the present investigation we focus on obtaining stellar metallicities for a significant number of individual RGB stars in these four fields using spectra obtained with the HYDRA spectrograph at the CTIO 4m telescope. In Section \ref{targetselection} we present our target selection. The observations and data reduction are presented in Section \ref{datareduction}. The radial velocities of the stars in our sample are obtained in Section \ref{radialvelocities}. In Section \ref{cat} we discuss the calculation of the CaT equivalent widths and the determination of metallicities. Section \ref{agedetermination} presents the method used to derive ages for each star by combining the information on their metallicity and position on the CMD. The analysis of the data is presented in Section \ref{analysis}, where the possible presence of a kinematically hot halo is discussed and the AMR and the SFH in each of our fields are described and compared. In the final section, the AMRs are compared to theoretical models and conclusions are drawn about the chemical evolution of the LMC. | \label{conclusions} Using infrared spectra in the CaT region, we have obtained metallicities and radial velocities for a sample of stars in four LMC fields. Metallicities have been calculated using the relationships between the equivalent width of the CaT lines, $\Sigma Ca$, and metallicity derived in Paper II. In addition, we have estimated the age of each star using a relationship derived from a synthetic CMD which, from the color, magnitude and metallicity of a star, allow us to estimate its age. The main results of this paper are: \begin{itemize} \item The velocity distribution observed in each field agrees with the rotational thick disk kinematics of the LMC. The velocity dispersion is slightly larger for the most metal-poor stars. However, the values obtained do not indicate the presence of a kinematically hot halo. \item The metallicity distribution of each field has a well defined peak with a tail toward low metallicities. The mean metallicity is constant until the field at 6\arcdeg\@ ([Fe/H]$\sim$-0.5 dex), and is a factor two more metal-poor for the outermost field ([Fe/H]$\sim$-0.8 dex). \item The AMR observed in each disk field is compatible with a single global relationship for the disk. We conclude that the outermost field is more metal-poor on average because it contains a lower fraction of relatively young stars (age$\leq$5 Gyr), which are also more metal-rich. \item The disk AMR shows a prompt initial chemical enrichment. Subsequently, the metallicity increased very slowly until about 3 Gyr ago, when the rate of metal enrichment increased again. This AMR is similar to that of the cluster system, except for the lack of clusters with ages between 3 and 10 Gyr. The recent fast enrichment observed in the disk and in the cluster system is not observed in the bar. \item The $\psi(t)$ of the three innermost fields show a first episode of star formation until about 10 Gyr ago, followed by a period with a low SFR until $\sim$5 Gyr ago, when the SFR increases, and reaches its highest values $\sim$2-3 Gyr ago. The outermost field does not show the recent increase of SFR. The second main episode is also observed, and is more prominent in the bar, where an increased SFR at old ages is not observed. The lower SFR between 5 and 10 Gyr ago is probably related to the age-gap observed in the clusters. \item Under the assumption of a solar yield, the disk AMR is well reproduced either by a chemical evolution model with outflow with $\lambda$ between 1 and 2, so the disk losses the same amount of gas that has taken part in the star formation, or by models combining infall and outflow with $\alpha$=0.2 and $\lambda$=0.05, which means that the galaxy is almost a closed-box system. With a smaller yield, the AMR could also be reproduced with a closed-box model. The bar AMR is well reproduced by models with a combination of inflow and outflow with $\alpha$=1.2 and $\lambda$ between 0.4 and 0.5. This suggests that the amount of infalling gas was larger than the amount that participated in the star formation in the bar, and also that the amount of gas that escaped the bar was 50\% of the total that participated in star formation. \end{itemize} | 7 | 10 | 0710.3076 |
0710 | 0710.1073_arXiv.txt | We determine the ratio of helium- to hydrogen-atmosphere white dwarf stars as a function of effective temperature from a model atmosphere analysis of the infrared photometric data from the Two Micron All Sky Survey combined with available visual magnitudes. Our study surpasses any previous analysis of this kind both in terms of the accuracy of the $\Te$ determinations as well as the size of the sample. We observe that the ratio of helium- to hydrogen-atmosphere white dwarfs increases gradually from a constant value of $\sim0.25$ between $\Te=15,000$~K and 10,000~K to a value twice as large in the range $10,000 > \Te > 8000$ K, suggesting that convective mixing, which occurs when the bottom of the hydrogen convection zone reaches the underlying convective helium envelope, is responsible for this gradual transition. The comparison of our results with an approximate model used to describe the outcome of this convective mixing process implies hydrogen mass layers in the range $M_{\rm H}/M_{\rm tot}=10^{-10}$ to $10^{-8}$ for about 15\% of the DA stars that survived the DA to DB transition near $\Te\sim 30,000$~K, the remainder having presumably more massive layers above $M_{\rm H}/M_{\rm tot}\sim10^{-6}$. | The Two Micron All Sky Survey (2MASS) represents one the largest set of homogeneous infrared photometric data for all areas of the sky applicable to white dwarf stars of various spectral types. Indeed, the 2MASS survey contains $JHK_S$ magnitudes for almost all point sources of the sky up to the magnitude thresholds of $J \sim 16.8$, $H \sim 16.5$, and $K_S\sim 16.2$ for detections with any formal uncertainty \citep[see for instance Fig.~1 of][]{tremblay07}. This yields a nearly complete magnitude limited survey of white dwarfs that allows us to trace a portrait of the local population of white dwarfs, assuming we can properly identify the white dwarfs in this survey. To that effect, the most complete compilation is that of the Villanova White Dwarf Catalog\footnote{http://www.astronomy.villanova.edu/WDCatalog/index.html} (WDC), which contains more than 5500 spectroscopically identified white dwarfs reported in the literature. The interest of such a sizeable sample of homogeneous infrared photometric observations is to reach an unprecedented level in the statistical analysis of white dwarf stars. In addition, \citet{tremblay07} have demonstrated the overall reliability of the 2MASS Point-Source Catalog (PSC) photometry for the purpose of detailed comparisons with the predictions from white dwarf model atmospheres. Studies of the spectral evolution of white dwarfs attempt to understand how the surface composition and the corresponding spectral signature evolve along the white dwarf cooling sequence (see \citealt{fontaine87,fontaine91} for a review). Because of the intense gravitational field present in these stars, hydrogen will float to the surface in the absence of competing mechanisms, while heavier elements will sink below the photospheric regions. Indeed, the majority of white dwarfs have optical spectra that are completely dominated by strong hydrogen Balmer lines --- the DA stars --- with hydrogen-dominated atmospheres. We also find that almost 25\% of white dwarfs have helium-dominated atmospheres, mostly DO, DB, DQ, and DZ stars, and some of the DC stars as well. The fact that not all white dwarfs have hydrogen-rich atmospheres imply that either some stars are born with hydrogen-deficient atmospheres and remain that way throughout their evolution, or alternatively, that several physical mechanisms such as diffusion, accretion, convection, radiation pressure, and stellar winds are competing with gravitational settling in determining their surface composition as they cool off. This particular issue can be resolved by analyzing the distribution of spectral types as a function of effective temperature (or any temperature index) for a large sample of white dwarf stars, in a way similar to the analysis of \citet{sion84} nearly twenty five years ago. Before discussing the physical processes that will affect the spectral evolution of hydrogen- and helium-atmosphere white dwarfs, it is important first to consider the constraints brought about by the studies of the hottest white dwarfs. The canonical mass fractions of light elements in white dwarfs considering the mass threshold for residual nuclear burning are $M_{\rm He}/M_{\rm tot} \sim 10^{-2}$ and $M_{\rm H}/M_{\rm tot} \sim 10^{-4}$. However, multiple excursions on the AGB nuclear burning phase, during which very late helium flashes could remove essentially all of the hydrogen \citep{werner06}, could also produce white dwarfs with total hydrogen masses much lower than the canonical value of $10^{-4}$. The first class would correspond to the progenitors of the hottest DA stars while the second class would represent the progenitors of hot white dwarf stars with hydrogen-deficient atmospheres --- the PG 1159 stars at very high effective temperatures, and the DO stars at slightly lower temperatures, whose spectra are dominated by He\textsc{ii} lines. It has traditionally been believed that by the time DO stars cool off to $\Te\sim45,000$~K, residual amounts of hydrogen ($M_{\rm H}/M_{\rm tot} \sim 10^{-16}$) thoroughly mixed in their helium-rich envelope would gradually accumulate to the surface, turning all helium-atmosphere white dwarfs into DA stars \citep{fontaine87,fontaine91}. This scenario would account for the existence of the so-called DB gap, a range of temperature between $\Te=45,000$~K and 30,000~K where all white dwarfs have a DA spectral type. However, the recent discovery in the Sloan Digital Sky Survey of several DB stars within the gap \citep{eisenstein06a} implies that the total mass of hydrogen left in the envelope of DO stars can be even smaller than previously believed. It is then generally admitted that the significant increase in the number of helium-atmosphere DB white dwarfs below $\Te\sim30,000$~K can be explained in terms of the convective dilution of the superficial hydrogen atmosphere by the underlying helium convective envelope, provided that the hydrogen layer is sufficiently thin \citep[$M_{\rm H}/M_{\rm tot} \sim 10^{-15}$,][]{macdonald91}. At even lower effective temperatures ($\Te\lesssim12,000$~K), hydrogen-atmosphere white dwarfs get a second opportunity to evolve into stars with helium-dominated atmospheres when the superficial hydrogen layer becomes convective over a significant fraction of its depth. This process is illustrated in Figure \ref{fg:f1} where the extent of the hydrogen convection zone is displayed as a function of decreasing effective temperature for a 0.6 \msun\ DA white dwarf, based on evolutionary models with thick hydrogen layers similar to those described by \citet{fon01} and kindly provided to us by G.~Fontaine and P.~Brassard. These calculations show that if the hydrogen envelope is thin enough, the bottom of the hydrogen convection zone may eventually reach the underlying and more massive convective helium layer, resulting in a mixing of the hydrogen and helium layers \citep{strittmatter71,shipman72,baglin73,koester76,vauclair77}. Figure \ref{fg:f1} also indicates that the effective temperature at which this mixing occurs will depend on the thickness of the hydrogen envelope. The thicker the envelope, the lower the mixing temperature; if the hydrogen layer is more massive than $M_{\rm H}/M_{\rm tot} \sim 10^{-6}$, mixing will never occur. The simplest physical model that can be used to describe this convective mixing process is to assume that hydrogen and helium are homogeneously mixed. Since the helium convection zone is much more massive ($M_{\rm He-conv}/M_{\rm tot}\sim10^{-6}$) than the hydrogen layer when mixing occurs, it is generally assumed that a DA star would be transformed into a helium-atmosphere white dwarf of type DB, DQ, DZ, or DC with only a trace abundance of hydrogen. More detailed calculations discussed by \cite{fontaine91} confirm these predictions. If such a process takes place in cool white dwarfs, we then expect the ratio of helium- to hydrogen-atmosphere white dwarfs to increase at low effective temperatures. A comparison between the observed ratio as a function of $\Te$ and the theoretical expectations would make it possible to estimate the thickness of the hydrogen layers in DA white dwarfs. In turn, these determinations could be compared to independent measurements of the hydrogen layer mass inferred from ZZ Ceti asteroseismology. The best approach to study the occurrence of this convective mixing process begins with the statistical analysis of large samples of cool white dwarfs for which we can determine the main atmospheric constituent and the effective temperature, or any analogue temperature index such as the absolute $V$ magnitude. \citet{sion84} was the first to study this problem by estimating the non-DA to DA ratio at low effective temperatures using a proper motion sample of 695 white dwarfs drawn from the catalog of spectroscopically identified white dwarfs available at that time. The results from Figure 1 of \cite{sion84} are reproduced here in Figure \ref{fg:f2} in terms of the non-DA to DA ratio (rather than the absolute number of DA and non-DA stars) as a function of the absolute visual magnitude $M_V$. Looking at the results, we notice a first increase in this ratio at $M_V>11.25$, which corresponds to $\Te\sim15,000$ K for a 0.6 \msun\ white dwarf, a temperature much in excess of the theoretically predicted value for the convective mixing process. A second increase occurs at $M_V\sim 12.5$, or $\Te \sim 10000$ K, but the observed ratio drops suddenly afterwards and this increase is probably not significant. Then a third increase occurs at $M_V> 13$, or $\Te<8000$~K. Clearly, the evidence based on these results are not exactly convincing, and our current picture of the situation below 12,000 K is at best sketchy. The global portrait did not change significantly since the analysis of \citet{sion84}, and this remains the only available proof quoted in the literature of the transformation of some DA stars into non-DA stars (see also \citealt{greenstein86}). Our goal is to improve upon the analysis of \citet{sion84} by using the 2MASS photometric sample combined with the WDC database to identify white dwarfs and determine their atmospheric composition, and most importantly to obtain more accurate temperature determinations using a full model atmosphere analysis of the $VJHK_S$ photometry. We first describe in \S~2 the 2MASS photometric white dwarf sample used in our analysis. The determination of the atmospheric parameters is then presented in \S~3 and the uncertainties related to our approach are discussed at length in \S~4. The evolution of the ratio of helium- to hydrogen-atmosphere white dwarfs as a function of effective temperature is then determined in \S~5. Our conclusions follow in \S~6. | We have presented a detailed study of the spectral evolution of cool white dwarfs following the pioneering effort of \citet{sion84}. In particular, we have determined the ratio of helium- to hydrogen-atmosphere white dwarfs below $\Te=15,000$~K as a function of effective temperature by increasing the size of the sample by a factor of 2, and by obtaining more precise measurements of the effective temperature of these stars through model atmosphere fits to Johnson $V$ (or Str\"omgren $y$) and 2MASS $JHK_S$ photometric observations. We have confirmed that in this temperature range, the evolution of hydrogen- and helium-atmosphere white dwarfs is not independent, and we have revealed in greater detail the role of the convective mixing process responsible for the coupling between these two white dwarf types. We have found that about 15$\%$ of cool hydrogen-atmosphere DA white dwarfs between $\Te=15,000$~K and 10,000~K are transformed into helium-atmosphere non-DA white dwarfs at lower temperatures. Using a basic model for convective mixing, our results imply in turn that 15\% of the DA stars that have survived the DA to DB transition have hydrogen layer masses in the range $M_{\rm H}/M_{\rm tot}\sim10^{-10}$ to $10^{-8}$, although the exact values depend on the assumed convective efficiency in the evolutionary model calculations. The remaining DA stars would presumably have more massive hydrogen layers, probably in excess of $M_{\rm H}/M_{\rm tot}\sim10^{-6}$. Further analysis of the convective mixing problem would require a significantly larger photometric sample, such as the {\it ugriz} photometric sample of the Sloan Digital Sky Survey (SDSS). This sample covers a few areas of the sky with a completeness of around 95\% up to $g=20$. However, most spectroscopic follow-ups have relied on color-color diagrams to identify white dwarfs \citep{kl04,eisenstein06b}, and this method is strongly biased towards hot white dwarfs ($\Te > 12000$~K). Therefore, the SDSS white dwarfs identified so far cannot be used to study the convective mixing process at low effective temperatures. One way to proceed to identify cooler white dwarfs in the SDSS is through proper motion-color diagrams. \citet{harris06} have used this method to identify nearly 6000 white dwarfs but a complete spectroscopic analysis similar to that of \citet{kilic06} has yet to be performed to take full advantage of this sample. | 7 | 10 | 0710.1073 |
0710 | 0710.1135_arXiv.txt | The Cygnus Loop was observed from the northeast to the southwest with XMM-Newton. We divided the observed region into two parts, the north path and the south path, and studied the X-ray spectra along two paths. The spectra can be well fitted either by a one-component non-equilibrium ionization (NEI) model or by a two-component NEI model. The rim regions can be well fitted by a one-component model with relatively low \kTe~whose metal abundances are sub-solar (0.1--0.2). The major part of the paths requires a two-component model. Due to projection effects, we concluded that the low \kTe~($\sim$0.2\,keV) component surrounds the high \kTe~($\sim$0.6\,keV) component, with the latter having relatively high metal abundances ($\sim$5 times solar). Since the Cygnus Loop is thought to originate in a cavity explosion, the low \kTe~component originates from the cavity wall while the high \kTe~component originates from the ejecta. The flux of the cavity wall component shows a large variation along our path. We found it to be very thin in the south-west region, suggesting a blowout along our line of sight. The metal distribution inside the ejecta shows non-uniformity, depending on the element. O, Ne and Mg are relatively more abundant in the outer region while Si, S and Fe are concentrated in the inner region, with all metals showing strong asymmetry. This observational evidence implies an asymmetric explosion of the progenitor star. The abundance of the ejecta also indicates the progenitor star to be about 15\,\Msun. | A supernova remnant (SNR) reflects the abundance of the progenitor star when the remnant is young and that of the interstellar matter (ISM) when it becomes old. In this way, we can study the evolution of the ejecta and the ISM. The Cygnus Loop is a proto-typical middle-aged shell-like SNR. The angular diameter is about 2$^{\circ}$.4 and it is very close to us (540\,pc; Blair et al.\ 2005), implying a diameter of $\sim$23\,pc. The estimated age is about 10000\,years, less than half that based on the previous distance estimate of 770\,pc \cite{minkowski58}. Since the Cygnus Loop is an evolved SNR, the bright shell mainly consists of a shock-heated surrounding material. Its supernova (SN) explosion is generally considered to have occurred in a preexisting cavity \cite{mccray79}. Levenson et al. (1997) found that the Cygnus Loop was a result of a cavity explosion that was created by a star no later than B0. It is almost circular in shape with a break-out in the south where the hot plasma extends out of the circular shape. Miyata et al. (1994) observed the northeast (NE) shell of the Loop with ASCA and revealed the metal deficiency there \cite{miyata94}. Since Dopita et al. (1977) reported the metal deficiency of the ISM around the Cygnus Loop, they concluded that the plasma in the NE-shell is dominated by the ISM. Due to the constraints of the detector efficiency, they assumed that the relative abundances of C, N and O are equal to those of the solar value \cite{anders89}. More recently, Miyata et al. (2007) used the Suzaku satellite \cite{mitsuda07} to observe one pointing position in the NE rim. They detected emission lines from C and N and determined the relative abundances \cite{miyata07}. They concluded that the relative abundances of C, N and O are consistent with those of the solar values whereas the absolute abundances show depletion from the solar values \cite{anders89}. Katsuda et al. (2007) observed four pointings in the NE rim and detected a region where the relative abundances of C and N are a few times higher than that of O. Hatsukade \& Tsunemi (1990) detected a hot plasma inside the Cygnus Loop that is not expected in the simple Sedov model \cite{hatsukade90}. They reported that the hot plasma was confined inside the Loop. Miyata et al. (1998) detected strong emission lines from Si, S and Fe-L from inside the Loop \cite{miyata98}. They found that the metal abundance is at least several times higher than that of the solar value \cite{anders89}, indicating that a few tens of higher than that of the shell region. They concluded that the metal rich plasma was a fossil of the SN explosion. The abundance ratio of Si, S and Fe indicated the progenitor star mass to be 25\Msun. Miyata \& Tsunemi (1999) measured the radial profile inside the Loop and found a discontinuity around 0.9\,R$_\mathrm{s}$ where R$_\mathrm{s}$ is the shock radius. They measured the metallicity inside the hot cavity and estimated the progenitor mass to be 15\Msun. Levenson et al. (1998) estimated the size of the cavity and the progenitor mass to be 15\Msun. Therefore, the progenitor star of the Cygnus Loop is a massive star in which the triple-$\alpha$ reaction should have dominated rather than the CNO cycle. If the surrounding material of the Cygnus Loop is contaminated by the stellar activity of the progenitor star, it may explain the C abundance inferred for this region with Suzaku \cite{katsuda07}. In order to study the plasma condition inside the Cygnus Loop, we observed it from the NE rim to the south-west (SW) rim with the XMM-Newton satellite. We report here the result covering a full diameter by seven pointings. | We have observed the Cygnus Loop along the diameter from the NE rim to the SW rim employing XMM Newton. The FOV is divided into two paths: the north path and the south path. Then it is divided into many small annuli so that each annulus contains a similar number of photons to preserve statistics. The spectra from the rim regions can be expressed by a one-\kTe~component model while those in the inner region require a two-\kTe~component model. The low \kTe~plasma shows relatively low metal abundance and covers the entire FOV. It forms a shell that originates from the preexisting cavity. The high \kTe~plasma shows high metal abundance and occupies a large part of the FOV. The origins of these two components are different: the high \kTe~plasma with the high metal abundance must come from the ejecta while low \kTe~plasma with low metal abundance must come from the cavity material. We find that the thickness of the shell is very thin in the south west part where, we guess, the ejecta plasma is blow out in the direction of our line of sight. We estimate the mass of the metals. Based on the relative metal abundance, we find that the Cygnus Loop originated from a 15\,\Msun~star. The distribution of the ejecta is asymmetric, suggesting an asymmetric explosion. | 7 | 10 | 0710.1135 |
0710 | 0710.3130_arXiv.txt | We have considered precession in accretion disks in which a second moment of inertia relative to an axis perpendicular to the axis of rotation may be very important. This formalism, that takes into account the precession contribution to the angular momentum, is based on the existence of a parameter $\it p$ which determines three characteristic densities resulting from the averaging process and imposes constraints on the actual disk density. It is shown that the precession velocity will lie in a three branch solution, and depends on how large is the disk actual density as compared to the characteristic densities. Besides the large spread on the solution for the precession velocity, depending on the density strength, it may be prograde and retrograde. It is shown that the keplerian thin disk, with very large density values compared to characteristic ones , only precesses very far away from the primary object, which implies very large precession periods. For other models, the disk will thicken, with large deviations from the keplerian approximation. Constraints on the density only will be effective for very large values of the ratio ${{\dot M} \over M_{p}}$, respectively, accretion rate and mass of the primary. Under this condition, the structure of the precessing region is found. Lower bounds on the precession period are found for not so large values of this ratio. Deviations from the mean precessional motion are considered. It is shown that these deviations result in periodic motions as long as the time scales associated to them are comparable to the remaining time scales. Otherwise, they result in misalignment motions, forcing the plane of the disk to become normal to the orbital plane of the secondary. | Precessional activity has long been invoked to explain time variations occurring in the spectra of galactic X-ray binaries like LMC X-4, Her X-1, SS 443, Cygnus X-1, etc. Observational evidences give support to the idea that these systems, and quite a lot of suspected others, consist of a disk-like system plus a third body. This is very suggestive of precession as being the effect due to the perturbing torque of a distant companion star in the the disk or, even, as the result in the disk of the secondary star precessing in the tidal field of the central compact object. As a matter of fact, the scenario is not unique, and according to \citet{prie87}, there are, at least , five possibilities. This is so because there is uncertainty concerning the reason why the disk precesses, where and the extent on it that precesses. The possibilities explicitly dependent upon tidal torques under which the disk will precess can be summarized by the following two scenarios: a- the disk has a permanently tilted edge that freely precesses in the tidal field of the companion star. This kind of scenario has been first exploited by \citet{ka73}, trying to explain the 35 day period on the light curve of Her X-1. This same model was used by \citet{ka80} to explain observational features in SS 443; b- the companion star precesses in the tidal field of the central compact star, slavishly followed by the the outer tilted disk. This kind of scheme was proposed by \citet{rob74}, also to explain the 35 day period observed in Her X-1. However, precession is not a exclusiveness of galactic X -ray binaries sources, its presence being suspected in a lot of other astrophysical systems. Jet structures observed in AGN, and also in protostars seem to be associated with a precessing disk ( \citet{li99}). In the context of disks around supermassive objects, the reason why the disk precesses is still more unclear. \citet{ka97} suggested that observational data in OJ 287 could be explained by a tidal torque due to the presence of another massive companion star, in a way similar to that employed to Her X-1 and SS 443. \citet{prin96} and \citet{prin97} suggests precession as a result of a warping instability caused by the central source irradiation. \citet{rom00} argue, in analogy with \citet{ka82}and \citet{ka97},, aiming to explain the pattern of ejection observed in the Quasar 3C 273, that precession is due to a massive secondary black hole tidal torque on the disk,a model, subsequently, applied to 3C 345 by \citet{ca02}. It should be argued, however, that precession is a kind of solid-like response of the disk. A fluid-like response would be a more appropriate one for a fluid system like the disk. Besides, in the case of fluid-like response, the tidal torque acts as a perturbation, which main consequence is the generation of different sort of waves. Quite a variety of physical issues, suitable for a wave approach, such as several tidally driven instabilities, disk tidal deformation (warping), horizontally and vertically driven resonances, angular momentum transport, and so on (\citep{lu91, lua92, lu93, olga02, olgb03, vis92}) have been tackled quite successfully. In addition, it has been shown that, under this approach, for disks in systems with extreme mass ratio, precession may occur due to a coupling between an eccentric instability and Lindblad resonances (\citet{lua92}). \citet{pat95}, extending \citet{lua92} result, argue that the axisymmetric part of the tidal potential is related to the disk solid-like response, while the disk fluid-like response is related to the non-symmetric part of the potential. In these works, the the disk solid body-like response comes as an integrabil;ity condition. \citet{ka82} proposal to compute frequencies treats the disk as a ring, with no allowance given for the width in the plane of the ring, nor to the thickness perpendicular to that plane. By employing a suitable averaging procedure, he obtains for the for the precession angular velocity $${\Omega}_{p}=-\frac{3}{4} \, \frac{{{\omega}_{s}}^2}{{\Omega}_{d}} \, cos{\delta},$$ where ${\omega}_{s}$ is the keplerian angular velocity of the secondary, ${\Omega}_{d}$ is the ring angular velocity, and $\delta$ is the inclination angle. This formalism, besides treating the disk as a rigid body, not taking into account its structure, assumes the disk angular velocity parallel to the disk angular momentum . Though keeping $\vec{{\Omega}_{p}}$ parallel to ${\vec L}$, assuming thin disks obeying polytropic equation of state, the structure has been considered by \citet{pat95}, \citet{laral96}, \citet{larp97} and \citet{lar97}. Decomposing the secondary tidal potential into odd and even z parts, with subsequent Fourier decomposition, and arguing that the odd axisymmetric part is responsible for rigid precession, these authors,assuming ${\ell} \over r$ constant, obtain $${\Omega}_{p}=-\frac{3}{4} \, \, {\left( \frac{7-2 \, n}{5-n} \right)} \, \frac{{{\omega}_{s}}^2}{{\Omega}_{d}} \, cos{\delta},$$ as an integrability condition. In this expression, n is the polytropic index and ${\Omega}_{d}$ is the disk angular velocity evaluated at the outer edge of the disk, $R_{0}$. For $n=2$, they recover \citet{ka82} result. According to them, the validity of this result holds whenever the disk doesn't thicken, the perturbative tidal potential is weak, and, above all, the sound crossing time is shorter than the precession period. It should be argued, however, that the disk angular momentum and the precession velocity are not parallel. Besides, heat generation, cooling and angular momentum transport do occur in accretion disks, a fact that makes not realistic the structure obtained from a polytropic equation of state. This imposes severe restrictions on the model of disk we want to consider.It is reasonable to expect the polytropic index varying from $n=\frac{3}{2}$, for a (monoatomic) gas pressure dominated disk, to $n=3$ for a radiation pressure dominated one. Disks with negligible temperature gradients would require a polytropic index much greater than 3: $\gamma =1$ is appropriate for an isothermal atmosphere ( $n \rightarrow \infty $). In that situation, the mass content in the disk, the moment of inertia and the precession velocity are determined by values close to the inner radius rather than close to the outer radius, the disk structure being very susceptible to relativistic corrections and gravitational radiation losses may become important. Besides, if one is interested in a region in parameter space $M-{\dot M}$ where the disk thickens, the dynamics can not be dissociated from the structure, as is the case for the keplerian thin disk. In other words, the azimuthal velocity depends on the disk height scale . In addition, for a given tidal potential, its relative strength increases as the disk thickens. For actual radiation pressure dominated, the height scale is constant, the polytropic index is negative, implying negative specific heat. Besides this, in this case, one can not satisfy null boundary condition at the surface of the disk Finally, the precession period diminishes relative to the sound crossing time as the disk gets thick. A final possible question concerning the applicability of of these results to disks is related to the precession velocity dependence on the outer radius of the disk, which yields proportionality with the mass of the disk. Despite implying larger inertia moments, the larger the mass the faster the disk will precess. \citet{rom00} work consists essentially of the application of \citet{ka82} and \citet{ka97} to a binary system composed by a disk around a primary supermassive black hole, plus a supermassive secondary orbiting the primary in a keplerian way, not coplanar to the disk. It should be argued, however, that the formulation followed by these authors, besides suffering from the drawbacks, we just mentioned, of any precessing disk model, relies heavily on the above cited authors results for the precession angular velocity, which was obtained without considering the precession contribution to the total angular momentum. Besides, this model is supposed to apply to disks eventually dominated by radiation pressure. In addition, account of the precession contribution to the angular momentum introduces some non-linearity in the problem, which, in turn, introduces conditions on the averaging procedure, imposing constraints on the density, on the disk unperturbed angular velocity, on the viscosity parameter as well as on the height scale of the disk. In that case, there will be a contribution from the precession velocity to a motion in a plane perpendicular to the plane containing the normal to the disk and the normal to the plane of the orbit of the secondary. This motion, depending on the magnitude of its time scale, will be periodic or not . In the following we shall address the question of precession in accretion disks in systems similar to that considered by \citet{rom00}. In our formulation, however, we will be concerned to treat the secondary tidal torque on the disk as a perturbation, which means that the secondary never crosses the disk, being far away from any point in the disk. Granted this, we will take into account the precession contribution to the total angular momentum and will discuss the conditions under which we may take an average of the Euler equations, in such a way as to keep algebraic solutions to the precessional motion. Deviations from the precessional motion will be treated by solving the equation for the time evolution of $\delta$, the angle that measures the misalignment between the normal to the plane of the disk and the normal to the plane of the secondary orbit. The time evolution will be obtained assuming the time scale for this motion is large or comparable to the remaining time scales. We obtain results for the precession velocity for different disk models, taking into account deviations from geometrical thinness due to a relation between the angular velocity of the disk and the height scale. One of the main results of this work is to show that, if precession is to occur in the inner region of disks around a supermassive black hole, the disk is not keplerian, being very thick and luminosity deficient. | We have treated the problem of precession in accretion disks taking into account its contribution to the total angular momentum. We have looked for conditions under which the problem may be treated by solving the algebraic Euler equations, i.e., constant angular velocities. We have found that the problem is characterized by the parameter p, given by $$p= {\frac{256}{27}} {{cos^2{\delta} \, cot^2{\delta}} \over {\beta}^4} \, {\left( {M_{p} + M_{s}}\over {M_{p}} \right)}^2 \, {\left( {r \over d} \right)}^4 ,$$ and the densities $${\rho}_{min}= 12 cos{\delta} cot{\delta} {\left({{M_{p} + M_{s}}\over {M_{p}}}\right)} {\frac{K}{{\beta}^2 \, {d}^3}}, $$ $${\rho}_{c}= {K \over r^3} ,$$ $${\rho}_{*}= 512 \, cos^3{\delta} cot^3{\delta} {\left( {M_{p} + M_{s}}\over {M_{p}} \right)}^3 {\left( {r \over {{\beta} d}} \right)}^6 {K \over r^3} .$$ For $p \leq 1$, the only constraint is $\rho \geq {\rho}_{min} $, and the solutions, as a function of the density, lie in three branches: the upper one, with prograde precession velocities, at the upper left of $\rho ={2 \over 3} {\rho}_{c} $; the lower branch, with retrograde velocities, at the lower right of $\rho ={2 \over 3} {\rho}_{c} $; and the middle branch, between the upper and the lower, in which the solution is retrograde for $\rho < {r \over d} {\rho}_{c} $, and prograde for $\rho > {r \over d} {\rho}_{c} $. For $ p \geq 1$, part of the $\rho $-space, $ {\rho}_{c} \leq {\rho} \leq {\rho}_{*} $, is not allowed. Using simple energy and angular momentum transport arguments, together with constraints obtained in that formulation, it is shown that, for $p >1$, $ \alpha \rightarrow 1$ and ${\ell} \rightarrow r$, leading to the breakdown of the keplerian and thin disk approximations. In that situation the disk is luminosity defficient. As a matter of fact, a very important contribution of this paper is to show the incompatibility between keplerian disk and thick disk. The procedure we have adopted is suitable for treating deviations from the mean precessional motion, the misalignment motion, as well. For different limiting expressions for the precession velocity, we have obtained the time evolution for $ \delta$, the angle that measures the misalignement of the plane of the disk and the orbital plane of the secondary star, the misalignment angle. It has been shown that the time evolution of ${\delta}$ depends on the magnitude of the time scale associated to this motion. Periodic motions only occur if this time scale is comparable to the remaining ones. Otherwise, there will be a tendency to force the plane of the disk to become normal to the secondary orbital plane. It is shown that, if the disk is keplerian, full misalignment, i.e. ${\delta}= 0.5 \, \pi$, is reached in a time comparable to the precession period, which means that the solution for the system is no longer given algebraically. We should make resort to the solutions of the differential Euler equations, a formidable task far beyond our goal in this paper. Finally, for different disk models we have shown that the constraints on the density will hold as long as ${{\dot M}_{1} \over M_{9}} >> 1$. In that situation, we have found the properties of the precessing region as a function of ${\dot M}_{1}$ and $ M_{9}$. An expression for the separation distance between the primary and the secondary is also found. Otherwise, if ${{\dot M}_{1} \over M_{9}} < 1$, lower bounds on the precession period and on the separation distance are found. These bounds depend on ${\dot M}_{1} \over M_{9}$. Application of this formalism to different disk models, may be summarized as follows: 1- for the keplerian thin $\alpha$ standard disk model, the precession period (in years) will be $$ T_{p} \geq 0.08 \, z_{d}^{1.5} \, M_{p}, $$ For accretion disks in binary systems, we may take the size of the disk as the truncation radius, the point where the Roche equipotentials first intersect. In our systems we are not allowed to do so but, in any case, we found in the literature that values for $ z \approx 100-1000 $ are not unusual. Besides, expecting precessiuon periods of about hundred years, the system will reach full misalignment in a time of the order of the precession period. The disk is not affected by constraints on the density. 2- For disk in which radiation dominates pressure and cooling, if ${{\dot M}_{1} \over M_{9}} > 75.76$ for a Schwarzschild black hole and ${{\dot M}_{1} \over M_{9}} > 12.62$ for a Kerr black hole, density in the disk may come close to the critical density ${\rho}_{*}$. In that situation, ${\ell} \approx r$, and the size of the disk can not exceed $$z_{d}= 0.0132 \, {{\dot M}_{1} \over M_{9}}.$$ Besides, if the system is not to be affected by gravitational radiation, ${{\dot M}_{1} \over M_{9}}$ should be much larger than the values given above. Precession period will be $$T_{p}= 2.81\times 10^{-3} \, cos{\delta} \, {\left( \frac{{{\dot M}_{1}}^3}{{M_{9}}^5} \right)}^{1/2} .$$ A reasonable value for $z_{d}$ would imply a very large ratio ${{\dot M}_{1} \over M_{9}}$. We are, therefore, led to the suspicion that the assumption $ {\ell} \approx r$, or $p_{0} \approx 1$, is very strong for this model. Abandon of this assumption leads to a lower bound on the precession period, given by $${T_{p}}^2=0.02 \, cos^2{\delta} \, \frac{{z_{d}}^3}{{M_{9}}^2} \, \frac{a^2}{{\left(4 \, a^2-{\left( {\left(1+4 \, a^2 \right)}^{1/2}-1\right)}^2 \right)}} .$$ 3-Finally, for a radiation pressure and advective cooling dominated disk, $ {\ell} \approx r$, and its size $$z_{d}= 0.75 \, {{\dot M}_{1} \over M_{9}}.$$ Again, if the system is not to be affected by gravitational radiation, otherwise it will live for a short period of time, ${{\dot M}_{1} \over M_{9}} >> 1$. If ${{\dot M}_{1} \over M_{9}}$ is not too large, ${\ell} << r$, and the constraints on the density will not be effective. Again, there will be a lower bound on the precession period given by $${T_{p}}^2=0.031 cos^2{\delta} \, \frac{z^3}{{M_{9}}^2} \, {{\left( {\left(1-4.4 \, a+8.84 \, a^2 \right)}^{1/2}+1-2.2 \, a \right)}^2\over {\left( {\left( {\left(1-4.4 \, a+8.84 \, a^2 \right)}^{1/2}+1-2.2 \, a \right)}^2-4 \, a^2 \right)}}.$$ To finalize, we would stress the main contribution of this work as pointing out the possibility of considering the disk structure, when studying disk precession, by means of an alternative procedure. | 7 | 10 | 0710.3130 |
0710 | 0710.4526_arXiv.txt | Halo coronal mass ejections (HCMEs) originating from regions close to the center of the Sun are likely to be geoeffective. Assuming that the shape of HCMEs is a cone and they propagate with constant angular widths and velocities, at least in their early phase, we have developed a technique (Michalek et al. 2003) which allowed us to obtain the space speed, width and source location. We apply this technique to obtain the parameters of all full HCMEs observed by the Solar and Heliospheric Observatory (SOHO) mission's Large Angle and Spectrometric Coronagraph (LASCO) experiment until the end of 2002. Using this data we examine which parameters determine the geoeffectiveness of HCMEs. We show that in the considered period of time only fast halo CMEs (with the space velocities higher than $\sim 1000{km\over s}$ and originating from the western hemisphere close to the solar center could cause the severe geomagnetic storms. We illustrate how the HCME parameters can be used for space weather forecast. It is also demonstrated that the strength of a geomagnetic storm does not depend on the determined width of HCMEs. This means that HCMEs do not have to be very large to cause major geomagnetic storms. | Coronal mass ejections (CMEs) originating from regions close to the central meridian of the Sun and directed toward Earth cause the most severe geomagnetic storms (Gosling, 1993; Kahler, 1992; Webb et al., 2001). Many of these Earth-directed CMEs appear as an enhancement surrounding the occulting disk of coronagraphs. We call them halo CMEs (Howard et al. 1982). The measured properties of CMEs include their occurrence rate, direction of propagation in the plane of the sky, angular width, and speed (e.g. Kahler, 1992; Webb, 2000; St. Cyr et al., 2000, Gopalswamy et al., 2003a; Gopalswamy, 2004; Yashiro et al., 2004). It is well known that the geoffective CMEs originate mostly within a latitude $\pm30^o$ (Gopalswamy et al., 2000a, 2001; Webb et al., 2000, 2001; Wang et al., 2002; Zhang et al. 2003). Srivastava and Venkatakrishan (2002) showed that the initial speed of the CMEs is correlated with the $D_{ST}$ index strength of the geomagnetic storm, although their conclusion was based only on the study of four events. This tendency was also suggested earlier by Gosling et al. (1990) and Tsurutani $\&$ Gonzalez (1998). On the other hand, Zhang et al. (2003) demonstrated that both slow and fast HCMEs can cause major geomagnetic disturbances. They showed that geoeffective CMEs are more likely to originate from the western hemisphere than from the eastern hemisphere. They also demonstrated a lack of correlation between the size of X-ray associated with a given CME and the importance of geomagnetic storms. Unfortunately, these studies were based on the sky plane speeds of CMEs without consideration of the projection effects. The parameters describing properties of CMEs, especially for HCMEs, are affected by projection effects (Gopalswamy et al., 2000b). Assuming that the shape of HCMEs is a cone and they propagate with constant angular widths and speeds, at least in their early phase of propagation, we have developed a technique (Michalek et al. 2003) which allows us to determine the following parameters: the linear distance $r$ of source location measured from the solar disk center, the angular distance $\gamma$ of source location measured from the plane of sky, the angular width $\alpha$ (cone angle =$0.5\alpha$) and the space velocity $V$ of a given HCME. A similar cone model was used recently by Xie et~al.~(2004) to determine the angular width and orientation of HCMEs. The present paper is divided into two parts. First, in the Section~2 we applied the cone model (Michalek et al. 2003) to obtain the space parameters of all HCMEs observed by the Solar and Heliospheric Observatory (SOHO) mission's Large Angle and Spectrometric Coronagraph (LASCO) until the end of 2002. In the Subsection~2.2 a short statistical analysis, based on the derived parameters, of HCMEs is presented (Fig.~1~-~Fig.~4). In the Section~3, we use these parameters to identify the most important factors determining geoeffectiveness of HCMEs and how they could be used for space weather forecast (Fig.~5~-~Fig.~17) | In this study we considered the geoeffectiveness of all full HCMEs observed by SOHO/LASCO coronagraphs from the launch in 1995 until the end of 2002. For $101/144$ ($70\%$) of full HCMEs we were able to find the source location, width and space velocity using the cone model (Michalek et al., 2003). We must be aware that the cone model is only rough simplification of real events. We know that not all CMEs are perfectly symmetric (Moran and Davila, 2004; Jackson et al., 2004). Most of CMEs could be approximate using cone model but probably for some of them this assumption is unrealistic. Fortunately technique presented by Michalek et al. does not demand perfect symmetry for CMEs. This approach requires measurements of sky-plane speeds and the moments of the first appearance of the halo CMEs above limb at only two opposite points. We are able to determine, with good accuracy, the space velocity and with of a given CME at least in the plane symmetry crossing CMEs at these points. When a given CME could be approximated by the cone model these derived parameters are valid for the entire CME. HCMEs originating very close to the disk center (mostly within a latitude of $\pm40^o$), are very wide (the average angular width $=120^o$) and are very fast (the average space speed $=1291km/s$). We find significant ($40\%$) increase in the average space velocities of HCMEs during the maximum of solar activity. These results could suggest that the HCMEs represent a special class of CMEs which are very wide and fast. It is important to note that this "class" of CMEs is defined due to artificial effect caused by coronagraphic observations. Events originating close to the disk center (from SOHO/LASCO point of view) must be wide and fast to appear as HCMEs in LASCO observations. This is not due to localization on the solar disk but due to oculting disk which not only blocks bright photospheric light but also eliminates some narrow and slow events. We have to emphasize that this effect mostly depends on the dimension of oculting disk but in less degree on the sensitivity of instrument. More sensitive instrument can record some poorer events (halos and also not halos so statistic will be similar) but could not register these events which never appear behind occulting disk. Potentially more sensitive instrument could register less energetic events (narrower and slower) and the average velocities and widths (for halos and whole population of CMEs) could be slightly lower but the main relation between the halos and whole population of events will be the same. Fortunately, poor events do not cause a big concern because they are not geoeffective. We do not expect, in the near future, any special programs devoted to looking for less energetic CMEs. Next scientific mission (STEREO) will be mostly dedicated to recognize 3D structure of CMEs. Such fast and wide CMEs are known to be associated with electron and proton acceleration by driving fast mode MHD shocks (e.g., Cane et al., 1987; Gopalswamy et al., 2002a). Using observations from Wind spacecraft, interplanetary magnetic clouds (MC) and geomagnetic disturbances associated to HCMEs were identified. The strength of geomagnetic storms, described by $D_{ST}$ and $Ap$ indices, is highly correlated with the source location and space velocity of a given event. Only HCMEs originating in the western hemisphere, close to the solar center and very fast (space velocity $\geq 1100km/s$) are likely to cause major geomagnetic storms ($D_ST<-150nT$). Slow HCMEs (space velocity $\leq 1100km/s$), even originating close to the solar center, may not cause severe geomagnetic disturbances. We have to note that there was one event (04 April 2000), which originated far from disk center and produced a severe geomagnetic storm ($D_{ST}=-288nT$). Probably this storm was not due to an ICME. It was caused by the sheath region ahead of the CME as was reported by Gopalswamy (2002b). We illustrated, using contour maps, how the derived HCME parameters can be useful for space weather forecast. We have to note that geoeffectiveness of events does not depend on their widths. During our study period we recognized $56/144$ ($30\%$) FHCMEs without any geomagnetic signature at Earth. This is significant population of FHCMEs. To distinguish them from the geoeffective events we considered the source locations and space velocities of HCMEs. When both the parameters are available, it becomes easier to assess the geoeffectiveness of HCMEs. We may say that fast FHCMs ($V>1200km/s$) originating close to the disk center ($|longitude|<30^o$) must be geoeffective. For such events there were no false alarms. But, even very fast events originating far from the disk center can be non-geoeffective. | 7 | 10 | 0710.4526 |
0710 | 0710.1303_arXiv.txt | We propose to study cosmic reionization using absorption line spectra of high-redshift Gamma Ray Burst (GRB) afterglows. We show that the statistics of the dark portions (gaps) in GRB absorption spectra represent exquisite tools to discriminate among different reionization models. We then compute the probability to find the largest gap in a given width range $[W_{\rm max}, W_{\rm max}+dW]$ at a flux threshold $F_{\rm th}$ for burst afterglows at redshifts $6.3 \le z \le 6.7$. We show that different reionization scenarios populate the $(W_{\rm max},F_{\rm th})$ plane in a very different way, allowing to distinguish among different reionization histories. We provide here useful plots that allow a very simple and direct comparison between observations and model results. Finally, we apply our methods to GRB~050904 detected at $z=6.29$. We show that the observation of this burst strongly favors reionization models which predict a highly ionized intergalactic medium at $z\sim 6$, with an estimated mean neutral hydrogen fraction $x_{\rm HI}=6.4\pm 0.3 \times 10^{-5}$ along the line of sight towards GRB~050904. | In the last few years, our knowledge of the high-$z$ Universe and in particular of the reionization process has been enormously increased mainly owing to the observation of quasars by the SDSS survey (Fan 2006) and CMB data (Hinshaw et al. 2007, Page et al 2007). Long gamma ray bursts (GRB) may constitute a complementary way to study the reionization process avoiding the proximity effects and possibly probing even larger redshifts. This has now become clear after the detection of five GRBs at $z\gsim 5$, over a sample of about 200 objects observed with the {\it Swift} satellite (Gehrels et al. 2004). The current record holder is GRB~050904 at $z = 6.29$ (Tagliaferri et al. 2005, Kawai et al. 2006). Totani et al. (2006) have used this object to constrain the ionization state of the intergalactic medium (IGM) at high redshift by modeling its optical afterglow spectrum. They report the evidence that the IGM was largely ionized already at $z=6.3$. The best-fit neutral hydrogen fraction is consistent with zero with upper limit $x_{\rm HI}<0.17$ ($<0.6$) at 68\% (95\%) C.L. Several authors (Natarajan et al. 2005, Daigne, Rossi \& Mochkovitch 2006, Bromm \& Loeb 2006, Salvaterra et al. 2007a) have computed the number of high--$z$ GRBs detectable by {\it Swift}. In spite of model details, all these different studies consistently predict that a non-negligible fraction (up to $\sim 10$\%) of all observed GRBs should lie at very high redshift. On the basis of these results, several GRBs at $z\ge 6$ should be observed by {\it Swift} in the near future. Depending on the {\it Swift}/BAT trigger sensitivity and on model details, Salvaterra \& Chincarini (2007) found that $\sim 2-8$ GRBs can be detected above this redshift during every year of mission. From the observational point of view, Ruiz-Velasco et al. (2008) find that the fraction of {\it Swift} GRBs at $z>6$ should not exceed the conservative upper limit of $19\%$. Moreover, future X-- and Gamma--ray missions will increase rapidly the sample of high--$z$ GRBs, possibly up to $z\sim 10$ (Salvaterra et al. 2007b). Once detected, spectroscopic follow-up observations of high--$z$ GRBs require a rapid trigger of 8-meter, ground based telescopes. This can be done by pre-selecting reliable candidates on the basis of some promptly-available information provided by {\it Swift}, such as burst duration, photon flux, the lack of detection in the UVOT $V$-band, and the low Galactic extinction (Campana et al. 2007, Salvaterra et al. 2007a). We note that cosmological time dilation helps keeping the flux bright, since observations will sample the afterglow early phases, even a few days after trigger. The spectra of high-redshift sources (as QSOs and GRBs) bluewards of the Ly$\alpha$ are characterized by dark portions (gaps) produced by intervening neutral hydrogen along the line of sight. The use of various gap statistics in QSO spectra has been recently recognized as a very powerful tool to constrain the IGM ionization state (Paschos \& Norman 2005; Fan et al. 2006; Gallerani, Choudhury \& Ferrara 2006, hereafter G06; Gallerani et al. 2007, hereafter G07). For example, by comparing the statistics of these spectral features in a sample of 17 observed QSOs with Ly$\alpha$ forest simulations, G07 concluded that the HI fraction, $x_{\rm{HI}}$, evolves smoothly from $10^{-4.4}$ at $z=5.3$ to $10^{-4.2}$ at $z=5.6$, with a robust upper limit $x_{\rm{HI}} < 0.36$ at $z=6.3$. These results encourage the application of such analysis to GRBs. There are several advantages promised by such an attempt. First, GRBs are soon expected to be found at redshifts higher than those typical of QSOs; second, and contrary to the massive hosts of QSOs, they reside in ``average" cosmic regions only marginally affected by local ionization effects or strong clustering; finally, they are bright and their afterglow spectra closely follows a power-law, making continuum determination much easier In agreement with the WMAP3 results (Spergel et al. 2007) we assume a flat universe with $\Omega_m=0.24$, $\Omega_{\Lambda}=0.76$, $\Omega_bh^2=0.022$, $h=0.73$. The parameters defining the linear dark matter power spectrum are $n=0.95$, $\de n/\de \ln k=0$, $\sigma_{8}=0.82$. Mpc are physical unless differently stated. | We have proposed to investigate cosmic reionization using absorption line spectra of high--$z$ GRB afterglows. The evolution of the largest gap width (LGW) as a function of the flux threshold $F_{\rm th}$ used to define gaps shows marked differences in the two reionization models favored by present data. In particular, (i) the LGW is typically $\sim 2$ times wider in the LRM than in the ERM, for a fixed $F_{\rm th}$; (ii) the difference between ERM and LRM in terms of $(W_{\rm max}, F_{\rm th})$ increases with $z$; (iii) an overall shift towards larger $W_{\rm max}$ values for a fixed $F_{\rm th}$ is found in both models towards higher redshifts. This analysis has shown that we can robustly distinguish among different reionization histories: improved results can be obtained if data is collected promptly after burst detection and by using GRBs at the highest available redshifts. A direct comparison of the model with data can be carried on in terms of the probability to find the largest gap in a given width range for burst afterglows at $z=z_{\rm GRB}$. When applied to the only known GRB at $z>6$, i.e. GRB~050904 at $z_{\rm GRB}=6.29$, a clear indication is obtained that reionization must have occurred well before $z=6$. We find that the observed LGW in the GRB~050904 afterglow spectrum are consistent with $x_{\rm HI}=6.4\pm 0.3\times 10^{-5}$. This result is in agreement with previous measurements by Totani et al. (2006), who find that $x_{\rm HI}$ is bound to be $x_{\rm HI}<0.17$ ($<0.6$) at 68\% (95\%) C.L. Some gaps/peaks in absorption spectra could be due to DLAs/HII regions intervening along the LOS towards the background source (e.g. the DLAs/transverse proximity effect detection by Totani et al. 2006/Gallerani et al. 2007). Such contaminants could affect neutral hydrogen measurements, if the foreground sources are not properly removed. In this work we have applied the gap statistics to the GRB 050904 observed flux. However, the same analysis can be done by using the transmitted flux $F_{\rm obs}/F(\nu)$, i.e. normalizing the observed flux to the continuum. Following this approach, we provide useful plots (Fig. 2) which allow a straightforward comparison between our model results and future observations of $z>6$ GRBs afterglow spectra. It is worth noting that the gap statistics with a variable flux threshold can be readily applied to other background sources, as QSOs. We plan to apply this analysis to the Fan et al. (2006) data, thus improving the results obtained by Gallerani et al. (2007). The advantage of varying the flux threshold is that this technique allows to study the response of the transmitted flux both to the high and low end tail of the IGM density field. In fact, high (low) $F_{\rm th}$ values probe regions corresponding to low (high) overdensities. In spite of the fact that QSO data are presently more abundant, GRBs are expected to be found at higher redshift, as they should result from the death of (early) massive stars (e.g. Abel et al. 2002, Schneider et al. 2002). Larger samples of GRB afterglow spectra at $z\ge 6$, likely available in the near future, will allow a statistically significant analysis and possibly to reconstruct the cosmic reionization history. At redshifts approaching reionization epoch, gaps in afterglow absorption spectra are expected to become as large as the spectral region between the Ly$\alpha$ and Ly$\beta$ emission lines, thus obscuring completely the GRBs optical counterpart. | 7 | 10 | 0710.1303 |
0710 | 0710.5300_arXiv.txt | Using big bang nucleosynthesis and present, high-precision measurements of light element abundances, we constrain the self-gravity of radiation pressure in the early universe. The self-gravity of pressure is strictly non-Newtonian, and thus the constraints we set provide a direct test of this prediction of general relativity and of the standard, Robertson-Walker-Friedmann cosmology. | Certain aspects of general relativity are well tested. For example, the Schwarzschild metric has been quantitatively verified in the weak-field limit on small scales, e.g., the Solar system \citep{Shapiro64, Bertotti03} and binary radio pulsars \citetext{e.g., \citealp{Hulse75, Taylor79, Weisberg84}}; and on galaxy scales \citep[e.g.,][]{Bolton06}. In another fundamental test of general relativity, the existence of gravity waves has been established \citetext{e.g., \citealp{Taylor82}, \citealp{Weisberg84}}. General relativity theory, utilizing the Robertson--Walker metric \citetext{\citealp{Friedmann22}, \citealp{Robertson35}, \citealp{Walker36}} leads to the Friedmann equations \citetext{\citealp{Friedmann22, Lemaitre27}} which govern the expansion behavior of a homogeneous, isotropic Universe. However, it is probably fair to say that the Friedmann equations, while providing a self-consistent and highly successful framework for cosmology, have not been subjected to extensive, independent testing. In this paper we show that one particular aspect of the Friedmann equations, the self-gravity of pressure, can be tested quantitatively. The development of big-bang nucleosynthesis (BBN) codes \citep{Wagoner67, Kawano92}, coupled with measurements of the relevant nuclear reaction rates \citep{Caughlan88, desc04}, have allowed observations of light element abundances to become powerful tools with which to investigate the early evolution of the universe. Computational predictions over a wide range of parameter space, when compared with primordial abundances inferred from observations, have yielded constraints on the current-epoch baryon density \citep{Wagoner73, Yang79, KS04, Steigman07}, neutrino physics \citep{SSG77, Yang79, KS04, Steigman07}, the fine structure constant \citep{Bergstrom99}, the gravitational constant \citep{Steigman76, Yang79, Copi04, Steigman07}, primordial magnetic fields \citep{Kernan96}, the universal lepton asymmetry \citep{Wagoner67, KS04, Steigman07} and other parameters of astrophysical interest. Increasingly accurate measurements of element abundances, as well as improved understanding of the processes (i.e., stellar and galactic nucleosynthesis) which have altered the original abundances, allow these restrictions to be continually refined. Deuterium abundances \citep{GeissReeves72, Omeara06}, helium abundances \citep{HoyleTayler64, IT04, Izotov07}, and lithium abundances \citep{RNB00, Asplund06} have all been well measured, although the inferred primordial abundances are subject to large and often difficult to quantify systematic uncertainties. More recently, observations of the cosmic microwave background (CMB) have yielded an independent estimate of $\eta$, the baryon to photon ratio at a much later epoch in the evolution of the Universe \citep{Spergel07}. | As illustrated in Figure 1, the combined constraints are, within the uncertainties, consistent with the general relativity prediction of $\chi = 1$ and the independent (of BBN) WMAP constraint on $\eta_{10}$ of $6.1 \pm 0.2$ \citep{Spergel07} which corresponds to $\chi = 1$. For the \citet{PLP07} choice of Y$_{P}$, $\chi = 1.00\pm 0.14$, while for the Steigman \cite{Steigman07} helium abundance, $\chi = 0.84\pm 0.25$. Note that the data strongly exclude $\chi = 0$. The current light element observations and BBN computations have provided a test of the general relativistic self-gravity of pressure. Since the modification of GR we are testing corresponds, for the radiation-dominated evolution appropriate for BBN, to an overall multiplicative factor of the product of Newton's gravitational constant and the radiation density, $G\rho \rightarrow G\rho \left({\frac{1 + \chi}{ 2}}\right)$, our result is equivalent to the BBN constraint on the variation of Newton's constant or, alternatively, to a modification of the radiation energy density as parameterized by the effective number of neutrinos (see \S1 for references). \begin{equation} \frac{1 + \chi}{2} = 1 + \frac{\Delta G}{G} = 1 + \frac{7\Delta N_{\nu}}{ 43} \end{equation} Assuming that these other parameters take on their standard-model values ($\Delta G = \Delta N_{\nu} = 0$), the self-gravity of the radiation (photons and neutrinos) pressure during the BBN epoch has been constrained quantitatively. | 7 | 10 | 0710.5300 |
0710 | 0710.1715_arXiv.txt | AA~Dor is an eclipsing, close, post common-envelope binary (PCEB). We present a detailed spectral analysis of its sdOB primary star based on observations obtained with the Far Ultraviolet Spectroscopic Explorer (FUSE). Due to a strong contamination by interstellar absorption, we had to model both, the stellar spectrum as well as the interstellar line absorption in order to reproduce the FUV observation well and to determine the photospheric parameters precisely. | The eclipsing binary system AA~Dor consists of a sdOB with \mbox{$T_{\rm{eff}}=42\,\rm{kK}$} and an unseen low-mass companion. Recent studies of its primary star have shown discrepancies in the determination of the surface gravity $g$. \citet{Rauch00} obtained \mbox{$\log g\hspace{-0.5mm} =\hspace{-0.5mm} 5.21$} by a NLTE spectral analysis of optical and ultraviolet data. The results of radial-velocity and lightcurve analyses by \citet{Hilditch96,Hilditch03} indicate a higher value of \mbox{$\log g\hspace{-0.5mm} =\hspace{-0.5mm} 5.45\,-\,5.51$}. Twelve high S/N-ratio FUSE observations were performed in August 2003 and June 2004. With about 200 secs each, the exposure times were chosen short in order to reduce the smearing by orbital motion. The FUV range includes the hydrogen Lyman series which in general is sensitive to changes in $\log g$. However, a precise determination of the surface gravity and further stellar parameters is hampered by strong interstellar absorption lines. We used our NLTE code TMAP \citep[T\"ubingen Model Atmosphere Package, ][]{Werner03} to calculate plane-parallel, line-blanketed model atmospheres in radiative, hydrostatic and statistical equilibrium. We included calcium and all elements of the iron group (Sc, Ti, V, Cr, Mn, Fe, Co, Ni) into our calculations. Due to the large amount of transitions, those elements were treated in a statistical approach \citep{rd2003}. Atomic data were taken from NIST (National Institute of Standards and Technology, \texttt{http://physics.nist.gov}), the Opacity Project \citep{Seaton94}, and the Iron Project \citep{Hummer93}. In case of the iron-group elements, energy levels and oscillator strengths stem from Kurucz's line lists \citep{Kurucz91}. Only weak lines of those elements are detected in the FUSE spectra, which form absorption troughs due to rotational broadening (Fig. \ref{fig:POSLIN}). \begin{figure}[ht!] \begin{center} \includegraphics[width=0.7\textwidth]{fleig_01.eps}\vspace{-2mm} \caption[]{ Identification of iron-group lines. \citet{Kurucz91} gives only a few lines with reliably measured wavelengths (POS lines, blue, shifted to the top) which makes the identification of isolated lines difficult. The addition of iron-group lines with theoretical positions (Kurucz's LIN lines, red) improves the overall-fit to the observation. The synthetic spectra are convolved with a rotational profile of $v_{\rm{rot}}=35\,\rm{km\,s^{-1}}$.}\vspace{-5mm} \label{fig:POSLIN} \end{center} \end{figure} | The effective temperature was derived using the ionization equilibrium of N\,{\scriptsize III} and N\,{\scriptsize IV}. The best fit is achieved for $T_{\rm{eff}}\hspace{-0.5mm}=\hspace{-0.5mm}40 \pm 3\,\mathrm{kK}$ (Fig\@. \ref{fig:Teff}). Due to the interstellar contamination of the spectra (Fig\@. \ref{fig:ISM_MOD}), a precise determination of the surface gravity was not possible (Fig\@. \ref{fig:logg}). However, a hint towards a higher value of $\log g$ is given by the rotational velocity of the primary (Fig\@. \ref{fig:vrot}). With $v_{\rm{rot}}=35 \pm 5\,\rm{km\,s}^{-1}$, the primary rotates slower than bound ($v_{\rm{rot}}=45.7\,\rm{km\,s}^{-1}$). \citet{Rauch00} assumed bound rotation for the calculation of the primary's spectrum and thus, one can expect a slightly higher value of $\log g$. Calculating the surface gravity directly via $\log\,g=\log\, [(GM_1)/r_1^2]$, $r_1\propto v_{\rm{rot}}\cdot P$, returns a surface gravity of 5.44, which is in agreement with the value found by \citet{Hilditch03}. \begin{figure}[ht!] \begin{center} \includegraphics[width=0.7\textwidth]{fleig_02.eps}\vspace{-2mm} \caption[]{Determination of $T_{\rm{eff}}$ from the \mbox{N\,{\scriptsize{III}} / N\,\scriptsize{IV}} ionization equilibrium. While the \mbox{N\,\scriptsize{IV}} line (left panel) appears saturated within the relevant $T_{\rm{eff}}$, the \mbox{N\,\scriptsize{III}} multiplet (right) matches the observation best at $T_{\rm{eff}} = 40\,\mathrm{kK}$. } \label{fig:Teff} \end{center} \end{figure} \begin{figure}[ht!] \begin{center} \includegraphics[width=0.7\textwidth]{fleig_03.eps}\vspace{-2mm} \caption[]{Comparison of a pure photospheric model spectrum (blue) and the combined photospheric + ISM model (red) within a section of the FUSE observation. Interstellar lines are labeled with an asterisk.} \label{fig:ISM_MOD} \end{center} \end{figure} \begin{figure}[ht!] \begin{center} \includegraphics[width=0.9\textwidth]{fleig_04.eps}\vspace{-2mm} \caption[]{Log $g$ - dependence of the first four lines of the Lyman series. Lyman $\alpha$ is compared to IUE observations, Lyman $\beta - \delta$ to FUSE observations. An interstellar H\,{\scriptsize I} column density of $\log N_\mathrm{H\,{\footnotesize I}} = 20.3$ and an extinction of $E_\mathrm{B-V}=0.04$ are considered. Interstellar absorption prevents a precise determination of the continuum and the course of the line wings. All continua are normalized to that of the synthetic spectra with $\log g=5.20$.} \label{fig:logg} \end{center} \end{figure} \begin{figure}[ht!] \begin{center} \includegraphics[width=0.7\textwidth]{fleig_05.eps}\vspace{-2mm} \caption[]{ Comparison of the \mbox{C\,{\scriptsize III} $\lambda\lambda 1175\,\mathrm{\AA}$} multiplet with the FUSE observation at three different rotational velocities $v_{\rm{rot}}$. The best fit is achieved with $v_{\rm{rot}}=35\,\rm{km\,s}^{-1}$.} \label{fig:vrot} \end{center} \end{figure} \begin{figure}[ht!] \begin{center} \includegraphics[width=0.9\textwidth]{fleig_06.eps}\vspace{-2mm} \caption[]{Detection of phosphorus and sulfur. The resonance doublets of \mbox{S\,{\scriptsize VI} $\lambda\lambda 933.38, 944.52\,\mathrm{\AA}$} and \mbox{P\,{\scriptsize V} $\lambda\lambda 1117.98, 1128.01\,\mathrm{\AA}$} are well reproduced at solar abundances. Since interstellar absorption contributions cannot be excluded, these abundances are upper limits.} \label{fig:doublets} \end{center} \end{figure} | 7 | 10 | 0710.1715 |
0710 | 0710.3704_arXiv.txt | {} { % We study the evolution of the galaxy population up to $z\sim 3$ as a function of its colour properties. In particular, luminosity functions and luminosity densities have been derived as a function of redshift for the blue/late and red/early populations.} {We use data from the GOODS-MUSIC catalogue which have typical magnitude limits $z_{850}\leq 26$ and $Ks\leq 23.5$ for most of the sample. About 8\% of the galaxies have spectroscopic redshifts; the remaining have well calibrated photometric redshifts derived from the extremely wide multi-wavelength coverage in 14 bands (from the U band to the Spitzer $8 \mu$m band). We have derived a catalogue of galaxies complete in rest-frame B-band, which has been divided in two subsamples according to their rest-frame U-V colour (or derived specific star formation rate, SSFR) properties.} {We confirm a bimodality in the U-V colour and SSFR of the galaxy sample up to $z\sim 3$. This bimodality is used to compute the LFs of the blue/late and red/early subsamples. The LFs of the blue/late and total samples are well represented by steep Schechter functions evolving in luminosity with increasing redshifts. The volume density of the LFs of the red/early populations decreases with increasing redshift. The shape of the red/early LFs shows an excess of faint red dwarfs with respect to the extrapolation of a flat Schechter function and can be represented by the sum of two Schechter functions. Our model for galaxy formation in the hierarchical clustering scenario, which also includes external feedback due to a diffuse UV background, shows a general broad agreement with the LFs of both populations, the larger discrepancies being present at the faint end for the red population. Hints on the nature of the red dwarf population are given on the basis of their stellar mass and spatial distributions.} {} | The evolution of galaxy Luminosity Function (LF) is one of the main tools to study the structure evolution through the cosmic time. The advent of large surveys has allowed the analysis of sub-samples of galaxies selected as a function of their morphological, spectroscopic or colour properties \citep[][elsewhere G05]{strateva2001,norberg2002,madgwick2002,wolf2003,willmer,bell2004,baldry2004,hogg2004,weiner2005,blanton2005,ilbert2006,marchesini2006,baldry2006,driver2006,cucciati06,cirasuolo2006,arnouts2007,giallo2005}. In fact these kind of studies allow us to probe the evolution of galaxies having different star formation histories. Of special interest are the studies concerning the statistical properties of galaxies selected on the basis of their intrinsic colour distribution. This distribution appears bimodal up to $z\sim 2-3$ \citep{baldry2004,blanton2005,giallo2005} and separates the galaxies in two populations, red early types vs. blue late types. It has been shown that this spectral classification is roughly consistent with the correspondent morphological classification (bulge vs. disc dominated) at least at low and intermediate redshifts \citep{strateva2001,weiner2005}. This bimodal colour distribution can find a natural explanation in hierarchical models for galaxy formation \citep{menci2005,menci2006} where two distinct populations arise in the colour distribution based on two different star formation histories affected by the feedback effects produced by the SN and AGN activities \citep{menci2006}. However the effect of environmental density on the paths of galaxy evolution can have a fundamental role. In this context it is not clear whether the evolutionary history of galaxies is originated by a {\it nurture} senario (galaxy properties are affected by environment through physical mechanisms acting on galaxies) or by a {\it nature} scenario \citep[the evolution is driven by the initial condition established during the formation epoch of galaxies, e.g. ][]{mateus2007,cooper2007}. Recent studies have estimated the shape and evolution of the LF of galaxies selected according to their bimodal colour distribution using both the large local Sloan survey \citep{baldry2004,blanton2005}, and other surveys at intermediate and high redshifts \citep{bell2003,giallo2005,faber,willmer,ilbert2006}. Their results show the red LF evolving mildly in density up to $z \sim 1$ with a quite flat shape at the faint-end, although the evaluation of the faint-end slope of the red LF remains an open issue especially at intermediate and high redshifts where the present surveys do not constrain the faint slope very well \citep{faber,bell2003}. In G05 we studied the red and blue LFs, using the properties of bimodality in colour and in specific star formation rate (SSFR), with a complete but relatively small sample of galaxies selected in the rest-frame B-band from low to high redshifts. We showed that the bimodality extends at least up to $z\sim 2.5$. We also found that the red/early galaxies decrease in their luminosity density by a factor $\sim 5-6$ from $z\sim 0.5 $ to $z\sim 2.5-3$ in broad agreement with the hierarchical cold dark matter model. These results provided a first picture of the evolution of the red and blue LFs up to high redshifts relaying on a relatively deep but small sample. For a more reliable picture a wider sample at high redshift is clearly needed. For this purpose larger areas with deep near-IR imaging are required. Thanks to the wide area ($\sim 140\ arcmin^2$) and to the deep near-IR observations, the GOODS-South survey provides a good starting point for the study of the galaxy properties at high redshift. In particular, the inclusion of the deep IR observations obtained with the Spitzer telescope represent a useful constraint for the estimate of the physical properties of galaxies at high redshift. Last but not least the extensive spectroscopic follow up obtained in this field provides a wide set of spectroscopic redshifts. From this public data set we have obtained a multi-colour catalogue of galaxies we named GOODS-MUSIC \citep[GOODS MUlticolour South Infrared Catalog, ][]{grazian}. This catalog, where galaxies are selected both in the $z_{850}$ and $Ks$ bands, contains information in 14 bands from the U to the Spitzer 8 $\mu m$ band, and all the available spectroscopic information. For all the objects without spectroscopic information we have obtained well calibrated photometric redshifts by means of our photometric redshift code \citep{fontana00}. The GOODS-MUSIC catalogue has been already used to investigate the physical and clustering properties of high redshift galaxies \citep{fontana06,grazian2006_eros,grazianc,pentericci2007,castellano2007}. With the present paper we study the galaxy LFs of the red and blue populations, enlightening evolutionary features which are characteristic of the two populations. The paper is organised as follows: in section 2, we describe the basic features of our dataset. In section 3, we describe the bimodality properties of the sample and we define the loci of minimum for the selection of the red/early and the blue/late sub-samples as a function of $z$. In section 4, we compute the shape and evolutionary properties of the LFs and the luminosity density of both populations. Section 5 is devoted to the analysis of the physical properties of the faint early population. All the magnitudes used in the present paper are in the AB system. An $\Omega_\Lambda=0.7$, $\Omega_M=0.3$, and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ cosmology is adopted. | We have used a composite sample of galaxies selected in deep NIR images, obtained from the GOODS public survey, to study the evolution of the galaxy LFs of red/early and blue/late galaxy populations selected using the colour and SSFR statistical properties of the sample. The observed $U-V$ colour and SSFR distributions show a clear bimodality up to $z\sim 3$, confirming the results obtained in G05 at a higher level of statistical confidence. We found a trend with redshift for the colour magnitude distribution with an intrinsic blueing of about 0.15 mag. in the redshift interval $z=0.4-2.0$ for both populations. This observed bimodality can be explained in a hierarchical clustering scenario as due to the different star formation histories of the red/early and blue/late galaxies, see e.g. \cite{menci2006}. For the total and the blue/late sample the LF is well described by a Schechter function and shows a mild luminosity evolution in the redshift interval $z=0.2-1$ (e.g. $\Delta M^* \sim 0.7$ for the total sample; $\Delta M^*\sim 1$ for the blue/late fraction), while at higher redshifts the LFs are consistent with no evolution. A comparison with our hierarchical CDM model shows a good agreement at bright and intermediate magnitudes. A better agreement of the model has also been found at fainter magnitudes due to the suppression of star formation in small objects by the action of an ionising UV background. The shape of the red/early luminosity function is better constrained only at low and intermediate redshifts and it shows an excess of faint red dwarfs with respect to the extrapolation of a flat Schechter function. In fact a minimum around magnitude $M_B(AB)=-18$ is present together with an upturn at fainter magnitudes. This peculiar shape has been represented by the sum of two Schechter functions. We found that the bright one is constant up to $z\sim 0.7$ beyond which it decreases in density by a factor $\sim 5$ (10 for the early galaxies) up to redshift $z\sim 3.5$. The comparison with our hierarchical CDM model shows that, although the predicted LF has a slight flattening at intermediate luminosity, the model still overpredicts the LF at faint magnitudes. The bright end of blue and red LFs at low and intermediate redshifts is in good agreement with recent estimates from the DEEP2 spectroscopic survey. As a consequence of this complex evolutionary behaviour, the luminosity densities of the relatively bright ($M_B(AB)<-20.2$) red/early and blue/late galaxies show a bifurcation beyond redshift $z\sim 1$. Indeed the LD of the blue/late population keeps increasing up to $z\sim 3.5$, while the luminosity density of red/early galaxies decreases by a factor $\sim 2-3$ respectively in the $z=1-3.5$ interval. To derive hints on the nature of the galaxies responsible for the peculiar shape of the red/early LF, we have performed an analysis of their stellar masses and spatial distribution. We found that the early galaxies have systematically higher stellar masses with respect to the late ones for a given B band luminosity. Brighter early galaxies have a spatial distribution more concentrated in higher density regions if compared to the late ones of the same luminosity class. On the contrary, fainter early and late galaxies show a very similar spatial distribution. Thus, the different environmental properties do not seem to be the main responsible for the difference in shape at intermediate magnitudes between the blue and red LFs. The latter seem to stem from the different star formation and feedback histories corresponding to different possible merging trees (evolutionary paths) leading to the final assembled galaxy; this specific history, driving the evolution of the star formation, leads to the different $M/L$ ratios characterising the different properties of blue/late and red/early galaxies. In summary, the peculiar shape of the red LF is mainly driven by the {\it nature} of the galaxy merging tree rather than by the {\it nurture} where the galaxy has grown. | 7 | 10 | 0710.3704 |
0710 | 0710.1537_arXiv.txt | {The rapidly varying non-thermal X-ray emission observed from Sgr A$^{\star}$ points to particle acceleration taking place close to the supermassive black hole. The TeV $\gamma$-ray source HESS\,J1745$-$290 is coincident with Sgr A$^{\star}$ and may be closely related to the X-ray emission. Simultaneous X-ray and TeV observations are required to elucidate the relationship between these two objects. Here we report on joint H.E.S.S./Chandra observations in July 2005, during which an X-ray flare was detected. Despite a factor $>10$ increase in the X-ray flux of Sgr~A$^{\star}$, no evidence is found for an increase in the TeV $\gamma$-ray flux. We find that an increase of the $\gamma$-ray flux of a factor 2 or greater can be excluded at a confidence level of 99\%. This finding disfavours scenarios in which the bulk of the $\gamma$-ray emission observed is produced close to Sgr A$^{\star}$. } \begin{document} \newcommand{\astar}{Sgr~A$^{\star}$} | The existence of a supermassive ($3.6\pm0.3 \times 10^{6}$ solar mass) black hole at the centre of our galaxy has been inferred using measurements of stellar orbits in the central parsec (see e.g. ~\cite{GC:Eisenhauer05}). The supermassive black hole (SMBH) is coincident with the faint radio source: Sgr~A$^{\star}$. The compact nature of Sgr~A$^{\star}$ has been demonstrated both by direct VLBI measurements~\cite{GC:Shen05} and by the observation of X-ray and near IR flares with timescales as short as a few minutes (see for example~\cite{GC:Eckart06,GC:Porquet03}). Variability on such short timescales limits the emission region (via causality arguments) to within $<10$ Schwarzchild radii of the black hole. X-ray flares from \astar\ have reached fluxes of $4\times10^{35}$ erg s$^{-1}$, two orders of magnitude brighter than the quiescent flux~\cite{GC:Porquet03, GC:Baganoff03}, and exhibit a range of spectral shapes~\cite{GC:Porquet03}. Several models exist for the origin of this variable emission, all of which invoke non-thermal processes close to the event horizon of the central black hole to produce a population of relativistic particles. Model independent evidence for the existence of ultra-relativistic particles close to Sgr~A$^{\star}$ can be provided by the observation of TeV $\gamma$-rays from this source. Indeed, TeV $\gamma$-ray emission has been detected from the Sgr~A region by several ground-based instruments~\cite{GC:Whipple04,GC:CANGAROO,HESS:gc04, GC:MAGIC06}. The most precise measurement of this source, HESS\,J1745$-$290, are those made using the H.E.S.S. telescope array. The centroid of the source is located $7'' \pm 14_{\mathrm{stat}}'' \pm 28_{\mathrm{sys}}''$ from Sgr~A$^{\star}$, and has an rms extension of $<1.2'$~\cite{HESS:gcprl}. TeV emission from \astar\ is expected in several models of particle acceleration in the environment of the black hole. In some of these scenarios \cite{Levinson, AhNer1} TeV emission is produced in the immediate vicinity of the SMBH and variability is expected. In alternative scenarios particles are accelerated at \astar\ but radiate in within the central $\sim$10~parsec region~\cite{AhNer2}, or are accelerated at the termination shock of a wind driven by the SMBH \cite{AtDer}. However, several additional candidate objects exist for the origin of the observed $\gamma$-ray emission. The radio centroid of the supernova remnant (SNR) Sgr~A East lies $\sim 1'$ from Sgr~A$^{\star}$, only marginally inconsistent with the position of the TeV source given in \cite{HESS:gcprl}. Shell-type SNR are now well established TeV $\gamma$-ray sources~\cite{HESS:rxj1713,HESS:velajnr} and several authors have suggested Sgr A East as the origin of the TeV emission (see for example \cite{Melia}). However, recent improvements in the statistical and systematic uncertainties of the centroid of HESS\,J1745$-$290 effectively exclude Sgr~A East as the dominant $\gamma$-ray source in the region~\cite{VanEldik}. The recently discovered pulsar wind nebula candidate G\,359.95-0.04~\cite{GC:Wang06} lies only 9 arcseconds from \astar\ and can plausibly explain the TeV emission~\cite{GC:Hinton07}. Particle acceleration at stellar wind collision shocks within the central young stellar cluster has also been hypothesised to explain the $\gamma$-ray source~\cite{GC:Quataert05}. Finally, an origin of this source in the annihilation of WIMPs in a central dark matter cusp has been extensively discussed~\cite{GC:Hooper04,GC:Profumo05,HESS:gcprl}. Given the limited angular resolution of current VHE $\gamma$-ray telescopes, the most promising tool for identification of the TeV source is the detection of \emph{correlated variability} between the $\gamma$-ray and X-ray and/or NIR regimes. A significant increase of the flux of HESS\,J1745$-$290 simultaneous with a flare in wavebands with sufficient angular resolution to isolate \astar, would provide an unambiguous identification of the $\gamma$-ray source. Therefore, whilst not all models for TeV emission from Sgr~A$^{\star}$ predict variability of the VHE source, coordinated IR/keV/TeV observations can be seen as a key aspect of the ongoing program to understand the nature of this enigmatic source. | For the first time simultaneous TeV $\gamma$-ray observations have been presented for a period of X-ray activity of \astar. The non-detection of an increase in the TeV flux provides an important constraint on scenarios in which the source HESS\,J1745-290 is associated with the supermassive black hole. | 7 | 10 | 0710.1537 |
0710 | 0710.4254_arXiv.txt | The Atacama Large Millimeter Array (ALMA) will consist of up to 64 state-of-the-art sub-mm telescopes, subject to stringent performance specifications which will push the boundaries of the technology, and makes testing of antenna performance a likewise challenging task. Two antenna prototypes were evaluated at the ALMA Test Facility at the Very Large Array site in New Mexico, USA. The dynamic behaviour of the antennas under operational conditions was investigated with the help of an accelerometer system capable of measuring rigid body motion of the elevation structure of the antenna, as well as a few low-order deformation modes, resulting in dynamic performance numbers for pointing stability, reflector surface stability, path length stability, and structure flexure. Special emphasis was given to wind effects, one of the major factors affecting performance on timescales of seconds to tens of minutes. Though the accelerometers could not directly measure antenna performance on timescales longer than 10 seconds, it was possible to use them to derive antenna properties which allowed extrapolation of the wind-affected performance to timescales of 15 minutes and longer. This paper describes the accelerometer system, its capabilities and limitations, and presents the dynamic performance results of the two prototype antennas investigated. In addition to verification of the performance requirements, we investigated the vibration environment on the antennas, relevant for vibration-sensitive equipment for use on the ALMA antennas, the lowest eigenfrequencies for the antennas, and the sensitivity to vibration generated by similar antennas operating nearby. This work shows that seismic accelerometers can successfully be used for characterisation of antenna dynamics, in particular when complemented with simultaneous wind measurements and measurements for static performance. The versatility of the accelerometers also makes them a valuable tool for troubleshooting of unforeseen antenna features. | \label{intro} The enormous growth of radio astronomy in the millimeter and submillimeter wavelength regime (frequencies from 100 - 1000 GHz) over the last 25 years has been made possible both by the emergence of sensitive detectors like SIS-diodes and HEB-devices and the construction of ever larger and more accurate radio telescopes. Reflector antennas for this wavelength region have been built from relatively small (up to 15 m diameter), extremely accurate (reflector accuracy \mbox{15 - 25 $\mu$m}) submillimeter telescopes to larger (20 - 45 m) millimeter antennas with surface accuracy from \mbox{75 - 150 $\mu$m}. Current projects in this area are the 50 m diameter Large Millimeter Telescope (LMT) under construction in Mexico \cite{Kaercher2000}, which is aiming to reach \mbox{75 $\mu$m} rms surface and 1 arcsec pointing accuracy and the Atacama Large Millimeter Array (ALMA). ALMA is a global collaboration of North America (USA and Canada) and Europe (European Southern Observatory) with contributions from Japan to build a powerful millimeter wavelength aperture synthesis array at a 5000 m high plateau in Northern Chile. The instrument will consist of up to 64 high accuracy Cassegrain reflector antennas of 12 m diameter with a surface rms accuracy of \mbox{20 $\mu$m} and a pointing and tracking accuracy and stability of 0.6 arcseconds, all under the severe operational conditions of the high site. This telescope will operate over the entire frequency range from 30 to 950 GHz. These specifications are among the most severe ever realised in radio telescopes and force the designer and manufacturer to push the boundaries of the technology. At the same time, it is becoming increasingly difficult for the contractor and the customer to quantitatively and reliably evaluate the performance characteristics of these instruments. At the longer wavelengths radio astronomers have developed and used methods of antenna evaluation which are based on the use of strong cosmic radio sources and astronomical observing techniques \cite{Baars1973}. These measurements are only of limited use at the short millimeter wavelengths at frequencies above 100 GHz. The number of suitable cosmic test sources is severely limited because of their intrinsic weakness and the sensitivity limitations of the relatively small antennas. \pubidadjcol For the ALMA Project the partners decided early to obtain two prototype antennas from different design and fabrication groups to increase the chance of achieving the desired performance. The prototype antennas tested here, one designed and constructed by VertexRSI, the other by a consortium of Alcatel and European Industrial Engineering, hereafter referred to as AEC, are similar in overall design and built to meet identical requirements, though significant differences exist in the approaches taken to meet these requirements. The antennas, together with a third one from Japan, are located at the ALMA Test Facility (ATF), on the site of the Very Large Array (VLA) in New Mexico, USA (Figures~\ref{fig:vertexrsi} and \ref{fig:atf}). For more information on the antennas and the evaluation program, see \cite{Mangum2006}. \begin{figure} \resizebox{\hsize}{!}{ \includegraphics[scale=0.48,angle=0]{Vertex-astroph.jpg}} \caption{The VertexRSI ALMA prototype antenna.} \label{fig:vertexrsi} \end{figure} \begin{figure} \resizebox{\hsize}{!}{ \includegraphics[scale=0.65,angle=0]{DSC01896-astroph.jpg}} \caption{The ALMA Test Facility, with the AEC antenna (left), VertexRSI antenna (middle) and Mitsubishi antenna (right).} \label{fig:atf} \end{figure} An international team of radio astronomers was formed to subject the antennas to an extensive evaluation program. In preparing their tasks this group determined that additional measurements and test methods and instruments beyond the usual astronomical testing would be needed to check the very stringent specification of the antennas. A particularly important, but difficult to measure quantity is the accuracy with which the antenna can be pointed at arbitrary positions on the sky and the stability with which such a pointing can be maintained under the influence of variations in temperature and wind forces. Given our need to check these parameters independently of the availability of celestial radio sources, we looked into the possibility of using accelerometers on the antenna structure to establish its dynamical behaviour. The use of seismic accelerometers for performance characterization has been successfully demonstrated on optical telescopes \cite{Smith2004} and mm-antennas \cite{Ukita2002}, \cite{Ukita2004}. Using a set of 10 seismic accelerometers, installed on the antenna back-up structure (BUS), subreflector support structure (apex), and receiver cabin, we have measured accelerations allowing determination of rigid body motion of the elevation structure, and a few low-order distortions of the BUS. The nature of the accelerometers used in this work limits accurate displacement measurement to time scales of at most 10 seconds or frequencies of at least 0.1 Hz. Since this is well below the lowest eigenfrequencies of the antennas, this is sufficient to determine dynamic antenna behaviour. For the frequency range covered accurately by the accelerometers, approximately 0.1 to 30 Hz, it is possible to check the following antenna specifications: \begin{enumerate} \item Variations in surface shape, restricted to large scale effects like focal length and astigmatism \item Variations in pointing in elevation and cross-elevation direction \item Translation of apex structure with respect to the BUS \item Variations in path length along the boresight direction \end{enumerate} Using detailed long term wind studies, and wind measurements simultaneous with accelerometer measurements, antenna performance can be extrapolated to longer time scales under the assumption that wind effects dominate antenna performance on these time scales. Antenna performance should be met for all modes in which the antenna will be used to perform astronomical observations. For this paper, we have considered \begin{enumerate} \item pointing where the antenna is commanded to remain fixed at an azimuth and elevation position, \item sidereal tracking, where azimuth and elevation are updated continuously, \item fast switching mode, where the antenna is switched quickly between two neighbouring points, \item interferometric mosaicking, in which areas of sky are mapped at slow speed (0.05 deg/s), \item on-the-fly mapping (OTF), in which large areas of sky are mapped at high scan speed (0.5 deg/s). \end{enumerate} | \subsection{Accelerometers} Seismic accelerometers mounted on the back-up structures of large reflector antennas are capable of characterising all relevant BUS rigid body motions, and a few of its low-order deformation modes. Independently of any external sources to the antenna, and even for antennas without receivers, it is possible to derive the dynamical behaviour of the performance parameters of an antenna, such as pointing accuracy, primary reflector surface stability, and path length stability. The accuracy at which performance parameters can be measured is a function of frequency, and typically at the sub-$\mu$m and sub-arcsecond level for frequencies above 0.1 Hz. In combination with wind measurements, antenna performance for windy conditions can safely be extrapolated to frequencies well below the limit of 0.1 Hz imposed by noise in the accelerometer data. The validity of this extrapolation has been demonstrated with the use of an optical pointing telescope, not limited by the low-frequency noise. In addition to performance testing, the robust accelerometer system has proven to be well suited for troubleshooting of unexpected antenna behaviour, such as large and off-axis apex structure rotation for the AEC antenna, and servo-tuning issues for the VertexRSI antenna. \subsection{Wind-Driven Performance} Performance of the antennas for windy conditions was a major design driver. In spite of the very different wind conditions at the antenna test site and the ALMA site, it was possible to accurately predict performance for the ALMA site based on measurements performed at the ATF. The key to this extrapolation is careful characterisation of the wind at the sites, in particular at the ATF site. Wind-driven performance as measured on the antennas was combined with the wind characteristics during the time of measurement, which allowed calculation of wind-independent antenna properties, allowing calculation of antenna performance for any wind condition. As a spin-off of the investigations, antenna wake turbulence at a typical distance of 30 - 50 m was determined; which is useful information for calculation of antenna performance in the compact configuration. \subsection{Antenna Performance} Antenna performance as measured with the accelerometers complements other performance measurements, such as optical and radio pointing, and holography. The full performance of the antennas could not be determined with any of the individual methods, but a combination of them gives confidence in the completeness of the test results. During design, antenna performance was calculated from the sum of individual contributions to the total error budget. The performance numbers presented in this paper are best interpreted in the context of the corresponding error budget contributions used for antenna design. Table~\ref{tab:pointing} gives the measured performance for each antenna, and the calculated error budget entry as taken from the design documentation where available \cite{Mangum2006}. % \begin{table*} \centering \caption{Wind Pointing} \begin{tabular}{lll} & VertexRSI & AEC \\ \\ Pointing accuracy (wind only) & 0.81 arcsec& 0.45 arcsec\\ Pointing accuracy (wind only) error budget & 0.035 arcsec& 0.35 arcsec\\ Pointing accuracy (wind + tracking) & 0.94 arcsec& 0.50 arcsec\\ Pointing accuracy offset pointing requirement& 0.6 arcsec& 0.6 arcsec\\ \\ Primary reflector surface stability, wind effects (astigmatism + % AAM, at edge of BUS) & 5.3 + 2.2 $\mu$m& 6 + 5 $\mu$m\\ Primary reflector surface stability, wind effects error budget& 8.4 $\mu$m& 2.1 $\mu$m\\ Primary reflector surface stability, overall requirement& 25 $\mu$m& 25 $\mu$m\\ \\ Path length stability, wind effects & 6 $\mu$m& 6 $\mu$m\\ Path length stability, wind effects error budget& 7.6 $\mu$m& 3.5 $\mu$m\\ Path length stability, requirement total non-repeatable residual delay & 15 $\mu$m& 15 $\mu$m\\ \\ Structure flexure cross-elevation & 2.1 arcsec/(deg/s$^2$)& 1.6 arcsec/(deg/s$^2$)\\ Structure flexure elevation & 2.8 arcsec/(deg/s$^2$)& 2.1 arcsec/(deg/s$^2$)\\ \\ Lowest eigenfrequencies & 5.57 Hz & 6.8 Hz \\ \end{tabular} \label{tab:pointing} \end{table*} | 7 | 10 | 0710.4254 |
0710 | 0710.2191_arXiv.txt | We have studied non-axisymmetric standing accretion shock instability, or SASI, by 3D hydrodynamical simulations. This is an extention of our previous study on axisymmetric SASI. We have prepared a spherically symmetric and steady accretion flow through a standing shock wave onto a proto-neutron star, taking into account a realistic equation of state and neutrino heating and cooling. This unperturbed model is supposed to represent approximately the typical post-bounce phase of core-collapse supernovae. We then have added a small perturbation ($\sim$1\%) to the radial velocity and computed the ensuing evolutions. Not only axisymmetric but non-axisymmetric perturbations have been also imposed. We have applied mode analysis to the non-spherical deformation of the shock surface, using the spherical harmonics. We have found that (1) the growth rates of SASI are degenerate with respect to the azimuthal index $m$ of the spherical harmonics $Y_{l}^{m}$, just as expected for a spherically symmetric background, (2) nonlinear mode couplings produce only $m=0$ modes for the axisymmetric perturbations, whereas $m \neq 0$ modes are also generated in the non-axisymmetric cases according to the selection rule for the quadratic couplings, (3) the nonlinear saturation level of each mode is lower in general for 3D than for 2D because a larger number of modes are contributing to turbulence in 3D, (4) low $l$ modes are dominant in the nonlinear phase, (5) the equi-partition is nearly established among different $m$ modes in the nonlinear phase, (6) the spectra with respect to $l$ obey power laws with a slope slightly steeper for 3D, and (7) although these features are common to the models with and without a shock revival at the end of simulation, the dominance of low $l$ modes is more remarkable in the models with a shock revival. | Many efforts have been made for the multi-dimensional modeling of core-collapse supernovae (see \cite{tomas,kotake_rev} for reviews), urged by accumulated observational evidences that core-collapse supernovae are globally aspherical commonly \citep{wang96,wang01,wang02}. Various mechanisms to produce the asymmetry have been discussed so far: convection (e.g., \citet{herant_94,burrows_95,jankamueller96}), magnetic field and rapid rotation (see, e.g., \cite{kotake_rev} for collective references), standing (/stationary/spherical) accretion shock instability, or SASI, \citep{blondin_03,scheck_04,blondin_05,ohnishi_1,ohnishi_2,fog06}, and g-mode oscillations of protoneutron stars \citep{burr_new}. Note, however, that most of them have been investigated with two-dimensional (2D) simulations. Recently a 3D study on SASI was reported by \citet{blondin_nat}. In 2D, the shock deformation by SASI is described with the Legendre polynomials $P_{l}(\theta)$, or the spherical harmonics $Y_{l}^{m}(\theta, \phi)$ with $m=0$. Various numerical simulations have demonstrated unequivocally that the $l=1$ mode is dominant and a bipolar sloshing of the standing shock wave occurs with pulsational strong expansions and contractions along the symmetry axis \citep{blondin_03,scheck_04,blondin_05,ohnishi_1,ohnishi_2}. In 3D, on the other hand, \citet{blondin_nat} observed the dominance of a non-axisymmetric mode with $l =1$, $m=1$, which produces a single-armed spiral in the later nonlinear phase. They claimed that this "spiral SASI" generates a rotational flow in the accretion (see also \cite{blondin_shaw} for 2D computations in the equatorial plane) and that it may be an origin of pulsar spins. There are many questions on 3D SASI remaining to be answered yet, however. For example, we are interested in how the growth of SASI differs between 3D and 2D among other things. In particular, the change in the saturation properties should be made clear. This will be the focus of this paper. Another issue will be the generation of rotation in the accretion flow by SASI~\citep{blondin_nat}. Its efficiency and possible correlation with the net linear momentum should be studied more in detail and will be the subject of our forthcoming paper~\citep{iwa07}. Incidentally, it is noted that the neutrino heating and cooling were entirely ignored and the flow was assumed to be isentropic in \citet{blondin_nat}, but the implementation of these physics is helpful for considering the implications not only for the shock revival but also for the nucleosynthesis~\citep{kifo}. In this paper, we have performed 3D hydrodynamic simulations, employing a realistic equation of state \citep{shen98} and taking into account the heating and cooling of matter via neutrino emissions and absorptions on nucleons as done in our 2D studies \citep{ohnishi_1,ohnishi_2}. The inclusion of the neutrino heating allows us to discuss how the critical luminosity for SASI-triggered explosion could be changed in 3D from those in 2D. To answer the questions we raised above, we have varied the initial perturbations as well as the neutrino luminosity and compared the growth of SASI between 2D and 3D in detail, conducting a mode analysis not only for the linear phase but also for the nonlinear saturation phase. The plan of this paper is as follows: In Section \ref{sec2}, we describe the models and numerical methods. We show the main numerical results in Section \ref{sec3}. We conclude the paper in Section \ref{sec4}. | } In this paper we have studied the non-axisymmetric SASI by 3D hydrodynamical simulations, taking into account the realistic EOS and neutrino-heating and cooling. We have added various non-axisymmetric perturbations to spherically symmetric steady flows that accrete through a standing shock wave onto a proto-neutron star, being irradiated by neutrinos emitted from the proto-neutron star. The mode analysis has been done for the deformation of the shock surface by the spherical harmonics expansion. After confirming that our 3D code is able to reproduce for the axisymmetric perturbations the previous results on 2D SASI that we obtained in \citet{ohnishi_1}, we have done genuinely 3D simulations and found the followings. First, the model of the initial perturbation with $(l, |m|)=(1,0)$ and $(1,1)$ have demonstrated that the non-axisymmetric SASI proceeds much in the same manner as the axisymmetric SASI: the linear phase, in which the initial perturbation grows exponentially, lasts for about $100$ms and is followed by the nonlinear phase, in which various modes are produced by nonlinear mode couplings and their amplitudes are saturated. It has been found that the critical neutrino luminosity, above which the shock revival occurs, is not much different between 2D and 3D. For the neutrino luminosities lower than the critical value, the SASI is settled to a quasi-steady turbulence. We have found that the saturation level of each mode in the non-axisymmetric SASI is lower in general than that of the axisymmetric counterpart. This is mainly due to the fact that the number of the modes contributing to the turbulence is larger for the non-axisymmetric case. The sequence of the mode generation, on the other hand, strongly suggests that the nonlinear mode coupling is chiefly of quadratic nature. Second, the simulations with the random multi-mode perturbations being imposed initially have shown that the dynamics in the linear phase is essentially a superimposition of those of single-modes. Toward the end of the linear phase, high entropy blobs are generated continuously and grow, starting the nonlinear phase. We have observed that these blobs repeat expansions and contrations, being merged and splitted from time to time. In the non-explosion models, these nonlinear dynamics lead to the saturation of mode amplitudes and the quasi-steady turbulence. For the explosion models, on the other hand, the production of blobs proceeds much more quickly, followed by the oligarchic evolution, with a relatively small number of great blobs absorbing smaller ones, and as a result the shock radius increases almost monotonically. The spectral analysis has clearly demonstrated that low $l$ modes are predominant in the nonlinear phase just as in the axisymmetric case. We have also shown that the equi-partition is nearly established among different $m$ modes in the quasi-steady turbulence and that the spectrum summed over $m$ in the non-axisymmetric case is quite similar to the axisymmetric counterpart. This implies that the larger number of modes in the non-axisymmetric case is the main reason why the amplitude of each mode is smaller in 3D than in 2D. The power spectrum is approximated by the power law for $l \gtrsim 10$. Although the slope is slightly steeper for the non-axisymmetric models, whether the difference is significant or not is unknown at present. We have seen in the explosion models, on the other hand, that the oscillation period of each mode becomes longer in the late nonlinear phase as the shock radius gets larger and the advection-acoustic cycle becomes longer. What is more interesting is the fact that in this late phase the dominance of low $l$ modes becomes even more remarkable. Although this may be related with the volume-filling thermal convection, the mechanism is yet to be revealed. This spectral evolution leads to the global deformation of the shock surface to an ellipsoidal shape, whose major axis is not necessarily aligned with the symmetry axis, however. Finally, we have presented the neutrino heating in the 3D SASI. It has been seen that the volume-integrated heating rate is affected mainly in the nonlinear phase. The comparison with the spherically symmetric counterparts has confirmed that the SASI enhances the neutrino heating also in the non-axisymmetric case. Although the critical neutrino luminosity in the non-axisymmetric SASI is not much changed from that for the axisymmetric case, the spatial distribution of neutrino heating is different in the non-linear phase. Relatively narrow regions surrounding high entropy blobs are efficiently heated for the non-axisymmetric case, while wider regions near the symmetry axis are heated strongly in accord with the sloshing motion of shock wave along the symmetry axis for the axisymmetric case. For the non-explosion models, the high entropy blobs produced by neutrino heatings occupy the inside of shock wave, repeating expansion and contraction and being splitted and merged intermittently, but the flow is finally settled to a quasi-static state. For the explosion models, on the other hand, the high entropy blobs are generated much more quickly and extend the heating region, pushing the shock outwards and narrowing the cooling region near the proto-neutron star. In the present study we have not observed a persistent segregation of angular momentum in the accretion flow such as found by \citet{blondin_nat, blondin_shaw}) although instantaneous spiral flows are frequently seen. As discussed above, the equi-partition is nearly established among different $m$ modes in our models. It should be emphasized here again, however, that we have added the initial perturbations only to the radial velocity in this study and, as a result, the modes with $m=\pm1$ are equally contributing. We will defer the analysis of the models with the initial non-axisymmetric perturbations being added also to the azimuthal component of velocity to our forthcoming paper~\citep{iwa07}, in which we will also discuss a possible correlation between the kick velocity and spin of neutron stars if they are produced by the 3D SASI indeed. A lot of issues on SASI are still remaining to be studied. In particular, the influences of rotation and magnetic field are among the top priorities and will be addressed elsewhere in the near future. | 7 | 10 | 0710.2191 |
0710 | 0710.0477_arXiv.txt | We reexamined the gravitational time delay of light, allowing for various models of modified gravity. We clarify the dependence of the time delay (and induced frequency shift) on modified gravity models and investigate how to distinguish those models, when light propagates in static spherically symmetric spacetimes. Thus experiments by radio signal from spacecrafts at very different distances from Sun and future space-borne laser interferometric detectors could be a probe of modified gravity in the solar system. | 7 | 10 | 0710.0477 |
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0710 | 0710.5778_arXiv.txt | We validate the performance and accuracy of the current SEGUE (Sloan Extension for Galactic Understanding and Exploration) Stellar Parameter Pipeline (SSPP), which determines stellar atmospheric parameters (effective temperature, surface gravity, and metallicity) by comparing derived overall metallicities and radial velocities from selected likely members of three globular clusters (M~13, M~15, and M~2) and two open clusters (NGC~2420 and M~67) to the literature values. Spectroscopic and photometric data obtained during the course of the original Sloan Digital Sky Survey (SDSS-I) and its first extension (SDSS-II/SEGUE) are used to determine stellar radial velocities and atmospheric parameter estimates for stars in these clusters. Based on the scatter in the metallicities derived for the members of each cluster, we quantify the typical uncertainty of the SSPP values, $\sigma (\rm [Fe/H])$ = 0.13 dex for stars in the range of 4500 K $\le T_{\rm eff} \le 7500$ K and $2.0 \le \log g \le 5.0$, at least over the metallicity interval spanned by the clusters studied ($-2.3 \le {\rm [Fe/H]} < 0 $). The surface gravities and effective temperatures derived by the SSPP are also compared with those estimated from the comparison of the color-magnitude diagrams with stellar evolution models; we find satisfactory agreement. At present, the SSPP underestimates [Fe/H] for near-solar-metallicity stars, represented by members of M~67 in this study, by $\sim$ 0.3 dex. | The Sloan Extension for Galactic Understanding and Exploration (SEGUE) is one of three key projects (LEGACY, SUPERNOVA SURVEY, and SEGUE) in the current extension of the Sloan Digital Sky Survey, known collectively as SDSS-II. The SEGUE program is in the process of obtaining $ugriz$ imaging of some 3500 square degrees of sky outside of the SDSS-I footprint (Fukugita et al. 1996; Gunn et al. 1998, 2006; York et al. 2000; Stoughton et al. 2002; Abazajian et al. 2003, 2004, 2005; Pier et al. 2003), with special attention being given to scans of lower Galactic latitudes ($|b|$ $<$ 35$^{\circ}$) in order to better probe the disk/halo interface of the Milky Way. SEGUE is also obtaining $R$ $\simeq$ 2000 spectroscopy over the wavelength range 3800 $-$ 9200\,{\AA} for some 250,000 stars in 200 selected areas over the sky available from Apache Point, New Mexico. The SEGUE Stellar Parameter Pipeline (hereafter, SSPP) processes the wavelength- and flux-calibrated spectra generated by the standard SDSS spectroscopic reduction pipeline (Stoughton et al. 2002), obtains equivalent widths and/or line indices for 77 atomic or molecular absorption lines, and estimates $T_{\rm eff}$, log $g$, and [Fe/H] through the application of a number of approaches. The current techniques employed by the SSPP include a minimum distance method (Allende Prieto et al. 2006), neural network analysis (Bailer-Jones 2000; Willemsen et al. 2005; Re Fiorentin et al. 2007), auto-correlation analysis (Beers et al. 1999), and a variety of line index calculations based on previous calibrations with respect to known standard stars (Beers et al. 1999; Cenarro et al. 2001a,b; Morrison et al. 2003). The SSPP employs five different methods for estimation of $T_{\rm eff}$, eight for estimation of log $g$, and nine for estimation of [Fe/H]. Details of the methods used are discussed in detail by Lee et al. (2007a, hereafter Paper I). The use of multiple methods allows for empirical determinations of the internal errors for each parameter, based on the range of reported values -- typical internal errors for stars in the temperature range $4500$~K $\le$ $T_{\rm eff}$ $\le$ 7500~K are $\sim$ 73~K, $\sim$ 0.19 dex, and $\sim$ 0.10 dex, in $T_{\rm eff}$, log $g$, and [Fe/H], respectively. Allende Prieto et al. (2007, hereafter Paper III) point out that the internal uncertainties provided by the SSPP underestimate the typical random errors at high signal-to-noise ($S/N$) ratios because most methods in the SSPP make use of similar parameter indicators (e.g., hydrogen lines for effective temperature) and similar atmospheric models. Paper III empirically determines empirically external uncertainties of $\sim$ 130~K, $\sim$ 0.21 dex, and $\sim$ 0.11 dex, for $T_{\rm eff}$, log $g$, and [Fe/H], respectively, by comparison with high-resolution spectroscopy ($7000 < R < 45,000$) of brighter SDSS-I/SEGUE that have been obtained with 8m$-$10m class telescopes. Somewhat larger errors apply to stars with temperatures near the extremes of the range above. The present study of Galactic globular and open cluster stars tests the SSPP's ability to derive accurate results for stars with a wide range of temperatures and gravities appropriate for metal-poor and near-solar-metallicity stellar populations in the Galaxy, and demonstrates that the derived metallicity scale is identical for dwarfs and giants. Although the SSPP will continue to evolve in the near future, it has been frozen for now at the version used for obtaining results for stars with suitable data from SDSS Data Release 6 (DR-6; Adelman-McCarthy et al. 2007b). Previous versions of the SSPP have already been used for the analysis of SDSS-I observations. For example, Allende Prieto et al. (2006) report on the application of one of the methods included in the SSPP to some 20,000 F- and G-type stars from SDSS-I DR-3 (Abazajian et al. 2005). Beers et al. (2006) have compiled a list of over 6000 stars with [Fe/H] $< -2.0$ (including several hundred with [Fe/H] $< -3.0$), based on application of the present SSPP to some 200,000 stars from SDSS-I DR-5 (Adelman-McCarthy et al. 2007a). Carollo et al. (2007) reports on an analysis of the kinematics of relatively bright stars from SDSS-I that have been used as calibration objects during the main survey. In this paper, the second in the SSPP series, we show that estimates of the atmospheric parameters and radial velocities obtained by the SSPP for stars with a reasonable likelihood of membership in previously studied Galactic globular and open clusters are sufficiently accurate to justify the use of the present SSPP parameters for carrying out detailed studies of the halo and thick-disk populations of the Milky Way. In deriving the overall iron abundance for each cluster, we assume it comprises a chemically homogenous population. In \S 2, the photometric and spectroscopic data obtained for M~13, M~15, M~2, NGC~2402, and M~67 are described. Section 3 presents the methods used to separate likely cluster members from field stars in the directions toward these clusters. Best estimates of the overall [Fe/H] and radial velocity of each cluster are derived in \S 4. In \S 5 we compare the SSPP determinations of $T_{\rm eff}$ and log $g$ for selected member stars in each cluster with their expected positions on color-magnitude diagrams. A summary and brief conclusions are provided in \S 6. | Based on photometric and spectroscopic data reported in SDSS-I and SDSS-II/SEGUE, we have examined estimates of stellar atmospheric parameters and heliocentric radial velocities obtained by the SEGUE Stellar Parameter Pipeline (SSPP) for likely members of three Galactic globular clusters, M~13, M~15, and M~2, and two open clusters, NGC~2420 and M~67, and compared them with those obtained by external estimates for each cluster as a whole. From the derived scatters in the metallicities and radial velocities obtained for the likely members of each cluster, we quantify the typical external uncertainties of the SSPP-determined values, $\sigma (\rm [Fe/H])$ = 0.13 dex, and $\sigma (\rm RV)$ = 7.7 km~s$^{-1}$, respectively. These uncertainties apply for stars in the range of 4500 K $\le T_{\rm eff} \le 7500$ K and $2.0 \le \log g \le 5.0$, at least over the metallicity interval spanned by the clusters studied ($-2.3 \le {\rm [Fe/H]} < -0.4$). Therefore, the metallicities and radial velocities obtained by the SSPP appear sufficiently accurate to be used for studies of the kinematics and chemistry of the metal-poor and moderately metal-rich stellar populations in the Galaxy. We have also confirmed that $T_{\rm eff}$ and log $g$ are sufficiently well-determined by the SSPP to distinguish between different luminosity classes through a comparison with theoretical predictions. A comparison of the analysis of the available high-resolution spectroscopy of SDSS-I/SEGUE stars (Paper III) with the SSPP predictions indicates that the uncertainty in radial velocities adopted by the SSPP is no more than 5 km~s$^{-1}$ (after adjusting for an empirical offset of +7.3 km~s$^{-1}$). The empirically determined precisions in estimated atmospheric parameters are $\sim$ 130~K for effective temperature, $\sim$ 0.21 dex for surface gravity, and $\sim$ 0.11 dex for [Fe/H]. These errors apply to the brightest stars obtained by SDSS-I/SEGUE observations, on the order of $14.0 \le g \le 15.5$, and are expected to degrade somewhat for fainter stars. We also found that the SSPP tends to underestimate [Fe/H] for near-solar-metallicity stars (represented by members of M~67 in this study), by $\sim$ 0.3 dex. In future papers we will compare the predictions of the SSPP with intermediate-metallicity clusters ([Fe/H] $\sim$ $-$0.7) and with additional near-solar-metallicity populations, as sampled by metal-rich globular clusters and nearby open clusters. Additionall metal-poor clusters will also be examined. Further refinements in the SSPP, which hopefully will be better able to recover accurate abundances for near-solar-metallicity stars, are anticipated. | 7 | 10 | 0710.5778 |
0710 | 0710.2702_arXiv.txt | The effect of viscosity and of converging flows on the formation of blobs in the slow solar wind is analysed by means of resistive MHD simulations. The regions above coronal streamers where blobs are formed \citep{sheeley} are simulated using a model previously proposed by \citet{einaudijgr}. The result of our investigation is two-fold. First, we demonstrate a new mechanism for enhanced momentum transfer between a forming blob and the fast solar wind surrounding it. The effect is caused by the longer range of the electric field caused by the tearing instability forming the blob. The electric field reaches into the fast solar wind and interacts with it, causing a viscous drag that is global in nature rather than local across fluid layers as it is the case in normal uncharged fluids (like water). Second, the presence of a magnetic cusp at the tip of a coronal helmet streamer causes a converging of the flows on the two sides of the streamer and a direct push of the forming island by the fast solar wind, resulting in a more efficient momentum exchange. | The understanding of slow solar wind genesis has progressed considerably in the last few years following the observational discovery of plasma inhomogeneities (called {\em blobs}) formed and expelled from regions above coronal streamers~\citep{sheeley, wang}. The observational evidence has been provided by the Large Angle and Spectrometric Coronagraph (LASCO) instrument on the Solar and Heliospheric Observatory (SOHO). Approximately 4 blobs per day were observed in a relatively quiet period for the solar corona (year 1997, near solar minimum). After formation the blobs are carried away with the ambient slow solar wind. Generally, the plasma blobs move radially outward with their speed increasing from 0-250 km s$^{-1}$ in the region 2-6 ${\rm R}_\odot$ (C2 field of view) to 200-450 km s$^{-1}$ in the region 3.7-30 ${\rm R}_\odot$ (outer portion of the C3 field of view). Models of the formation of blobs have focused on mechanisms involving reconnection between open field lines and closed field lines \citep{einaudijgr, einaudi, endeve,fisk, vanaalst, wang, wu}. Through field line reconnection, plasma from the denser regions within the helmet streamer is liberated and becomes tied to open field lines and forms the slow solar wind \citep{einaudi}. The present work is particularly centred around the model proposed by \citet{einaudijgr}. In the region downstream of the cusp in the streamer belt, \citet{einaudijgr} approximate the configuration as a 1D reversed field configuration with a 1D wake velocity profile. Under these assumptions, they can reproduce both the acceleration of the slow solar wind and the formation of blobs, formulating a complete picture of the genesis of the slow solar wind, in accordance with the observations \citep{einaudijgr}. The scenario has recently been further extended~\citep{rappazzo} to include the so called {\em melon seed effect} due to the diamagnetic force caused by the overall magnetic field radial gradients in the spherical geometry of the Sun~\citep{schmidt} and to include the effects of the cusp magnetic configurations atop the closed field lines in coronal streamers~\citep{lapenta-wind}. The fundamental conclusion of the model by \citet{einaudijgr} and in successive refinements is that a blob is formed by the tearing instability and that by virtue of the interaction with the fast solar wind the blob is accelerated and ejected, carrying with it the plasma that forms the slow solar wind. The fundamental question posed by the present work is: ``what is the basic physical process for the plasma acceleration''? We consider specifically two possibilities. First, that the momentum transfer causing the blob acceleration is due to viscous drag. Second, that the acceleration is due to the insurgence of the tearing instability and its non-linear effects. The acute reader will not find the suspense spoiled by learning that, as usual, Nature turns out not to be black or white. We discover that both mechanisms are present in a interplay mediated by the electric field. Furthermore, the presence of the converging flows at the cusp of the coronal streamer causes a direct momentum transfer by the action of the plasma flowing against the forming island and pushing it upward from the solar surface. The paper is organised as follows. Section 2 summarises the initial configurations considered distinguishing the 1D initial configuration proposed by \citet{einaudijgr} and its 2D extension to include the presence of a magnetic cusp~\citep{lapenta-wind}. Section 3 very briefly provides the methodology for our investigation: the resistive MHD model as implemented in the FLIP3D-MHD code~\citep{brackbill}. Section 4 and 5 present the non-linear evolution, respectively, of the case of the initial 1D reversed field and of the 2D configuration with a magnetic cusp. Conclusions are reached in Section 6. | 7 | 10 | 0710.2702 |
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0710 | 0710.1584_arXiv.txt | Recent observations with atmospheric Cherenkov telescope systems such as \HESS\ and MAGIC have revealed a large number of new sources of very-high-energy (VHE) $\gamma$-rays from 100~GeV -- 100~TeV, mostly concentrated along the Galactic plane. At lower energies (100 MeV -- 10 GeV) the satellite-based instrument EGRET revealed a population of sources clustering along the Galactic Plane. Given their adjacent energy bands a systematic correlation study between the two source catalogues seems appropriate. Here, the populations of Galactic sources in both energy domains are characterised on observational as well as on phenomenological grounds. Surprisingly few common sources are found in terms of positional coincidence and spectral consistency. These common sources and their potential counterparts and emission mechanisms will be discussed in detail. In cases of detection only in one energy band, for the first time consistent upper limits in the other energy band have been derived. The EGRET upper limits are rather unconstraining due to the sensitivity mismatch to current VHE instruments. The VHE upper limits put strong constraints on simple power-law extrapolation of several of the EGRET spectra and thus strongly suggest cutoffs in the unexplored energy range from 10~GeV -- 100~GeV. Physical reasons for the existence of cutoffs and for differences in the source population at GeV and TeV energies will be discussed. Finally, predictions will be derived for common GeV--TeV sources for the upcoming GLAST mission bridging for the first time the energy gap between current GeV and TeV instruments. | In recent years the knowledge of the Galactic VHE $\gamma$-ray sky above 100~GeV has been greatly improved through the detection and subsequent study of many sources, mostly by means of ground-based Imaging Atmospheric Cherenkov telescope systems such as the High Energy Stereoscopic System (\HESS) or the Major Atmospheric Gamma-ray Imaging Cherenkov Observatory (MAGIC). Currently known Galactic VHE $\gamma$-ray emitters include shell-type Supernova remnants (SNRs)~\citep{HESS1713_II, HESS1713_III, HESSVelaJr}, Pulsar Wind Nebulae (PWNe)~\citep{HESSMSH, HESS1825II, HESSKooka}, $\gamma$-ray binaries~\citep{HESSLS5039II, MAGICLSI}, Molecular clouds~\citep{HESSGCDiffuse} and possibly also clusters of massive stars~\citep{HESSWesterlund}. These various source classes were discovered both in pointed observations using \HESS\ and MAGIC as well as in a systematic survey of the inner Galaxy performed with the \HESS\ instrument. The highest energy photons detected from these source classes reach $\sim 100$~TeV~\citep{HESS1713_III}, currently representing the end of the observable electromagnetic spectrum for astrophysical objects. It is natural to investigate the relationship of these TeV sources to sources at lower energies as will be done in this work focusing on Galactic sources. The closest energy band for which data exist is that studied by the Energetic Gamma Ray Experiment Telescope (EGRET) aboard the Compton Gamma-Ray Observatory with an energetic coverage from 100~MeV -- 10~GeV~\citep{EGRET}. The GeV sky has a distinctively different overall appearance compared to TeV energies. In particular focusing on our Galaxy, the most prominent feature of the GeV sky is the dominant diffuse emission from cosmic ray (CR) interactions in the Galaxy, while the TeV sky due to the steeply falling energy spectrum of the diffuse component is dominated by individual sources. However, several prominent $\gamma$-ray sources are known to emit at both GeV and at TeV energies, the Crab Nebula being the most prominent example~\citep{WhippleCrab,EGRETCrab, HEGRACrab, HESSCrab, MAGICCrab}. In this paper the relationship between Galactic EGRET and VHE $\gamma$-ray sources will be assessed in a systematic way. For cases with a positional coincidence between a VHE and an EGRET source (in the following called ``coincident sources'') all currently known Galactic objects will be considered. For cases in which a source is detected only in one band -- the ``non-coincident sources'' -- we focus on the region covered by the \HESS\ Galactic plane survey (GPS) during 2004 and 2005~\citep{HESSScan, HESSScanII} (Galactic longitude $\pm 30^{\circ}$, Galactic latitude $\pm 3^{\circ}$) so that a statistical assessment of the ``non-connection'' can be made. EGRET was unable to perform detailed studies of the $\gamma$-ray sky above 10~GeV, partly due to back-splash of secondary particles produced by high-energy $\gamma$-rays causing a self-veto in the monolithic anti-coincidence detector used to reject charged particles. The upcoming Gamma Ray Large Area Space Telescope (GLAST) Large Area Telescope (LAT) will not be strongly affected by this effect since the anti-coincidence shield was designed in a segmented fashion~\citep{GLASTACD}. Moreover, the effective area of the GLAST-LAT will be roughly an order of magnitude larger then that of EGRET. The GLAST-LAT mission will therefore for the first time fully bridge the gap between the energy range of EGRET and current VHE instruments. Part of the study presented here can be seen as preparatory work for GLAST-LAT studies of sources in the largely unexplored energy band between 10 and 100~GeV. \begin{figure}[ht] \begin{center} \includegraphics[width=0.9\textwidth]{f1.eps} \end{center} \caption{{\bf{Left:}} Integral sensitivities for current, past and future $\gamma$-ray instruments (5-$\sigma$ sensitivity for $E>E_0$ multiplied with $E_0$ assuming a spectrum of $E^{-2}$). The solid lines show the nominal instrument sensitivities (for a typical observation time as specified below), the dashed curves show the actual sensitivities for the Inner Galaxy as appropriate for this work. INTEGRAL's (IBIS/ISGRI) sensitivity curve (solid green) shows the sensitivity for an observation time of $10^5$s, a typical value in the Inner Galaxy. The EGRET curves (brown) are shown for the whole lifetime of the mission (periods 1--9) for the Galactic anti-centre (solid) which received the largest exposure time and has a lower level of diffuse $\gamma$-ray emission than the Inner Galaxy and for the position of RX\,J1713.7--3946 (dashed), a typical position in the Inner Galaxy dominated by diffuse $\gamma$-ray background emission. The GLAST curves in red (taken from http://www-glast.slac.stanford.edu/software/IS/glast\_lat\_performance.html) show the 1-year sky-survey sensitivity for the Galactic North pole -- again a position with low diffuse emission (solid), and for the position of RX\,J1713.7--3946 (dashed). The H.E.S.S.\ curves (blue) are shown for a 50-hour pointed observation of a point-like source (solid) and for a 5-hour observation of a somewhat extended source as is typical for the Galactic Plane survey (angular cut of $\sqrt{0.05}$). The MAGIC curve (light blue) represents a 50-hour observation of a point-source. {\bf{Right:}} Energy-dependence of the angular resolution for selected $\gamma$-ray instruments expressed by the 68\%-containment radius of the point-spread function (PSF). As can be seen, the angular resolution of GLAST becomes comparable with current VHE instruments at high energies, whilst at the lower energy end GLAST and EGRET have comparable resolutions.}\label{fig::Sensitivity} \end{figure} From the 2004 and 2005 \HESS\ GPS 22 VHE $\gamma$-ray sources have been reported in the Inner Galaxy. The third EGRET catalogue~\citep{EGRET} represents the companion to the VHE source catalogue above an energy threshold of 100 MeV (with peak sensitivity between 150 and 400 MeV, depending on the $\gamma$-ray source spectrum). It lists 271 sources, 17 of which are located within the \HESS\ GPS region. Whilst the EGRET range currently represents the nearest energy band to VHE $\gamma$-rays, for the very few EGRET sources detected all the way up to $\sim 10$~GeV, there is still an unexplored energy band of roughly one decade before the VHE $\gamma$-ray energy range begins at $\sim 100$~GeV (it should be noted that EGRET does have some sensitivity beyond 10~GeV: \citet{EGRET10GeV} reported the detection of $\sim$ 1500 photons above that energy with 187 of these photons being found within $1^{\circ}$ of a source listed in the third EGRET catalogue). Comparing the instrumental parameters of VHE instruments and EGRET there is a clear mismatch both in angular resolution and in sensitivity as can be seen in Figure~\ref{fig::Sensitivity}. In a $\sim 5$ hour observation (as a typical value in the GPS region) \HESS\ is a factor of $\sim 50-80$ more sensitive (in terms of energy flux $E^2 dN/dE$) than EGRET above 1~GeV in the Galactic Plane for the exposure accumulated between 1991 and 1995 (corresponding to the third EGRET catalogue). Assuming a similar energy flux output in the two different bands this mismatch implies at first sight that \HESS\ sources are not likely to be visible in the EGRET data set. Conversely (again under the assumption of equal energy flux output), VHE $\gamma$-ray instruments should be able to detect the majority of the EGRET sources, as has been suggested in the past. Figure~\ref{fig::EnergyFluxDistribution} compares the energy fluxes $\nu F \nu$ for EGRET sources and \HESS\ sources in the inner Galaxy. Clearly, the EGRET sources do not reach down as low in energy flux as the \HESS\ sources, a picture that will change once the GLAST-LAT is in orbit as depicted by the GLAST-LAT sensitivity (dashed line). In reality the \naive\ expectation of equal energy flux output in the GeV and TeV band can easily be wrong in Galactic $\gamma$-ray sources for various reasons: EGRET sources may not emit comparable energy fluxes in the VHE $\gamma$-ray band but rather exhibit cut-offs or spectral breaks in the energy band between EGRET and \HESS\ \citep[this is certainly the case for pulsed magnetospheric emission from pulsars, see for example][]{HESSPulsar}. Furthermore, \HESS-like instruments are typically only sensitive to emission on scales smaller than $\sim 1^{\circ}$. If any of the EGRET sources are extended beyond $1^{\circ}$ without significant sub-structure on smaller scales (not precluded given the poor angular resolution of EGRET), current Imaging Cherenkov instruments may not be able to detect them since these sources would completely fill the field of view (FoV) and be removed by typical background subtraction methods ~\citep[see for example][]{BergeBackground}. Given the upcoming launch of GLAST and the recent \HESS\ survey it seems timely to study the relationship between GeV and TeV emitting sources in more detail. \begin{figure}[ht] \begin{center} \includegraphics[width=0.6\textwidth]{f2.eps} \end{center} \caption{Distribution of integrated energy flux $\nu F \nu$ for sources in the Inner Galaxy discussed here. For EGRET the energy flux between 1~GeV and 10~GeV, for the \HESS\ sources, the energy flux between 1~TeV and 10~TeV is shown. Also shown is the sensitivity prediction for the GLAST-LAT for a typical location in the Inner Galaxy (l=10, b=0).}\label{fig::EnergyFluxDistribution} \end{figure} Section~\ref{sec::analysis} describes the data and analysis methods used in this study, section~\ref{sec::connect} describes the sources detected in both energy bands, and section~\ref{sec::nonconnect} focuses on sources detected in only one of the two energy regimes. In section~\ref{sec::interpretation} astrophysical implications of the study are discussed. | The main results of the study of the relationship between GeV and TeV sources are: \begin{enumerate} \item There are rather few spatially coincident GeV-TeV sources for the considered Galactic region. \item Those few positional coincident GeV-TeV sources could occur by chance, the chance probability of detecting two coincident sources within the \HESS\ GPS region is $\sim 40$\%, thus no strong hint for a common GeV/TeV source population is detected. \item Spectral compatibility (based on a power-law extrapolation) seems present for most of the positionally coincident sources, but again, this is expected to occur by chance (as described in the text) given the sensitivity mismatch and the different energy bands. \item Dedicated \HESS\ limits at the position of the EGRET sources are constraining for a power-law extrapolation from the GeV to the TeV range for several of the EGRET sources, strongly suggesting cutoffs in the energy spectra of these EGRET sources in the unexplored region below 100~GeV. Power-law extrapolation of EGRET spectra seem to be ruled out for most of the EGRET sources investigated in this study. \item Dedicated EGRET limits at the position of the \HESS\ sources are not constraining for a power-law extrapolation from the TeV to the GeV range. This picture will dramatically change once the GLAST-LAT with its improved sensitivity over EGRET is in orbit. \item Several important mechanisms for cutoffs in the energy spectra of GeV sources have been discussed. There are well motivated physical reasons why the population of GeV and of TeV sources might be distinct. \item If a source can be detected with both GeV and TeV instruments, the huge energy ``lever arm'' over 5-6 decades in energy will undoubtedly provide stringent constraints on the $\gamma$-ray emission mechanism in these Galactic particle accelerators. \end{enumerate} Summarising, the study presented here shows that the GLAST-LAT will tremendously advance the study of the relationship between GeV and TeV sources by improving the sensitivity over EGRET by an order of magnitude and in particular by bridging the currently uncovered energy range between 10~GeV and 100~GeV. | 7 | 10 | 0710.1584 |
0710 | 0710.1067_arXiv.txt | % We present an analysis of the magnetic-field fluctuations in the magnetoionic medium in front of the radio galaxy 3C 31 derived from rotation-measure (RM) fits to high-resolution polarization images. We first show that the Faraday rotation must be due primarily to a foreground medium. We determine the RM structure functions for different parts of the source and infer that the simplest form for the power spectrum is a power law with a high-frequency cutoff. We also present three-dimensional simulations of RM produced by a tangled magnetic field in the hot plasma surrounding 3C 31, and show that the observed RM distribution is consistent with a spherical plasma distribution in which the radio source has produced a cavity. | Our knowledge of the structure and origin of magnetic fields in elliptical galaxies, groups and clusters is still rudimentary, but Faraday rotation of linearly-polarized radio emission can be used to probe the fields in ionized foreground gas. Here, we present an analysis of the magnetic-field fluctuations in the magnetoionic medium in front of the FR\,I radio galaxy 3C\,31 (z = 0.0169) derived from rotation-measure (RM) fits to high-resolution polarization images. Our analysis is based on VLA observations at 6 frequencies in the range 1.4 -- 8.4\,GHz with resolutions of 5.5 and 1.5\,arcsec FWHM (Figs~\ref{fig:rmlo} and \ref{fig:rmhi}). We show images of normalized polarization gradient $p^\prime(0)/p(0)$ from fits to $p(\lambda^2) = p(0) + p^\prime(0)\lambda^2$, where $p(\lambda^2)$ is the degree of polarization at wavelength $\lambda$ and a prime denotes differentiation with respect to $\lambda^2$, together with RM images derived from fits to $\chi(\lambda^2) = \chi(0) + {\rm RM}\lambda^2$ at 4 -- 6 wavelengths, where $\chi$ is the {\bf E}-vector position angle. The residual depolarization at 1.5\,arcsec resolution is very small and the rotation is accurately proportional to $\lambda^2$, indicating almost completely resolved foreground rotation. There is a large asymmetry across the nucleus: the lobe with the brighter jet shows a much smaller RM fluctuation amplitude than the counter-jet lobe on all scales, qualitatively as expected from relativistic jet models \citep{L88}. | 7 | 10 | 0710.1067 |
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0710 | 0710.3062_arXiv.txt | We systematically study the effects of collisions on the overall dynamical evolution of dense star clusters using Monte Carlo simulations over many relaxation times. We derive many observable properties of these clusters, including their core radii and the radial distribution of collision products. We also study different aspects of collisions in a cluster taking into account the shorter lifetimes of more massive stars, which has not been studied in detail before. Depending on the lifetimes of the significantly more massive collision products, observable properties of the cluster can be modified qualitatively; for example, even without binaries, core collapse can sometimes be avoided simply because of stellar collisions. | In dense stellar systems like massive young clusters, galactic centers, and old globular clusters (GC), the high densities of stars can give rise to many direct, single--single physical collisions, in addition to collisions mediated by dynamical interactions of binaries (\cite{2004MNRAS.352....1F}). In these systems stellar evolution can also be significantly modified through physical collisions. For example, in young star clusters, collisions can give rise to runaway growth of merger products, producing many exotic stellar populations, such as intermediate-mass black holes (\cite{2006ApJ...640L..39G}; \cite{2006astro.ph.12040T}). The cores of typical old Galactic GCs can also attain high enough densities so that most core stars undergo collisions during their lifetime (\cite{1976ApL....17...87H}; \cite{2004MNRAS.352....1F}). These collisions not only change the evolution of individual stars, but the increased stellar masses also change the overall GC properties like the core radius ($r_c$) and the half-mass radius ($r_h$). Moreover, at least some of the observed exotic stellar populations in dense clusters, like blue straggler stars (BSS) (\cite{2007ApJ...663..277F}; \cite{2001ApJ...548..323S}; \cite{1997ApJ...487..290S}), and compact binaries like ultracompact X-ray binaries (UCXB), are likely created through collisions (\cite{2006ApJ...640..441L}). A dense cluster of stars naturally evolves towards eventual core collapse through relaxation. At their present ages, most Galactic GCs are expected to have collapsed cores. However, observations show that the measured values of $r_c/r_h$ for most Galactic GCs are higher than predicted by theoretical models (\cite{1996ApJ...458..178V}). Many scenarios have been proposed to explain this apparent discrepancy. For example, the core can be supported against deep collapse by dynamically extracting the binding energy of hard primordial binaries (\cite{2007MNRAS.374..857T}; \cite{2007ApJ...658.1047F}). However, it is hard to explain most of the high observed values of $r_c/r_h$ in Galactic GCs purely through this ``binary burning'' process. Other mechanisms to halt core collapse like ejection of stellar mass BHs from the core (\cite{2007MNRAS.379L..40M}) or stellar captures by a central intermediate-mass black hole (\cite{2007MNRAS.374..857T}; \cite{2006astro.ph.12040T}) have also been discussed at this meeting (see contributions by Mackey and Trenti in this volume). It has also been suggested that $r_c/r_h$ could simply keep increasing over long timescales in clusters having fairly long relaxation times via ongoing mass segregation (\cite{2004ApJ...608L..25M}). Here we study stellar collisions as a possible mechanism for supporting clusters against core collapse. In a regime where collisions are important, they can produce many stars significantly more massive than those in the background population. Thus the subsequent evolution of collision products will be much faster than for normal stars in the cluster. Although the stellar evolution and observable properties of collision products have been extensively studied before (\cite{2001ApJ...548..323S}; \cite{1997ApJ...487..290S}), the feedback effects of collisions on the overall dynamical evolution of GCs has received less attention. The shorter evolution timescales of massive collision products may support the core against collapse even without any primordial binaries or other mechanisms for energy production, simply via indirect heating through stellar evolution mass loss (\cite{1991ApJ...378..637G}; \cite{1989Natur.339...40G}; \cite{1987ApJ...319..801L}). Using $N$-body simulations, we have begun studying numerically how collisions and the subsequent evolution of collision products can alter the overall properties of GCs. We use the Northwestern group's H\'enon-type Monte-carlo code, which provides a detailed, star-by-star representation of clusters with up to $N\sim 10^6-10^7$ stars (\cite{2007ApJ...658.1047F}; \cite{2001ApJ...550..691J}; \cite{2000ApJ...540..969J}). In \S2 we will first present two simple limiting cases, bracketing reality and illustrating the dramatic changes in global cluster properties depending on how the evolution of collision products is treated in the models. We also present the evolution of a more realistic GC model with a conventional ``rejuvenation'' prescription for determining the lifetimes of collision products. We discuss the implications of our study and planned future work in \S3. \begin{table}\def~{\hphantom{0}} \begin{center} \caption{Initial Conditions of Simulated GCs} \label{tab:ic} \begin{tabular}{cccccccc}\hline & Model & IMF & N & $n_{binary}$ & $r_c$ (pc) & Virial radius (pc) & $\rho_c$ ($M_{\odot}/pc^3$) \\ \hline% Cases 1,2 & Plummer & Single-Mass & $10^6$ & $0$ & $0.3$ & $0.85$ & $10^6$ \\ Case 3 & King, $w_0 = 6$ & Kroupa ($0.1$ -- $2.0 M_{\odot}$) & $10^6$ & $0$ & $0.87$ & $2.89$ & $1.7\times 10^4$ \\ \hline \end{tabular} \end{center} \end{table} | Although determination of stellar evolution after a collision is a subject of continuing research, the effects of collisions on cluster dynamics have not been studied in detail previously. This is a first report of an ongoing systematic study on this effect. Using typical initial conditions for old GCs and simple assumptions for rejuvenation we show that collisions between single stars not only alter the stellar properties and produce exotic stellar populations like some of the BSS, they also affect the overall GC properties. Most importantly, collisions can support the core of a cluster against collapse in typical old Galactic GCs. Even with our simple but reasonable assumptions, the values of $r_c/r_h$ obtained for our simulated GCs compare well with the observed $r_c/r_h$ values for Galactic GCs. Furthermore, we obtain a population of BSS candidates contained within the core, also consistent with observations (\cite{2007ApJ...661..210L}). Note that we do not find any BSSs well outside the core, consistent with the current understanding that those BSSs are most likely formed via primordial binary mergers (\cite{2006MNRAS.373..361M}). We have adopted many extreme simplifications for this first look at the problem. For example, at this stage of our study, once a star evolves off the MS, it is removed from the simulation, leaving no remnant. Since remnants are normally only a few percent of the total progenitor mass, we expect that this approximation will not affect the overall GC properties significantly, so far as the increase in energy is concerned from mass loss. However, some remnants from very massive stars at a very young age can remain in the cluster and sink into the core through dynamical friction. Dynamical interactions, including collisions, of these massive remnants can also alter the core properties of the GCs in certain regimes (\cite{2007MNRAS.379L..40M}; \cite{2007MNRAS.374..857T}; \cite{2006astro.ph.12040T}). Another possibly important effect left out of this study for now is the role of primordial binaries (\cite{2007ApJ...658.1047F}; \cite{2004MNRAS.352....1F}). On the one hand, the presence of primordial binaries will increase the rate of collisions through resonant encounters (\cite{2004MNRAS.352....1F}), on the other hand, binaries will provide further support of a cluster against collapse and hence may prevent the core from reaching high enough densities for significant collisions. \begin{center} \begin{figure} \includegraphics[height=2.7in,width=2.7in,angle=0]{chatterjee_f3.eps} \includegraphics[height=2.7in,width=2.7in,angle=0]{chatterjee_f4.eps} \caption{ {\em a)} Position vs mass scatter plot for all collision products still at their MS life at $14$ Gyr. Two times the present age turn-off mass ($1.6 M_{\odot}$) is shown as a horizontal solid line to guide the eye. The position of the core is shown with the vertical dotted line. {\em b)} Histogram showing the positions of the same population as in $(a)$. {\em c)} Histogram showing the positions of the collision products still in their MS life at present age having masses $\geqslant 1.6 M_{\odot}$. } \label{fig:king_BS} \end{figure} \end{center} | 7 | 10 | 0710.3062 |
0710 | 0710.2855_arXiv.txt | {One of the ways to determine the contribution of the dark halo to the gravitational potential of a galaxy is the study of non-circular (streaming) motions and the associated gas shocks in the bar region. These motions, determined by the potential in the inner parts, can break the disk-halo degeneracy. Here, two main fluid dynamical approaches have been chosen to model the non-circular motions in the bar region; a 2-D Eulerian grid code for an isothermal gas (FS2) and a 3-D smoothed particle hydrodynamic code (N-body/SPH)} % {The aim of this paper is to compare and quantify the differences of the gas flows in rotating barred potential obtained using two different fluid dynamical approaches. We analyse the effect of using 2-D and a 3-D codes in the calculation of gas flow in barred galaxies and to which extend the results are affected by the code. To do this, we derive the velocity field and density maps for the mass model of NGC~4123 using a 3-D N-body/SPH code and compare the results to the previous 2-D Eulerian grid code results.} {Numerical modelling, 3-D N-body/SPH simulations} {The global velocity field and the gas distribution is very similar in both models. The study shows that the position and strength of the shocks developed in the SPH simulations do not vary significantly compared to the results derived from the 2-D FS2 code. The largest velocity difference across the shock is 20\kms between the 2-D and 3-D fluid dynamical models.} {The results obtained in the studies deriving the dark matter content of barred galaxies using the bar streaming motions and strength and position of shocks are robust to the fluid dynamical model used. The effect of 2-D and 3-D modelling can be neglected in this type of studies.} | There have been several studies addressing the distribution of dark matter in galaxies using the non-axisymmetry of the potential of barred galaxies: a) the fluid dynamics approach with the work by~\citet[][hereafter WSW01]{weiner01} for the modelling of NGC~4123; ~\cite{perez3} and \cite[][hereafter PFF04]{perez4} for the study of the dark matter content of 5 barred galaxies (NGC~5505, NGC~7483, NGC~5728, and NGC~7267); and b) the sticky-particle approach by~\cite{rautiainen} modelling the dynamics of ESO~566-24, and~\cite{salo} who modelled the H$\alpha$ velocity field of IC~4214 to derive the halo contribution. The streaming motions and the associated gas shocks in the bar region are determined by the potential in the inner parts of galaxies. The velocity field of an axisymmetric galaxy does not allow to uniquely disentangle the contribution of the halo from that of the disk.The signatures of non-circular motions, however, can break the disk-halo degeneracy and be used to obtain the contribution of the dark halo to the potential. Common to all the modelling, the stellar potential is generated directly from the broad band galaxy images with some recipe to treat the vertical distribution and some assumptions about the mass-to-light (M/L) of the stellar populations. The velocity fields obtained in this way are then compared to the observed kinematic information to determine if they can be reproduced by the models. The main result, common to all the modelling carried out up to now, is the fact that the gravitational field in the inner region is mostly provided by the stellar luminous component. The bar pattern speeds found by the different groups are consistent with fast rotators, and the best fitting M/L ratios obtained are compatible with M/L ratios derived from current population synthesis models. Although it is a very powerful method, only a handful of galaxies have been modelled in this way. Two main fluid dynamical approaches have been chosen for this type of modelling: an Eulerian grid code for an isothermal gas (FS2) and a smoothed particle hydrodynamic code (N-body/SPH). SPH is a Lagrangian method for solving Euler's equation of motion. In this technique, the system is represented by a set of particles and the gas properties are calculated by a weighted average over the neighbouring particles. On the other hand, the Eulerian grid is fixed in space and time, the grid nodes and cells remain spatially fixed while the materials flow through the mesh. In the modelling carried out by WSW01, the scale height is introduced as a smoothing length for the potential while the approach chosen by PFF04 uses a full 3-D calculation of the forces. For future work and as a consistency check, it is important to understand whether the results obtained are dependent a) on the dynamical code used and b) 2D versus 3-D. No significant differences are expected between the SPH and the Eulerian grid approach. Previous simulations~\cite[][hereafter PA00]{englmaier,patsis00}, running 2-D SPH under the same conditions as used by~\cite{athanassoula92} with the FS2 Eulerian grid code indicated that the main features, such as the strength of the shocks in the bar are quite similar and small differences arise due to the statistical nature of SPH or to numerical and artificial viscosities used in the different codes. Both studies pointed out the importance of certain parameters used in the simulations for the outcome of the modelling, such as the sound speed and the way the non-axisymmetric part of the potential is introduced. The use of a 2-D versus a 3-D code might introduce large variations in the radial forces due to the different treatment of the vertical forces, which in turn might vary the outcome of the analysis and the conclusions derived from this methodology. In this paper, we investigate the effect of 2-D vs 3-D numerical modelling in the gas flows of rigidly rotating bar potentials. In order to test this effect, we have modelled the gas flow in NGC~4123 using the mass model and images provided by B.J. Weiner but using the 3-D N-body/SPH code used in PFF04. The parameters chosen for the SPH simulations have been taken from WSW01. In future work we will address the comparison between the sticky particle codes and the fluid dynamical methods. The modelling is presented in Section~\ref{modelling}, the results and comparison are presented in Section~\ref{comparison}. Finally, discussion and conclusions are presented in Section~\ref{conclusions} | \label{conclusions} The position and strength of the shocks in barred galaxies derived from dynamical models is robust to the chosen model. Here, we have compared 2-D to a 3-D code and despite the differences in the two codes and the small differences in the mass model, the resulting velocity and density maps for both models is very similar. However, for similar hydrodynamical parameters, the SPH code develops a central ring, associated to the x$_{2}$ family of orbits while the gas distribution in the FS2 model does not present this central ring. This difference seems to be due to differences between the two codes and not to the difference on the treatment of the vertical forces. In any case, the position and the strength of the shocks in both models remains very similar. We can therefore conclude that the methodology used to derive the halo contribution to the potential using dynamical modelling of the gas flow in the bar region of spiral galaxies is robust to the codes used and to wether they are 2 or 3-D codes. | 7 | 10 | 0710.2855 |
0710 | 0710.0257_arXiv.txt | \hspace{0.6cm} The thickness of the equilibrium isothermal gaseous layers and their volume densities $\rho_{gas}(R)$ in the disc midplane are calculated for 7 spiral galaxies (including our Galaxy) in the frame of self-consistent axisymmetric model. Local velocity dispersions of stellar discs were assumed to be close to marginal values necessary for the discs to be in a stable equilibrium state. Under this condition the stellar discs of at least 5 of 7 galaxies reveal a flaring. Their volume densities decrease with $R$ faster than $\rho_{gas}$, and, as a result, the gas dominates by the density at the disc periphery. Comparison of the azimuthally averaged star formation rate $SFR$ with the gas density shows that there is no universal Schmidt law $SFR \sim \rho_{gas}^n$, common to all galaxies. Nevertheless, $SFR$ in different galaxies reveals better correlation with the volume gas density than with the column one. Parameter $n$ in the Schmidt law $SFR \sim \rho_{gas}^n$, formally calculated by the least square method, lies within 0.8\,--\,2.4 range and it's mean value is close to 1.5. Values of $n$ calculated for molecular gas only are characterized by large dispersion, but their mean value is close to 1. Hence the smaller $\rho_{gas}$ the less is a fraction of gas actively taking part in the process of star formation. To be published in Astronomy Reports, 2008. | \hspace{0.6cm}The key question of galaxy evolution is the dependence of star formation rate $SFR$ on the gas density and other interstellar medium parameters, averaged by large enough area or volume for smoothing random fluctuations in gas and young stars distributions. Scmidt~\cite{Schmidt59} suggested a simple form of star formation rate parameterization: $SFR_v\sim\rho_{gas}^n$, usually called the Schmidt law. From the analysis of gas and young objects distribution in the solar vicinity, Schmidt~\cite{Schmidt59} obtained $n\approx 2$. Later a large number of papers were published aimed to check and to interpret the Schmidt law, based on the gas and $SFR$ distributions in the discs of spiral galaxies. In particular, it was found that both local, azimuthally averaged $SFR$ values as well as the total star formation rate in a galaxy correlate (although not too tightly) with the gas content in its disc (e.g. Madore et al., ~\cite{Madore74}, Kennicutt, ~\cite{Kennicutt89}, Wong\&Blitz, ~\cite{WB02}, Boissier et al., ~\cite{Bois03}, Shuster et al.,~\cite{Schu06}). Being empirical by it sence, the Schmidt law and its modifications open a possibility to calculate evolution models of galaxies, parameterizing the star formation history. This allows us to understand better mechanisms which regulate the rate of star formation. A situation is entangled by the fact that the dependence $SFR_v(\rho_{gas})$ (i.e. the local Schmidt law) and, in particular, the power $n$ can't be found directly from observations of other galaxies, because in order to estimate the volume-related values of $SFR_v$ and $\rho_{gas}$ it is necessary to know the gas layer thickness, which may vary significantly both along the galaxy radius and from one galaxy to another. Therefore in practice the Schmidt law is often replaced by the other one, outwardly alike empirical law $SFR_s\sim\sigma_{gas}^N$ (sometimes it is called Kennicutt\,--\,Schmidt law), where the compared values are scaled to unit disc surface area, or by the law $SFR_t\sim M_{gas}^N$ for the total values $SFR_{t}$ and gas mass $M_{gas}$ (so-called the global Schmidt law). These dependencies are more complicated for interpretation, because the compared parameters are the integrals of heterogeneous functions of distribution along of the line-of-sight (in the first case) or allover a whole disc (for the global law). In general case, parameters $n$ and $N$ must not coincide (see the discussion of the question in papers Madore et al., ~\cite{Madore74}, Tutukov, ~\cite{tut06}). Only if at any distance from disc plane $n=1$, the power $N$ is also equal to unit. In most cases values of $N$ that obtained for different galaxies lay within the limits of $1<N<2$ (Wong and Blitz, ~\cite{WB02}, Boissier et al.,~\cite{Bois03}, Shuster et al.,~\cite{Schu06}), but for some galaxies they proves to be more steep. For example, for M~33 $N>$3 (Heyer et al., ~\cite{HeyerAll04}). For the case of the global Schmidt law when different galaxies are compared, $N\approx 1.4-1.5$ (Kennicutt, ~\cite{Kennicutt98}, Li et al., ~\cite{li+06}). The scatter of $N$ remains large: for galaxies with similar gas masses $SFR_{t}$ may differ by an order of magnitude. Apparently, it is possible to reproduce the Schmidt law for different star formation models using some simplifying assumptions on mechanisms of self-regulation of large-scale star formation, or considering conversion of neutral gas into dense molecular clouds and formation of stars (see e.g., \cite{tut06,li+06,elm02,krumholz05,Gerristen97}). Note however that usually variation of the thickness of gas layer along galactic radius as well as from one galaxy to another is ignored, and $\rho_{gas}$ is accepted to be much smaller than the volume stellar density, that is not always correct. In the first part of this paper, the gas volume density in the plane of an equilibrium disc is calculated as a function of $R$ for several nearby galaxies. In the second part, the relation between the gas density and star formation rate is analyzed. Chosen galaxies strongly differs by their properties. They include: M~33 and M~101 --- multiarmed late-type galaxies (Sc), the first of them is rather small; interacting galaxy M~51 and galaxy M~100 which are distinguished by high molecular gas content; Seyfert galaxy M~106 where the ejections from nucleus and star formation burst in the inner region are observed; massive early-type Sab galaxy M~81 which possesses a large bulge and regular spiral structure, and, finally, our Galaxy. The main simplifications we use are the following. The gaseous layers in galaxies are assumed to be axisymmetric and being in hydrostatic equilibrium. The pressure of the gas is determined by its turbulent motion: $P_{gas}=\rho_{gas}\,C_z^2$, where $C_z$ is one-dimensional velocity dispersion, which is assumed to be constant (although different for atomic and molecular gas). These simplifying suggestions are definitely too tough for regions enveloped by the intensive star formation, for the inner discs deep inside dense bulges or in the neighborhood of active nucleus, and also for the far periphery of discs. Note however that gas velocity dispersion, although may slowly vary with the distance from the galaxy center, remains high enough even at large distances (see the discussion in paper Dib et al., ~\cite{dib06}). Magnetic field pressure gradient and thermal gas pressure play significantly lesser role in the formation of gas layer thickness, at least for the case of our Galaxy (see discussion of this question in Cox review,~\cite{Cox05}). It is essential that within the limits of approximations mentioned above the observed distribution of atomic (HI) and molecular (H$_2$) hydrogen disc thicknesses along the Galaxy's radius can be sufficiently explained (Narayan and Jog,~\cite{NJ02}). | \hspace{0.6cm}As it follows from this paper, the assumption of marginally stable stellar discs leads to conclusion that their thickness, at least in some galaxies, changes significantly (usually increases) along the radius. The thickness of equlibrium gas layer increases in all cases --- either nearly linearly within a wide range of $R$ (M~33, M~51, M~81 and M~101), or nonlinearly, with a positive second derivative (M~100 and the Galaxy). Volume gas density in the midplane in all cases decreases within the considered range of $R$ down to several units of $10^{-25}$ g/cm$^3$ . In all cases the stellar density decreases steeper, than the gas density, so at the peripheries of galaxies the gas midplane density becomes comparable with the stellar density or even exceeds it. To compare the half thicknesses (HWHM) with the densities of stellar discs and gaseous layers in the galaxies it is convenient to consider their parameters at a fixed radius, say, $R\approx 2R_0$ (see Table~\ref{tab5}). \begin{table}[h] \caption{Stellar disc and gas layers thicknesses and the densities of gas at $R\approx 2R_0$. \emph{Columns}: (1) --- a galaxy, (2) --- stellar disc thickness to radial scalelength ratio, (3) --- atomic gas disc thickness, (4) --- molecular gas disc thickness, (5) --- volume gas density (including helium).}\label{tab5} \begin{center} \begin{tabular}{ccccc} &&&&\\ Galaxy&$h_*(2\,R_0)/R_0$&$h_{\rm HI}(2\,R_0)$\,,\,pc& $h_{\rm H2}(2\,R_0)$\,,\,pc&$\rho_{gas}(2\,R_0)$\,,\,g/cm$^3$\\ &&&&\\ \hline &&&&\\ (1)&(2)\,&(3)\,&(4)\,&(5)\,\\ &&&&\\ \hline &&&&\\ M~33&0,18&87,8&57,6&$6,3\cdot10^{-24}$\\ &&&&\\ M~51&0,23&123,6&81,3&$3,5\cdot10^{-24}$\\ &&&&\\ M~81&0,34&129,2&85,7&$1,1\cdot10^{-24}$\\ &&&&\\ M~100&0,20&108,8&71,0&$6,1\cdot10^{-24}$\\ &&&&\\ M~101&0,19&173,8&114,8&$1,5\cdot10^{-24}$\\ &&&&\\ M~106$^*$&0,27&150,0&99,2&$2,2\cdot10^{-24}$\\ &&&&\\ Galaxy&0,11&100,3&61,4&$4,4\cdot10^{-24}$\\ &&&&\\ \hline \end{tabular} \end{center} $^*$For M~106 all the values are given for R\,=\,1,5\,$R_0$. \end{table} As it follows from the Table, the relative thickness of stellar disc $h_*/R_0$ at $R=2R_0$ lies within the range of 0.1\,--\,0.3. Our Galaxy is the most thin ($h_*/R_0 \approx 0.1$) in our sample, whereas the stellar discs of two giant early type spiral galaxies (M~81 and M~106) are almost three times thicker. The thicknesses of gas layers differ less than stellar ones in the galaxies. The most thin gas layer (in M~33) is less than twice thinner than the most thick one (in M~101). The gas density (at distance $2R_0$) in all cases is equal to several units of 10$^{-24}$\,g/cm$^{-3}$ with no obvious dependence on the morphological type: in M~101 (Sc) it is approximately the same as in M~81 (Sab), and in our Galaxy (Sb\,--\,Sbc) it is close to that in M~33. The comparison of the gaseous density with the star formation rates (figures\,\ref{fig9}a,~\ref{fig9}b, and~\ref{fig10}a) gives an evidence that the replacement of surface gas density (Figure~\ref{fig9}b) by the gas volume density leads to more tight dependencies both for the ``surface'' and ``volume'' star formation rates. The reason for it is rather evident: star formation takes place in a narrow layer near the midplane, hence this process is sensitive namely to the volume gas density, whereas the directly measured surface gas density depends, in addition, to the gas layer thickness (more precisely --- to the vertical density profile), which differs from galaxy to galaxy. It is worth to remind that the regions that are very close to the center (in all galaxies but M~33), or are located at the disc peripheries (beyond the well defined spiral arms) need a special study and were not considered in this paper. Indeed, here we ignored the influence of the bulge on the vertical density profiles. The accepted value of the turbulent velocity may also be not suitable for all radii. A formal inclusion of the most inner and outer regions into our diagrams leads to the resulting dependencies ``gas density\,--\,$SFR$''\,\, which are less regular and more different between galaxies. Of the galaxies we consider, M~106 stands out by almost constant density $\rho_{gas}$ between 4 and 10\,kpc. The increasing of stellar disc thickness along the radius in this galaxy accompanies by the increasing of the surface gas density out of the centre within the wide range of $R$. By this reason this galaxy differs from the others in the diagrams ``$SFR$\,--\,gas density\,''. The similar abnormal behavior of this galaxy on the Kennicutt-Schmidt diagram was also revealed by Boissier et al.~\cite{Bois03}. However the molecular gas density in M~106 behaves in the common manner i.e. steeply decreases with $R$, that explains its ``normal'' position at the diagram $SFR_v\,-\,\rho_{\rm H2}$. Note, that this galaxy is not typical in other aspects: it differs from the others by the presence of the active nucleus, as well as by the very active star formation in the central part of the disc. The values of $n$ in the Schmidt law $SFR_v\,\sim\rho_{gas}^n$ for the galaxies we consider are rather different (see Table~\ref{tab4}). For all galaxies except M~101 the exponent $n>1$, and its mean value is close to ``standard''\,\,value $\approx 1.4$ for ``global''\,\,Kennicutt-Schmidt law (column (3) of the Table). However, the mean value of $n$ approaches unit, if to compare the $SFR_v$ with the molecular gas density (column (4) in the Table). The most ``steep''\,\, relationships $SFR(\rho_{gas})$ are obtained for M~33, M~81 and the Galaxy. In the case of M~33 it is caused by the steep decreasing of the $SFR$ with the galactocentric distance (see Figure~\ref{fig8}). In the case of M~81, it is the consequence of the atypically slow decrease of the gas volume density along $R$ (see Figure~\ref{fig3}) when the radial gradient of $SFR$ is moderate. For these three galaxies, the exponent $n$ is close to the ``classical''\,\,value $n=2$, suggested earlier by Schmidt. In conclusion, parameter $n$ by no means can be considered as a universal one. The Schmidt law has a very approximate nature, and it fulfils much better when the volume instead of the surface gas density is used. In this case the exponent $n<2$ (with some possible exceptions). It is important that in spite of a large dispersion its value becomes closer to unit in the mean if to replace the total gas density $\rho_{gas}$ with the molecular gas density $\rho_{\rm H_2}$. It is evident that relationship between the volume gas density and $SFR$ for galaxies should not be too tight because the process of star formation depends on a number of parameters besides the mean gas density. The latter seems to be a crucial factor only for the most dense gas such as the gas in the nuclei of molecular clouds where the HCN radiation comes from. Indeed, as observations show, the star formation rate depends linearly on the gas mass determined by the HCN line intensity, i.e for the most dense molecular gas $n\approx 1$ (Gao and Solomon,~\cite{gao04a,gao04b}). Although the number of galaxies we considered is small, the results allow to propose that the value of $n$ is on average close to unit even for the less dense molecular gas, that reveals itself in CO line. Since $n>1$ for the total ($HI+H_2$) gas density, one may conclude that not only star formation rate, but also the efficiency of star formation ($SFR$ per gas mass unit) decreases along with $\rho_{gas}$. In other words, the less is the gas density, the longer time the gas remains in the rarified atomic state (that is a time scale of gas consumption is larger). It agrees with general conception, according to which the fraction of the interstellar gas which takes active part in the process of a star formation, decreases when the mean total density of gas becomes lower (see the discussion in~\cite{WB02,elm03}). The other important factor, besides the gas density, that determines the star formation rate at a given radius $R$ is the surface density of the old stellar disc $\sigma_{*}$ (see papers~\cite{ZasAbr06,BdeJ00,DR94}). As it was shown by Zasov and Abramova~\cite{ZasAbr06} on the example of four well-studied galaxies, local star formation efficiency defined as $SFE\,=\,SFR$/$\sigma_{gas}$ changes approximately as $\sigma_{*}^{\footnotesize 0,7}$ in a wide interval of $R$. Partially this relationship may be explained by higher volume gas density (for a given surface density) in those regions, where $\sigma_{*}$ is higher (that is closer to galaxy center). Nevertheless this relationship can not be reduced to the simple exponential Schmidt law (neither ``volume'' nor ``surface'' one) because the clearly defined relationship between the gas density and surface disc density is absent. It evidences the existence of more deep connection between the present-day star formation on the one hand and already formed stellar disc and the gas density on the other hand, which cannot be described by simple empiric laws. \bigskip The authors thank Igor V. Abramov for assistance in the numerical solutions and D.Bizyaev for helpful discussion. \bigskip This work was supported by the Russian Fond of Basic Researches grant 07-02-00792. | 7 | 10 | 0710.0257 |
0710 | 0710.2252_arXiv.txt | {}{To reexamine the implications of the recent HESS observations of the blazar 1ES0229+200 for constraining the extragalactic mid-infrared background radiation.}{We examine the effect of \gray ~absorption by the extragalactic infrared radiation on predicted intrinsic spectra for this blazar and compare our results with the observational data.}{We find agreement with our previous results on the shape of the IR spectral energy distribution (SED), contrary to the recent assertion of the HESS group. Our analysis indicates that 1ES0229+200 has a very hard intrinsic spectrum with a spectral index between 1.1 $\pm$ 0.3 and 1.5 $\pm$ 0.3 in the energy range between $\sim$0.5 TeV and $\sim$15 TeV.} {Under the assumptions that (1) the SED models of Stecker, Malkan \& Scully (2006) are reasonable as derived from numerous detailed IR observations, and (2) spectral indexes in the range $1 < \Gamma < 1.5$ have been shown to be obtainable from relativistic shock acceleration under the astrophysical conditions extant in blazar flares (Stecker, Baring \& Summerlin 2007), the fits to the observations of 1ES0229+200 using our previous IR SEDs are consistent with both the IR and \gray\ observations. Our analysis presents evidence indicating that the energy spectrum of relativistic particles in 1ES0229+200 is produced by relativistic shock acceleration, producing an intrinsic \gray ~spectrum with index $1 < \Gamma < 1.5$ and with no evidence of a peak in the SED up to energies $\sim$ 15 TeV.} { | Shortly after the first strong \gray ~ blazar 3C279 was discovered by the EGRET detector aboard the Compton Gamma Ray Observatory (CGRO), Stecker, De Jager \& Salamon (1992) proposed that the study of the spectra of such sources could be used to probe the intergalactic infrared radiation. At that time, there were no actual direct observations of the diffuse extragalactic IR background or extensive observations of the sources of such radiation. The idea was to look for the effects of photon-photon annihilation interactions into electron-positron pairs. The cross section for this process is exactly determined; it can be calculated using quantum electrodynamics (Breit \& Wheeler 1934). Thus, in principle, if one knows the emission spectrum of an extragalactic source at a given redshift, one can determine the column density of photons between the source and the Earth. In the last 15 years, great advances have been made in extragalactic infrared astronomy. The diffuse background at wavelengths not totally dominated by galactic or zodiacal emission has been measured by the Cosmic Background Explorer (COBE). In addition, there have been extensive observations of infrared emission from galaxies themselves, whose total emission is thought to make up the cosmic infrared background (see review by Hauser \& Dwek 2001). The latest extensive observations have been made by the Spitzer satellite. It is thus appropriate to use a synoptic approach combining the TeV \gray ~observations with the extragalactic infrared observations in order to best explore both the TeV emission from blazars and the diffuse extragalactic infrared radiation. Aharonian et al. (2007) have recently observed the spectrum of the BLLac object 1ES0229+200 up to an energy $\sim$~15 TeV with the High Energy Spectroscopic System (HESS). Then, by assuming that the intrinsic spectral index of this source is greater than 1.5, they drew conclusions regarding wavelength dependence and flux of the mid-IR extragalactic background radiation. Their conclusions regarding the mid-IR extragalactic background radiation appear to disfavor the results of the extensive semi-empirical calculations of the extragalactic IR background spectrum given by Stecker, Malkan \& Scully (2007) (SMS). In this paper, we will reexamine the assumptions and conclusions of Aharonian et al. (2007) and show that the observations of 1ES0229 are fully consistent with the diffuse IR background spectrum obtained by SMS. Furthermore, we will show that 1ES0229+200 is an example of a set of blazars which exhibit very hard spectra that are indicative of relativistic shock acceleration. | Aharonian et al. (2007), by assuming that blazar spectra have spectral indexes $\Gamma > 1.5$, concluded that the mid-IR spectral energy distribution (SED) must have a wavelength dependence steeper than $\lambda ^{-1}$, in their terminology $\lambda^{-\alpha}, \alpha > 1.1 \pm 0.25$ in the wavelength range between 2\mic ~ and 10\mic . On the contrary, galaxy emission models which take into account emission by warm dust and PAH and silicate emission in the mid-IR, as well as direct mid-IR observations of galaxy spectra give values for $\alpha$ in the range between $\sim$ 0.7 and $\sim$ 0.8 (Spinoglio et al 1995, Xu et al. 2001). These values are consistent with the 2-10\mic ~diffuse background SED given in the semi-empirical SMS models. Such values lead to an energy dependence for the \gray ~optical depth which is consistent with that obtained by both Totani \& Takeuchi (2002) and SMS. This result is also supported by the lower limit on the 15\mic ~diffuse background flux obtained by galaxy counts of 3.3 $\pm$ 1.3 nW m$^{-2}$sr$^{-1}$ (Altieri et al. 1999) which, under the conservative assumption that 80\% of the mid-IR flux is resolved out (SMS), yields a value for the total diffuse background flux at 15\mic ~ of 4.1 $\pm$ 1.6 nW m$^{-2}$sr$^{-1}$, higher than the upper limit of 3.1 nW m$^{-2}$sr$^{-1}$ derived by Aharonian et al. (2007) under the assumption that $\Gamma > 1.5$. Under the assumptions that (1) the SED models of Stecker, Malkan \& Scully (2006) (SMS) are reasonable as derived from numerous detailed IR observations, and (2) spectral indexes in the range $1 < \Gamma < 1.5$ have been shown by Stecker, Baring \& Scully (2007) (SBS) to be obtainable from relativistic shock acceleration under the astrophysical conditions extant in blazar flares, the fits to the HESS TeV observations of 1ES0229+200 using the SMS infrared SEDs are consistent with both the IR and \gray\ observations. The SBS simulations indicate with specific test runs that electron spectra with asymptotic spectral indexes between 1.26 and 1.62 can be obtained from acceleration by relativistic shocks with bulk Lorentz factors between 10 and 30 and these electrons can then Compton scatter to produce \gray~ spectra with indexes $\sim$1.1 and $\sim$ 1.3. The simulations show that larger bulk Lorentz factors lead to flatter spectra. Such Lorentz factors of 50 or more have been implied from studies of specific flares in BL Lac objects (Konopelko et al. 2003; Bagelman, Fabian \& Rees 2007) so that the existence of highly relativistic shocks in such sources with indexes as flat as $\sim$1 is not unreasonable. For the SMS fast evolution SED, the \gray~ index obtained is 1.11 and for the baseline SED, the index obtained is 1.45. Our analysis thus presents evidence for relativistic shock acceleration in 1ES0229+200 that results in a very hard intrinsic \gray ~spectrum with no evidence of a peak in the \gray ~ SED up to energies $\sim$ 10 TeV. Unfortunately, there are no simultaneous observations at other wavelengths that can be used to model the flare that occurred in 1ES0229+200 and, in any case, TeV orphan flares have been observed in other BL Lac objects. However, we note that hard X-ray spectra have been seen in other sources as discussed in Section 3. We also note that SBS have derived spectra for three other BL Lac objects at redshifts between 0.18 and 0.19 with indexes between 1 and 1.5, viz., 1ES1218+30, 1ES1101-232, and 1ES0347-121. It therefore appears that a whole class of blazars or blazar flares may exhibit the hard-spectrum characteristics of relativistic shock acceleration. | 7 | 10 | 0710.2252 |
0710 | 0710.2314_arXiv.txt | Deep \textit{Spitzer} IRAC images of L1157 reveal many of the details of the outflow and the circumstellar environment of this Class 0 protostar. In IRAC band 4, 8 $\mu$m, there is a flattened structure seen in absorption against the background emission. The structure is perpendicular to the outflow and is extended to a diameter of $\sim$2$\arcmin$. This structure is the first clear detection of a flattened circumstellar envelope or pseudo-disk around a Class 0 protostar. Such a flattened morphology is an expected outcome for many collapse theories that include magnetic fields or rotation. We construct an extinction model for a power-law density profile, but we do not constrain the density power-law index. | \label{dist} The L1157 dark cloud in Cepheus (IRAS 20386+6751) conceals a young protostar, a so-called Class 0 source \citep{andre1993}, which is deeply embedded within a large circumstellar envelope \citep{gueth2003,beltran2004}. L1157 has a large powerful molecular outflow that is the prototype of chemically active outflows \citep{bachiller2001}. Despite the attention that L1157 has received at radio wavelengths, few observations have been made in the near to mid-infrared outside of observations of the outflow in H$_2$ and K-band \citep[e.g.][]{davis1995,cabrit1998}. Only recently have sensitive instruments been available to observe these objects shortward of 10 $\mu$m \citep[e.g.,][]{tobin2007}. The outflow carves cavities in the circumstellar envelope, which allow photons from the embedded central source to escape and scatter off dust in the cavity at NIR wavelengths. The morphology of the scattered light can be used to probe many of the fundamental properties of the source such as opening angle, envelope mass, etc. \citep[e.g.,][]{whitney2003a,whitney2003b,tobin2007,robitaille2007,seale2008}. In this letter, we present new, deep \textit{Spitzer Space Telescope} observations of L1157. The IRAC continuum emission is dominated by molecular line emission in the outflow. Near the source there is a small amount of emission that may be attributed to scattered light and perhaps molecular line emission that is highly excited by the outflow jet. Other than the enormous outflow ($\sim$0.5 pc) to the north and south, the most prominent feature observed is a large, flattened absorption feature at 8.0 $\mu$m and less defined at 5.8 $\mu$m. This absorption feature is a flattened circumstellar envelope observed in silhouette against the Galactic infrared background. The distance to L1157 is important to any physical interpretation, but the distance is highly uncertain. The molecular clouds in Cepheus have three characteristic distances, 200, 300, 450 pc \citep{kun1998}; L1157 has a similar galactic latitude as the 200 pc and 300 pc absorbing clouds. Due to this, we disagree with the current accepted distance of 440 pc. This value was based upon a study of NGC 7023 in \citet{viotti1969}; L1157 is not in clear association with this cluster. In this letter, we use a distance of 250 pc. | The shape of the absorption feature is especially intriguing, as it looks like a disk structure perpendicular to the outflow axis. This is the first clear detection of a flattened envelope or pseudo-disk in a Class 0 object. \cite{galli1993a,galli1993b} have shown that a modest magnetic field structure modifies infall from the initial spherical cloud to form a so-called ``pseudo-disk''; a flat thin structure in the equatorial plane that is not rotationally-supported, thus collapsing. This type of structure is also seen from simple flat sheet models of collapse \citep[e.g.,][]{hartmann1994,hartmann1996}, as well as detailed ambipolar diffusion models \citep[e.g.,][]{fiedler1993}. On the other hand, this structure is large $\sim$15,000-30,000 AU, depending on the background used. That size is somewhat larger than the inner envelope size estimated from interferometric models of the dust continuum \citep{looney2003}. However, the single-dish dust emission \citep{gueth2003} is extended along the same axis as the absorption, which argues that the inner envelope in the equatorial plane either has higher density, and/or different dust opacity properties. Our modeling results show that the properties of the structure, as determined by the absorption model, are consistent with the above theoretical constructs, i.e. flattened envelopes and density profiles. Although we model a range of indexes ranging from p = 0.5 to 3, only the 0.5 to 2 provide acceptable fits at the 90\% confidence level with the vast majority of fits being p=1.5. To better explore the physical meaning the model, we assume dust opacities (dust plus gas) from \citet{lidraine2001} of $\kappa_{8.0\mu m} =5.912~cm^2~g^{-1}$. Although using interstellar dust opacities for $\kappa_\lambda$ is probably incorrect, as Class 0 sources are thought to have already experienced some grain growth \citep[e.g.,][]{looney2003,natta2007}, it is still a useful approximation. Using the assumed $\kappa_{8.0\mu m}$, the derived range for the density reference, $\rho_0$, or 1 pixel (1$\farcs$2) from the center of the envelope (i.e. 300 AU at 250 pc) is listed in Table \ref{fits}. In addition, we can estimate the absorbing mass of the flattened envelope component for each ``best fit'' model of Figure \ref{fit} and a height of 3 pixels (the size of the box we averaged over). We assume that the vertical density profile is constant for the mass estimate, even though the observed absorption falls off vertically with scale heights of $\sim$3-4 pixels, using a Gaussian vertical structure. Our mass estimate, listed in Table \ref{fits}, ranges from 0.08 to 0.16 M$_\odot$. Without using any model, we can also estimate the mass required for the observed extinction. We use the above value for $\kappa_{8.0\mu m}$ and a background of 0.54 MJy/sr to calculate the mass necessary for the extinction of all pixels in the central region below 0.458 MJy/sr (577 pixels). The total absorbing mass required for those pixels is 0.19 M$_{\odot}$. It is important to note that this mass is only in the absorbing pseudo-disk, but the mass is comparable to the 140$\arcsec$ extended envelope detected in the millimeter continuum as discussed in \citet{gueth2003} with an estimated mass of 0.7 M$_{\odot}$ at an assumed distance of 250 pc. This implies that although a large fraction of the mass of the envelope is in the flattened structure, most of the mass is more diffuse. Due to the log nature of the mass absorption, the density contrast from the center of the absorption feature to slightly offset from the absorption feature is approximately an order of magnitude. This is more than expected from the numerical models of \cite{galli1993b}, but somewhat consistent with the ambipolar models of \cite{fiedler1993} and the sheet collapse model of \cite{hartmann1996}. As can be seen in Table \ref{fits} and Figure \ref{fit}, these data can not constrain the model density profiles directly; multiple power-laws are allowed. On the other hand, although many theoretical models do suggest a flattened envelope structure, the power-law index in the pseudo-disk is expected to evolve. In that case, we would not expect a single power-law to well describe the envelope density. In fact, the ambipolar models \citep[e.g.][]{tassis2005a,tassis2005b} suggest that the power-law can be episodic in the flattened envelope. Further studies with increased sensitivity need to be compared directly to the theoretical density profiles in the flattened objects to say anything more completely. | 7 | 10 | 0710.2314 |
0710 | 0710.0311_arXiv.txt | We derive probability density functions for the projected axial ratios of the real and mock 2PIGG galaxy groups, and use this data to investigate the intrinsic three dimensional shape of the dark matter ellipsoids that they trace. As well as analysing the raw data for groups of varying multiplicities, a convolution corrected form of the data is also considered which weights the probability density function according to the results of multiple Monte-Carlo realizations of discrete samples from the input spatial distributions. The important effect observed is that the best fit distribution for all the raw data is a prolate ellipsoid with a Gaussian distribution of axial ratios with $\bar{\beta}=0.36$ and $\sigma=0.14$, whilst for the convolved data the best fit solution is that of an oblate ellipsoid $\bar{\beta}=0.22$ and $\sigma=0.1$. Previously only prolate distributions were thought compatible with the data, this being interprated as evidence of filamentary collapse at nodes. We also find that even after allowing for the sampling effects, the corrected data is better fit using separate multiplicity bins, which display a trend towards more spherical halos in higher multiplicity groups. Finally, we find that all results in the real data are in good agreement with the mock data from $\Lambda$CDM simulations, KS tests showing that all comparative data have been drawn from the same distributions within the $1\sigma$ confidence limits. | The spatial distribution of galaxies traces the shape of the dark matter potential in which they are embedded. Simulations show that dark matter halos are not spherical, as one would naively expect because of dark matter's non-dissipational nature, but are strongly flattened triaxial ellipsoids \citep{dubinski91}. Although prolate and oblate shapes are equally likely in the simulations, dissipative infall of baryonic gas eventually forces the halo shapes toward pronounced oblateness with axial ratios $b/a > 0.7$ \citep{katz91, dubinski94,combes02}. Because dark matter structures evolve self-similarly it is possible to derive the general shape of dark halos from analysis of single test objects. The first such attempts were carried out on rich clusters, whose shapes were found to be triaxial and strongly prolate \citep{plionis01}: however the non-linear evolution of these objects may play a role in shaping their density distributions \citep[e.g][]{binney79}. Indeed there is evidence that the more high-density systems are more spherical than low-density objects \citep[random gaussian field work of][]{bardeen86}. The observed trends are the opposite of what is observed in simulations \citep[e.g.][]{kasun05,allgood06}. The distribution of Milky Way satellites shows a strongly oblate and flattened distribution \citep{ruzicka07}, which is also consistent with analysis of the orbital planes of the Sagittarius dwarf \citep{johnston05} and the Monoceros stream \citep{penarrubia05}. However, star counts \citep{lemon04} and analysis of the stellar stream of the Sagittarius dwarf \citep{ibata01} support a spherical halo. It is also uncertain whether the dwarf satellites can be used as test particles, as they may not originate from cosmological sub-structures \citep{kroupa05}. Groups of galaxies are likely to be the best testbeds to study the shapes of dark matter halos. Large group catalogs are now available from redshift surveys (e.g., the 2dF Percolation Inferred Galaxy Groups [2PIGG] of \citealt{eke04} and the group catalog of \citealt{merchan05} from Data Release 3 of the Sloan Digital Sky Survey) and these allow a statistical approach to the study of the shapes of groups and the shape dependence on richness, multiplicity and dynamical evolution. Early studies marginally favoured prolate shapes \citep{fasano93,orlov01}, but not at the exclusion of oblate solutions. The most recent study of 2PIGG groups by \cite{plionis06} favours prolate groups, but is obtained by means of a multiplicity cut and as such represents a different method to the one presented here. Prolate results indicate that the original (oblate ?) shapes may have been strongly modified during gravitational collapse, or that filamentary collapse is in fact being witnessed. In this paper we carry out a detailed analysis of the shape of 2PIGG groups, coupled with extensive Monte Carlo simulations and comparison with groups extracted from the 2dF mock catalogs used by \cite{eke04} to optimize group selection. In section 2 we present our method for analysing the shape of observed groups and in section 3 apply this to the real and mock 2PIGG catalog \citep{eke04}. Comparisons are made between the raw data and data that is corrected for the finite (indeed, often sparse) sampling by factors which we determine from Monte-Carlo simulations of suitably populated groups. Furthermore, the mock and real data are considered separately and KS tests are used to confirm their distributional similarities. | We found good agreement between our results for raw data with a simple 20+ multiplicity cut and findings by other authors, indicating a strongly prolate distribution with a mean $\bar{\beta}=0.46$ and $\sigma_{\beta}=0.14$. However, when fits were applied to convolution corrected distributions which allow for the error introduced by finite sampling we found that for the higher multiplicity cuts both oblate and prolate distributions could produce reasonable fits. More significantly, evidence suggests that these large multiplicity populations share the same underlying distribution (if oblate: $\bar{\beta}\sim0.3$ $\sigma_{\beta}=0.1$, if prolate $\bar{\beta}\sim0.44$ $\sigma_{\beta}=0.14$). There is evidence in the data that low multiplicity groups have a highly elliptical sub population that is missing in other multiplicity cuts: for 5-9 multiplicity cuts oblate distributions produce a better quality of fit ($\bar{\beta}=0.2$, $\sigma_{\beta}=0.1$), this is despite the effect of the convolution correction being to reduce this signal. This is seen as evidence of high multiplicity groups being more spherical, and is consistant with them being located in nodes, or at least in a less filamentary structure than the extremely low multiplicity groups. Large multiplicity groups are better fit by a prolate distribution in the corrected data, but oblate fits are not as strongly rejected as in earlier work which did not account for the sampling bias. The effect of interlopers was extensively considered, and the results of this work imply all measurements obtained should be considered upper limits for both the under lying axial ratios (i.e. the true non-interloper distribution will be more elliptical than that measured) and the standard deviations of the populations. KS tests were utilised where appropriate and indicate that the real and mock catalogs are in very good agreement over all ranges. For all multiplicity comparisons the real and mock data share the same underlying distribution to within $1\sigma$ expectations, and the most self similar populations are for multiplicities 10-19 and 20+, which appear consistent with being part of the same overall population once the corrections have been applied. Further exploration of projected group shapes will be presented in a future paper which will consider the differences we find when colour cuts and different grouping algorithms are used. | 7 | 10 | 0710.0311 |
0710 | 0710.1691_arXiv.txt | Dynamical mass estimates of ultra-compact dwarfs galaxies and massive globular clusters in the Fornax and Virgo clusters and around the giant elliptical Cen\,A have revealed some surprising results: 1) above $\sim10^6 M_\odot$ the mass-to-light ($M/L$) ratio increases with the objects' mass; 2) some UCDs/massive GCs show high $M/L$ values (4 to 6) that are not compatible with standard stellar population models; and 3) in the luminosity-velocity dispersion diagram, UCDs deviate from the well defined relation of ``normal'' GCs, being more in line with the Faber-Jackson relation of early-type galaxies. In this contribution, we present the observational evidences for high mass-to-light ratios of UCDs and discuss possible explanations for them. | The so-called ultra-compact dwarf galaxies (UCDs) are very massive ($10^6 M_\odot<M<10^8 M_\odot$), old, compact stellar systems that were discovered in nearby galaxy clusters about a decade ago (\cite{hilk99}, \cite{drin00}). Their nature is unknown yet. Maybe they are remnant nuclei of disrupted galaxies, or maybe they are merged stellar super-clusters formed in interacting galaxies. Regardless of what UCDs actually are, some properties divide them from ``ordinary'' globular clusters (GCs). The half-light radii of UCDs scale with luminosity reaching $\sim90$ pc for the most massive UCDs. Unlike for GCs, their densities within the half-light radii are not increasing with mass but stay at a constant level or even decrease. Thus UCDs are not that compact at all when compared to $10^6 M_\odot$ GCs, but certainly much denser than dwarf ellipticals of comparable mass. | 7 | 10 | 0710.1691 |
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0710 | 0710.4282_arXiv.txt | {Charge transfer (or exchange) reactions between hydrogen atoms and protons in collisionless shocks of supernova remnants (SNRs) are a natural way of producing broad Balmer, Lyman and other lines of hydrogen.} {We wish to quantify the importance of shock-induced, non-thermal hydrogen emission from SNRs in young galaxies.} {We present a method to estimate the luminosity of broad ($\sim 1000$ km s$^{-1}$) Ly$\alpha$, Ly$\beta$, Ly$\gamma$, H$\beta$ and P$\alpha$ lines, as well as the broad and narrow luminosities of the two-photon (2$\gamma$) continuum, from existing measurements of the H$\alpha$ flux. We consider cases of $\beta=0.1$ and 1, where $\beta \equiv T_e/T_p$ is the ratio of electron to proton temperatures. We examine a modest sample of 8 proximate, Balmer-dominated SNRs from our Galaxy and the Large Magellanic Cloud. The expected broad Ly$\alpha$ luminosity per object is at most $\sim 10^{36}$ erg s$^{-1}$. The 2$\gamma$ continuum luminosities are comparable to the broad H$\alpha$ and Ly$\alpha$ ones. We restrict our analysis to homogenous and static media.} {Differences in the Ly$\alpha$/H$\alpha$ and Ly$\beta$/H$\alpha$ luminosity ratios between the $\beta=0.1$ and 1 cases are factors $\sim 2$ for shock velocities $1000 \lesssim v_s \lesssim 4000$ km s$^{-1}$, thereby providing a direct and unique way to measure $\beta$. In principle, broad, ``non-radiative'' Ly$\alpha$ from SNRs in young galaxies can be directly observed in the optical range of wavelengths. However, by taking into consideration the different rates between core collapse and thermonuclear supernovae, as well as the duration we expect to observe such Ly$\alpha$ emission from SNRs, we expect their contribution to the total Ly$\alpha$ luminosity from $z \sim 3$ to 5 galaxies to be negligibly small ($\sim 0.001 \%$), compared to the radiative shock mechanism described by Shull \& Silk (1979). Though broad, non-thermal Ly$\alpha$ emission has never been observed, these photons are produced in SNRs and hence the non-radiative Ly$\alpha$ luminosity is a part of the intrinsic Ly$\alpha$ spectrum of young galaxies.} {} | Observations of galaxies at high redshifts have revealed a broad class of Ly$\alpha$-emitting galaxies at $z \sim 3$ to 5 (e.g., Tapken et al. 2007). The Ly$\alpha$ emission from these objects is reaching us as light in the visible spectral band, enabling their study using large, ground-based optical telescopes, which in turn permits detailed spectroscopic studies of these galaxies. Observations of quasars at $z \sim 6$ (e.g., Fan et al. 2006) have revealed heavy elemental abundances exceeding solar values. We know that at least some of the galaxies at $z \sim 3$ to 5 have high abundances of heavy elements, facilitating the formation of dust. In homogenous and static media, the dust particles impede the escape of Ly$\alpha$ emission from gas-rich galaxies, due to the small mean free paths of the photons, low temperatures of the gas and ultimately high probabilities of absorption. In clumpy media, dust can enhance the escape of Ly$\alpha$ photons relative to the continuum (Neufeld 1991; Hansen \& Oh 2006). Broadening of Ly$\alpha$ lines due to multiple scatterings is a slow process requiring a long diffusion time (though velocity fields in the interstellar medium may broaden the Ly$\alpha$ lines and reduce the diffusion time). Hence, there is special interest in the physical processes that are able to naturally produce extremely broad wings in Ly$\alpha$ lines, which may permit the photons to leave the host galaxy without requiring many scatterings (but see \S\ref{sect:discussion}). Among obvious mechanisms is the one at work in the unique massive binary SS433 (for a recent review, see Fabrika [2004]), with strongly blue- and redshifted H$\alpha$ and H$\beta$ lines, due to cooling and recombination of hydrogen in the baryon-dominated, precessing jet moving with velocity $\sim 0.26 c$. Such objects are very rare --- SS433 is the only such example in our Galaxy. More well-known Galactic sources of H$\alpha$ emission with broad line wings are the supernova remnants (SNRs) of Type Ia, emitting due to charge transfer (or ``charge exchange'') reactions between hydrogen atoms and protons in the blast wave penetrating the low-density ($\sim 1$ cm$^{-3}$), ambient gas. The widths of the H$\alpha$ lines correspond to Doppler broadening with velocities up to $\sim 5000$ km s$^{-1}$. The same process should produce not only H$\alpha$ emission, but photons in the Lyman series of hydrogen as well. Recently, some of these SNRs were observed in Ly$\beta$ using the {\it FUSE} spacecraft (Korreck et al. 2004; Ghavamian et al. 2007, hereafter G07). Knowledge of the cross sections of charge transfers to excited levels and excitation of the fast-moving hydrogen atoms permit us to find simple formulae relating the luminosities of SNRs in the broad H$\alpha$ and Ly$\alpha$ lines. The Ly$\alpha$ line should have a similar spectral distribution to the observed H$\alpha$ one in the broad wings, because the optical depth of the SNR for broad photons is negligibly small and the optical depth for coherent scattering (in the distant Lorentzian wings) in interstellar gas is low. We compile the existing data for core collapse and thermonuclear SNRs, including SNR 1987A (where the reverse shock is bright in the broad H$\alpha$ line), and present their theoretically expected, broad Ly$\alpha$ and Ly$\beta$ luminosities. For two objects, we present their expected broad Ly$\gamma$, H$\beta$ and P$\alpha$ luminosities. Taking into account the supernova (SN) rates, the luminosities of the SNRs in H$\alpha$ and the duration of their active phase (for the charge transfer mechanism described), we find that --- even without discussing the cosmological evolution of the SN rates --- the expected broad Ly$\alpha$ is several orders of magnitude lower than the estimate of Shull \& Silk (1979), who treated fully radiative SNRs with low metallicities and velocities (20 to 120 km s$^{-1}$). We come to the conclusion that the contribution of both core collapse and thermonuclear SNRs to the Ly$\alpha$ luminosity of young galaxies is negligibly small. In \S\ref{sect:obs}, we gather a modest sample of 8 Galactic and Large Magellanic Cloud (LMC) remnants, and use them as a template for estimating the expected Ly$\alpha$, Ly$\beta$, Ly$\gamma$, H$\beta$ and P$\alpha$ production. In \S\ref{sect:ratios}, we compute the Ly$\alpha$/H$\alpha$, Ly$\alpha$/Ly$\beta$, Ly$\beta$/H$\alpha$, Ly$\gamma$/H$\alpha$, H$\beta$/H$\alpha$ and P$\alpha$/H$\alpha$ luminosity ratios. We present our results in \S\ref{sect:results} and discuss their implications in \S\ref{sect:discussion}. | \label{sect:discussion} SNR 1987A is a unique example of a Balmer-dominated SNR. By virtue of adiabatic expansion cooling, the SN ejecta comprises mostly neutral hydrogen; it rushes out at velocities $\gtrsim 12,000$ km s$^{-1}$ (Michael et al. 2003; Heng et al. 2006). The non-radiative H$\alpha$ and Ly$\alpha$ result from the interaction of the ejecta with the {\it reverse shock} and not the blast wave (Heng 2007). As SNR 1987A has a Type II origin, it is possible to produce Balmer and Lyman lines via this mechanism; this is obviously not possible with Type Ia's. Smith et al. (2005) have predicted that the H$\alpha$ and Ly$\alpha$ emission from the reverse shock of SNR 1987A is shortlived ($\sim$ 2012 to 2014) and will be extinguished by the increasing flux of extreme ultraviolet (EUV) and X-ray photons traveling into the pre-shock region and ionizing the atoms --- pre-ionization. This is marginal evidence that broad Ly$\alpha$ from SNRs of a core collapse origin will be short-lived, i.e., $\lesssim 100$ years. In general, for this scenario to work, some interaction of the blast wave with the ambient material is needed, but if it is too strong the pre-shock gas becomes ionized (R. Chevalier 2007, private communication). To further investigate the viability of the short-lived, non-radiative Ly$\alpha$ hypothesis, we examine the sample of optically identified SNRs by Matonick \& Fesen (1997), who studied an ensemble of 12 SNR samples from different galaxies, including the Small Magellanic Cloud (SMC), LMC, M31 and M33, with distances up to 7 Mpc. In galaxies like NGC 2403, M81 and M101, the SNRs are associated with star-forming regions and most of them probably have a Type Ib/c origin. In most cases, the measured H$\alpha$ flux is $\sim 10^{-15}$ erg cm$^{-2}$ s$^{-1}$ and the inferred luminosity is $\sim 10^{36}$ erg s$^{-1}$. Since Matonick \& Fesen (1997) did not provide H$\alpha$ line profiles, it is impossible to estimate the proportion of the H$\alpha$ emission that is non-radiative. Furthermore, their selection criterion is based on picking out objects with [S~{\sc ii}]/H$\alpha \ge 0.45$, which will not detect SNRs with predominantly non-radiative H$\alpha$ emission. Shull \& Silk (1979) computed the temporally-averaged Ly$\alpha$ luminosity from radiative shocks of a population of Type II SNRs, assuming low metallicities, to be \begin{equation} L_{\rm{SS79}} = 3 \times 10^{43} \mbox{ erg s}^{-1} E^{3/4}_{51} n^{-1/2}_0 \dot{N}_{\rm{SN}}, \end{equation} where $\dot{N}_{\rm{SN}}$ is the number of supernovae (SNe) a year. They considered SNRs in both the ST and the PDS stages, and $v_s = 20$ to 120 km s$^{-1}$. Charlot \& Fall (1993) remark that the numerical coefficient in the preceding equation is about 40\% lower if one assumes solar metallicity. A very conservative upper limit on the broad Ly$\alpha$ from the Matonick \& Fesen (1997) samples can be obtained if one generously allows for all of the H$\alpha$ to be broad, for the shock velocities to be low ($\sim 500$ km s$^{-1}$) such that $\Gamma_{\rm{Ly}\alpha/\rm{H}\alpha} \sim 100$, and for the non-radiative emission to last $\sim 10^4$ years. Even in this very unlikely scenario, $L_{\rm{Ly}\alpha} \sim 10^{42}$ erg s$^{-1}$ is only about $0.1 L_{\rm{SS79}}$. Hence, our charge transfer mechanism is not energetically competitive. There is the possibility a SNR can produce both radiative and non-radiative components of H$\alpha$. Well-known examples are Kepler (Fesen et al. 1989; Blair, Long \& Vancura 1991) and RCW 86 (Long \& Blair 1990; Smith 1997). There is also the possibility that the non-radiative emission from the SNR is inhibited. For example, Foster (2005) observed and studied the Galactic SNR 3C 434.1 ($t_{\rm{age}} \approx 25,000$ yr; $d = 4.5 \pm 0.9$ kpc; possible Type Ib/c), which formed inside the eastern portion of a pre-existing stellar-wind bubble of interior density $\sim 0.1$ cm$^{-3}$. Strong H$\alpha$ emission ($6.1 \pm 0.4 \times 10^{36}$ erg s$^{-1}$) is measured from the eastern side; it is believed to be from a radiative shock. Being farther away from the western wall of the bubble, the shock on the western side is essentially still in free expansion and produces no measurable, non-radiative H$\alpha$. Our SNR sample and the considerations of SNR 1987A lead us to believe that if the short-lived emission contribution from Type Ib/c and Type II SNRs in young galaxies exists, it has a luminosity of \begin{equation} L_{\rm{Ly}\alpha,\rm{CC}} \sim 10^{38} \mbox{ erg s}^{-1} ~t_{\rm{emit},2} \dot{N}_{\rm{SN}}, \end{equation} where $t_{\rm{emit}} = t_{\rm{emit},2} 100$ years is the length of time we expect core collapse SNRs to produce shock-induced Ly$\alpha$ emission. On the other hand, thermonuclear SNRs are expected to have $t_{\rm{emit}} = t_{\rm{emit},4} 10^4$ years $\sim t_{\rm{PDS}}$. However, they are also believed to be much scarcer at high redshifts. For example, Dahlen et al. (2004) estimate that only 5\% to 7\% of available progenitors explode as Type Ia SNRs. Therefore, the expected luminosity is \begin{equation} L_{\rm{Ly}\alpha,\rm{Ia}} \sim 10^{38} \mbox{ erg s}^{-1} ~t_{\rm{emit},4} \dot{N}_{\rm{SN},-2}, \end{equation} where $\dot{N}_{\rm{SN},-2}$ is the number of SN per year in units of 0.01. We conclude that for both core collapse and thermonuclear SNRs, the expected luminosity from broad Ly$\alpha$ is only a $\sim 0.001\%$ effect, compared to the mechanism of Shull \& Silk (1979). Ly$\alpha$ line luminosities from $z \sim 3$ to 5 galaxies have been observationally determined to be $\sim 10^{42}$ to $10^{43}$ erg s$^{-1}$ (e.g., Saito et al. 2007), in general agreement with theoretical expectations. In addition, the lifetime of an emitting atom is approximately the length of time corresponding to one atomic length scale, and is only $t_{\rm mfp} \sim l_a/v_s \sim 10^7 n^{-1}_0 v^{-1}_{s,8}$ s, where $v_{s,8} = v_s/1000$ km s$^{-1}$ (H07). We have restricted our analysis to homogeneous and static media. Though broad, non-thermal Ly$\alpha$ emission has never been observed, these photons {\it are} produced in SNRs and hence the non-radiative Ly$\alpha$ luminosity is a part of the intrinsic Ly$\alpha$ spectrum of young galaxies. The optical depth for a broad photon in the line wings is (Verhamme, Schaerer \& Maselli 2006) \begin{equation} \tau \sim 0.26 T^{-1/2}_4 N_{{\rm H},20} v^{-2}_{w,8} b_{12.85}, \end{equation} where $T = 10^4 T_4$ K and $N_{\rm H} = N_{{\rm H},20} 10^{20}$ cm$^{-2}$ are the temperature and obscuring hydrogen column density of the medium, respectively. The turbulent velocity in the interstellar medium is $b = 12.85 b_{12.85}$ km s$^{-1}$ (Verhamme, Schaerer \& Maselli 2006), while $v_w = 1000 v_{w,8}$ km s$^{-1}$ is the velocity of the emitting atom in the line wings. Multiple scattering is important for $\tau \gtrsim 0.3$ (Chevalier 1986) and any realistic treatment of non-thermal Ly$\alpha$ lines in a young galaxy has to include radiative transfer effects, which we have neglected in our analysis. | 7 | 10 | 0710.4282 |
0710 | 0710.3141_arXiv.txt | We present a study of the globular cluster (GC) systems of nearby elliptical and S0 galaxies at a variety of wavelengths from the X-ray to the infrared. Our analysis shows that roughly half of the low mass X-ray binaries (LMXBs), that are the luminous tracers of accreting neutron star or black hole systems, are in clusters. There is a surprisingly strong correlation between the LMXB frequency and the metallicity of the GCs, with metal-rich GCs hosting three times as many LMXBs as metal-poor ones, and no convincing evidence of a correlation with GC age so far. In some of the galaxies the LMXB formation rate varies with GC color even within the red peak of the typical bimodal cluster color distribution, providing some of the strongest evidence to date that there are metallicity variations within the metal-rich GC peak as is expected in a hierarchical galaxy formation scenario. We also note that any analysis of subtler variations in GC color distributions must carefully account for both statistical and systematic errors. We caution that some published GC correlations, such as the apparent 'blue-tilt' or mass-metallicity effect might not have a physical origin and may be caused by systematic observational biases. | High resolution Chandra X-ray images of nearby ellipticals and S0s have resolved large numbers of point sources, confirming a long-standing suggestion that the hard X-ray emission in many of these galaxies is predominantly from X-ray binaries. In early type galaxies most of the bright, L$_X$$\gtrsim$10$^{37}$ erg s$^{-1}$ sources seen in typical Chandra observations must be low mass X-ray binaries, binary systems comprising a neutron star or black hole accreting via Roche lobe overflow from a low mass companion, since they generally have stellar populations that are at least a few Gyrs old. An important characteristic of LMXBs is that they are disproportionately abundant in globular clusters. Even though GCs account for $\lesssim$0.1\% of the stellar mass in the Galaxy, they harbor about 10\% of the L$_X$$\gtrsim$10$^{36}$ erg s$^{-1}$ LMXBs , indicating a probability of hosting a LMXB that is at least two orders of magnitude larger than for field stars. This has long been attributed to efficient formation of LMXBs in clusters due to dynamical interactions in the core. Early type galaxies are ideal for studies of the LMXB-GC link as they are particularly abundant in globular clusters. The identification of LMXBs with these simple stellar systems that have well defined properties such as metallicity and age provides a unique opportunity to probe the effects of these parameters on LMXB formation and evolution. \begin{figure}[b] \begin{center} \includegraphics[width=5in]{kundu_fig1.ps} \caption{Top: The V-I colors of GCs vs. distance from the center of NGC 4472 and GC color distribution. LMXB-GC matches are indicated by filled circles/bins. Bottom: V magnitude of globular clusters vs. half light radius and the globular cluster luminosity function. LMXBs are preferentially located in the brightest, most metal-rich GCs. There is a weak anti-correlation with GC half-light radius and no obvious correlation with galactocentric distance. Each of these broad correlations (or lack thereof) have been confirmed in other galaxies. } \label{fig1} \end{center} \end{figure} | 7 | 10 | 0710.3141 |
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0710 | 0710.1528_arXiv.txt | We present the general relativistic calculation of the energy release associated with a first order phase transition (PT) at the center of a rotating neutron star (NS). The energy release, $E_{\rm rel}$, is equal to the difference in mass-energies between the initial (normal) phase configuration and the final configuration containing a superdense matter core, assuming constant total baryon number and the angular momentum. The calculations are performed with the use of precise pseudo-spectral 2-D numerical code; the polytropic equations of state (EOS) as well as realistic EOSs (Skyrme interactions, Mean Field Theory kaon condensate) are used. The results are obtained for a broad range of metastability of initial configuration and size of the new superdense phase core in the final configuration. For a fixed ``overpressure'', $\delta\overline{P}$, defined as the relative excess of central pressure of a collapsing metastable star over the pressure of the equilibrium first-order PT, the energy release up to numerical accuracy {\it does not depend} on the stellar angular momentum and coincides with that for nonrotating stars with the same $\delta\overline{P}$. When the equatorial radius of the superdense phase core is much smaller than the equatorial radius of the star, analytical expressions for the $E_{\rm rel}$ can be obtained: $E_{\rm rel}$ is proportional to $(\delta\overline{P})^{2.5}$ for small $\delta\overline{P}$. At higher $\delta\overline{P}$, the results of 1-D calculations of $E_{\rm rel}(\delta\overline{P})$ for non-rotating stars reproduce with very high precision exact 2-D results for fast-rotating stars. The energy release-angular momentum independence for a given overpressure holds also for the so-called ``strong'' PTs (that destabilise the star against the axi-symmetric perturbations), as well as for PTs with ``jumping'' over the energy barrier. | Many theories of dense matter predict that at some density larger than the nuclear saturation density, a phase transition (PT) to some ``exotic'' state (i.e. boson condensate or quark deconfinement) occurs; for review see e.g., \cite{NS1,Glend.book,weber99}). A first-order PTs are particularly interesting from the astrophysical and observational point of view, because are associated with a meta-stable state of dense matter. One can thus expect, in the case of PTs occurring in the interior of NSs, the release of non-negligible amount of energy. Here we will focus on the basic features of the energy release-angular momentum independence; for complete description we refer the reader to \cite{erot,erots}. The text is arranged as follows: in Sect.~\ref{calculations} we briefly describe the results of calculations. Sect.~\ref{discussion} contains conclusions and open questions. \begin{figure}[t] \begin{center} \psfig{file=bejger1.eps,width=5.in} \end{center} \caption{ Left panel: Strong and weak PTs on mass-radius diagram (dotted line marks unstable configurations). Right panel: the energy release $E_{\rm rel}$ vs the over-pressure parameter $\delta\bar{P}$ for strong (two upper curves) and weak (lowest curve) PTs. Points are colored differently for different total angular momenta $J=(0,0.1,...,1.3)\times GM^2_\odot/c$ and follow the curve for the $J=0$ (non-rotating) configurations.} \label{reffig1} \end{figure} | \label{discussion} Our numerical calculations show the $E_{\rm rel}(J)$ independence during the weak as well as strong first-order PTs, for large rotation rates and large oblatnesses of the stars: it is therefore not an property of slow rotating stars only. The independence holds also for PTs with negative over-pressure $\delta\bar{P}$ i.e. when the configuration ``jumps'' over the energy barrier to reach another stable configuration. For small positive $\delta\bar{P}$ analytical relations were found: $E_{\rm rel} \propto (\delta\overline{P})^{2.5}$. The energy release $E_{\rm rel}\sim 10^{51}-10^{52}$ erg is an absolute upper bound on the energy which can released in such PT. In astrophysical situation, the energy can be distributed between stellar pulsations, gravitational radiation, heating of stellar interior, X-ray emission from the neutron star surface, and even a gamma-ray burst. Currently there is no mathematical proof of the $J$-independence of $E_{\rm rel}$. It may be interesting to look at the problem from the ``thermodynamical'' point of view. We write the total energy of the star (gravitational mass $M$) as \be \label{e:M_A_J_p} M = \frac{\mu}{u^t} A + 2 \, \Omega J + 2 \int_{\Sigma_t} P\, N\sqrt{\gamma}\, d^3x , \ee where $\mu$ is the baryon chemical potential, $u^t$ the time component of the fluid 4-velocity, $A$ the number of baryons in the star, $\Omega$ the angular velocity, $P$ the fluid pressure, $N$ the lapse function and $\sqrt{\gamma}\, d^3x$ the covariant volume element in the constant $t$ hypersurface $\Sigma_t$. Eq.~(\ref{e:M_A_J_p}) is valid for any axi-symmetric stationary and rigidly rotating fluid star which obeys a barotropic EOS, as established by Bardeen \& Wagoner in 1971 \cite{BardeW71}. Since $\mathcal{C}$ and $\mathcal{C}^*$ have the same baryon number $A$ and the same angular momentum $J$, we get the energy release \be \label{e:DE} E_{\rm rel} = \Delta E = A \, \Delta\! \left( \frac{\mu}{u^t} \right) + 2\, J \, \Delta\Omega + 2 \left[ \int_{\mathcal{C}} P\, N\sqrt{\gamma}\, d^3x - \int_{\mathcal{C}^*} P\, N\sqrt{\gamma}\, d^3x \right] , \ee with \be \Delta\! \left( \frac{\mu}{u^t} \right) := \left. \frac{\mu}{u^t} \right| _{\mathcal{C}} - \left. \frac{\mu}{u^t} \right| _{\mathcal{C}^*} ,~~~ \Delta \Omega := \left. \Omega \right| _{\mathcal{C}} - \left. \Omega \right| _{\mathcal{C}^*} . \ee The terms on the right-hand-side of Eq.~\ref{e:DE} are of comparable magnitude in their contribution to the energy release, so the possible ``smallness'' of some of them in relation to the others is not responsible for the $E_{\rm rel}(J)$ independence. We will address the problem of mathematically proving this property in the near future. | 7 | 10 | 0710.1528 |
0710 | 0710.2237_arXiv.txt | {A wealth of observations of CO in absorption in diffuse clouds has accumulated in the past decade at $uv$ and mm-wavelengths} {Our aims are threefold: a) To compare the $uv$ and mm-wave results; b) to interpret \coth\ and \cotw\ abundances in terms of the physical processes which separately and jointly determine them; c) to interpret observed J=1-0 rotational excitation and line brightness in terms of ambient gas properties.} {A simple phenomenological model of CO formation as the immediate descendant of quiescently-recombining \hcop\ is used to study the accumulation, fractionation and rotational excitation of CO in more explicit and detailed models of \HH-bearing diffuse/H I clouds} {The variation of N(CO) with N(\HH) is explained by quiescent recombination of a steady fraction n(\hcop)/n(\HH) = $2\times 10^{-9}$. Observed N(\cotw))/N(\coth) ratios generally do not require a special chemistry but result from competing processes and do not provide much insight into the local gas properties, especially the temperature. J=1-0 CO line brightnesses directly represent N(CO), not N(\HH), so the CO-\HH\ conversion factor varies widely; it attains typical values at N(\cotw) $\la 10^{16}\pcc$. Models of CO rotational excitation account for the line brightnesses and CO-\HH conversion factors but readily reproduce the observed excitation temperatures and optical depths of the rotational transitions only if excitation by H-atoms is weak -- as seems to be the case for the very most recent calculations of these excitation rates.} { Mm-wave and $uv$ results generally agree well but the former show somewhat more enhancement of \THC\ in \coth . In any case, fractionation may seriously bias \TWC/\THC\ ratios measured in CO and other co-spatial molecules. Complete C$\rightarrow$CO conversion must occur over a very narrow range of \AV\ and N(\HH) just beyond the diffuse regime. For N(\HH) $< 7\times10^{19}\pcc$ the character of the chemistry changes inasmuch as CH is generally undetected while CO suffers no such break. } | 7 | 10 | 0710.2237 |
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0710 | 0710.2371_arXiv.txt | We analyze \WMAP 3 year data using the one-point distribution functions to probe the non-Gaussianity in the Cosmic Microwave Background (CMB) Anisotropy data. Computer simulations are performed to determine the uncertainties of the results. We report the non-Gaussianity parameter $\fnl$ is constrained to $26<\fnl<82$ for Q-band, $12<\fnl<67$ for V-band, $7<\fnl<64$ for W-band and $23<\fnl<75$ for Q+V+W combined data at 95\% confidence level (CL). | Non-Gaussianity is one of the most important tests of models of the inflation. Among the various theoretical models on the inflation, slow-roll inflation is currently most lively being studied. There are various predictions on the magnitude of non-Gaussianity based on the simple model of slow-roll inflation and its extensions, ranging from undetectably tiny values to large enough values to be detectable with currently available data \cite{barnabycline,battefeldeasther,calcagni,creminelli,bartolo.et.al}. On the other hand, observational works have claimed both detection and non-detection of non-Gaussianity (for reviews on recent works, see \cite{komatsu.et.al,Spergel.et.al,troia.et.al,creminelli.et.al,gaztanagawagg}). Among the popular techniques for detecting non-Gaussianity are one-point distribution function fitting, bispectrum, trispectrum and Minkowski functionals. Here, we investigate the one-point distribution functions to probe primordial non-Gaussianity in the CMB anisotropy data. An observed CMB anisotropy at a direction ($\delta T_{obs}$) can be regarded as the superposition of three parts: physical fluctuation of cosmic origin ($\delta T_p$), instrumental noise ($\delta T_n$), and foreground emissions ($T_{fg}$). Since the foreground templates are separately prepared, we start with foreground-removed data of which the CMB anisotropy can be decomposed into two uncorrelated components, \begin{equation}\label{eq1} \delta T=\delta T_{obs}-T_{fg}=\delta T_{p}+\delta T_n. \end{equation} The primary source for the cosmic fluctuation of CMB at the large scale is attributed to the Sachs-Wolfe effect which is again triggered by the primordial curvature perturbation. The curvature perturbation $\Phi$ by primordial seed during the inflation is transferred to CMB anisotropy with the relation \begin{equation}\label{eq2} \frac{\delta T_p\fpr{\bfx}}{T_0}=\eta_{t}\Phi\fpr{\bfx} \end{equation} where $T_0=2.725$ K, the thermodynamic temperature of the CMB today, and $\eta_{t}$ is the radiation transfer function. For the super-horizon scale, we take $\eta_{t}=-1/3$ from the Sachs-Wolfe effects. At the first-order of perturbation, we may replace $\Phi =\Phi_g$, where $\Phi_g$ is an auxiliary random Gaussian field with its mean $\fang{\Phi_g}=0$ and its variance denoted by $\fang{\Phi_g^2}$. When the second-order perturbation is considered, it is conventional to prescribe the nonlinear coupling of the curvature perturbation as \cite{komatsuspergel} \begin{equation}\label{eq3} \Phi (\bfx )\simeq\Phi_g(\bfx )+\fnl\fpr{\Phi_g^2(\bfx )-\fang{\Phi_g^2}} \end{equation} where $\fnl$ is the non-Gaussianity parameter. The second term in \eqref{eq3} is responsible for the non-Gaussianity of the primordial fluctuation. Then, the probability distribution function of the non-Gaussian field $\Phi$ can be derived as \begin{eqnarray} f_{\Phi}(\Phi)&=&\int f_G(\Phi_g)\delta_D \fsq{\Phi-\Phi_g-\fnl\fpr{\Phi_g^2-\fang{\Phi_g^2}}}d\Phi_g\nonumber\\ &=&\frac{1}{\sqrt{2\pi\fang{\Phi_g^2}\fnl^2\fpr{\Phi_+-\Phi_-}^2}}\nonumber\\ & &\times\fsq{\exp\fpr{-\frac{\Phi_+^2}{2\fang{\Phi_g^2}}} +\exp\fpr{-\frac{\Phi_-^2}{2\fang{\Phi_g^2}}}}\label{eq4} \end{eqnarray} where $\Phi_{\pm}$ are defined by \begin{equation}\label{eq5} \Phi_{\pm}=\frac{1}{2\fnl} \fsq{-1\pm\sqrt{1+4\fnl\Phi+4\fnl^2\fang{\Phi_g^2}}} \end{equation} and $\Phi$ has to be limited by the reality of $\Phi_{\pm}$ as \begin{equation}\label{eq6} \fnl\Phi > -\frac{1}{4}-\fnl^2\fang{\Phi_g^2}. \end{equation} $\fang{\Phi_g^2}$ can be expressed in terms of $\eta_t$, $T_0$ and $\sigma_{CMB}$, \begin{equation}\label{eq7} \fang{\Phi_g^2}=\frac{1}{4\fnl^2}\fsq{-1+\sqrt{1+8 \fpr{\frac{\fnl\sigma_{CMB}}{\eta_t T_0}}^2}}. \end{equation} For a pixelized CMB anisotropy data set, the probability distribution function for Gaussian instrumental noise becomes \begin{equation}\label{eq8} f_N(\delta T_n)=\frac{1}{N_{pix}}\sum_{i=1}^{N_{pix}} \frac{1}{\sqrt{2\pi\sigma_0^2/n_i}} \exp\fsq{-\frac{\delta T_n^2}{2\sigma_0^2/n_i}} \end{equation} where $n_i$ is the effective number of measurements at the $i_{th}$ pixel and $\sigma_0$ represents the dispersion of the instrumental noise per observation ($\sigma_0$=2.1898, 3.1249, 6.5112 mK for Q, V, W-band, respectively \cite{Limon.et.al}). Now, it is straightforward to express the probability density function for $\delta T$ in an integral form, \begin{eqnarray} f(\delta T)&=&\int f_{\delta T_p}(\delta T_p)f_N(\delta T_n)\nonumber\\ & &\times\delta_D\fpr{\delta T-\delta T_p-\delta T_n} d\delta T_p d\delta T_n\nonumber\\ &=&\int f_{\Phi}(\Phi )f_N(\delta T_n)\nonumber\\ & &\times\delta_D\fpr{\delta T-\eta_t T_0\Phi-\delta T_n}d\Phi d\delta T_n. \label{eq9} \end{eqnarray} The probability density function derived in \eqref{eq9} explicitly contains the non-Gaussianity parameter $\fnl$, and it can serve as the prediction of one-point distribution function with a given $\fnl$ for a (ideally) foreground-removed CMB anisotropy data set to estimate the magnitude of deviation from Gaussian distribution in a quantitative manner. | We developed an algorithm that uses the one-point distribution function to investigate the non-Gaussianity of CMB anisotropy data, and applied it to \WMAP 3 year data. We found that the null result ($\fnl$=0) is manifestly excluded at 95\% CL. The estimated magnitude of non-Gaussianity parameter is $23<\fnl<75$ at 95\% CL and $9<\fnl<88$ at 99\% CL for the (Q+V+W)-combined map. Since the quadratic term in \eqref{eq3} takes a generic form of Taylor series for a perturbative expansion, it is a good possibility that the observed non-Gaussianity in this work is a combined effects of various physical processes, while the primordial seeds are very likely to be the leading one. There are two premises we have taken in developing the algorithm, which, provided they are not precise enough, could cause non-Gaussianity of not cosmic but systematic origin: (1) the probability distribution function of the instrumental noise for each pixel is centered at zero, and (2) the foreground emissions are removed efficiently enough in the foreground-removed maps. The first condition can be broken when the thermal and radiation environments of the \WMAP satellite in its orbit are taken into account, while the \WMAP team assessed they are insufficient to influence the science data \cite{Limon.et.al}. So, we tested the effects of the alternative noise distributions with a random mean in each of the Gaussian distribution in \eqref{eq8} and the algorithm was not misled to show non-Gaussianity within the statistical error. It is difficult to directly estimate how much residual foreground emissions after foreground subtraction would affect the one-point distribution function. We solely rely on the quality of foreground templates and it is remarkably successful, showing that the observed total Galactic emission matches the model to less than 1\% \cite{bennett.et.al2,Hinshaw.et.al}. We also analyzed simulated maps which are (Gaussian map + foreground templates), and all the templates for Q, V and W-channel showed negative values of the non-Gaussianity parameter with $\fabs{\fnl}\sim\mathcal{O}\fpr{10^1}$ at the resolution $N_{\mathrm{side}}=512$. | 7 | 10 | 0710.2371 |
0710 | 0710.1552_arXiv.txt | The accreting millisecond pulsar SAX J1808.4-3658 may be a transition object between accreting X-ray binaries and millisecond radio pulsars. We have constrained the thermal radiation from its surface through XMM-Newton X-ray observations, providing strong evidence for neutrino cooling processes from the neutron star core. We have also undertaken simultaneous X-ray and optical (Gemini) observations, shedding light on whether the strong heating of the companion star in quiescence may be due to X-ray irradiation, or to a radio pulsar turning on when accretion stops. | 7 | 10 | 0710.1552 |
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0710 | 0710.0762_arXiv.txt | We use low-degree acoustic modes obtained by the BiSON to estimate the main-sequence age $t_\odot$ of the Sun. The calibration is accomplished by linearizing the deviations from a standard solar model the seismic frequencies of which are close to those of the Sun. Formally, we obtain the preliminary value $t_\odot=4.68\pm0.02\,$Gy, coupled with an initial heavy-element abundance $Z=0.0169\pm0.0005$. The quoted standard errors, which are not independent, are upper bounds implied under the assumption that the standard errors in the observed frequencies are independent. | Seismological calibration of stellar models against observed frequencies of low-degree modes was first discussed more than two decades ago \citep{jcd84, jcd88, ulrich86, dog87}, and can be regarded as a means of determining the main-sequence age of the Sun \citep{guenther89, dog_nov90, guenther_demarque97, weiss_schlattl98, wd99, dog01, bon_schlat_pat02}. The procedure is to match certain appropriate seismic signatures of theoretical frequencies determined on a grid of stellar models with corresponding signatures obtained from the observations. The signatures are chosen to reflect principally the properties of the energy-generating core, where nuclear transmutation leaves behind an augmenting concentration of helium, lowering the sound speed relative to the environs and thereby providing a diagnostic of age. But the signatures are also susceptible to other properties of the stellar interior, which must be eliminated before a robust outcome can be achieved. For example, although the so-called small frequency separation is sensitive predominantly to the evolving stratification of the core, its dependence on the zero-age chemical abundances plays a significant contaminating role. Therefore it behoves us to seek an additional diagnostic to attempt to measure abundance separately. For given relative abundances of the heavy elements, the total absolute heavy-element abundance $Z$ and the $^4\rm{He}$ abundance $Y$ are related by the requirement that the model has the observed luminosity and radius, principally the former. Therefore we need to % aim at detecting only one of them. Here we use a signature indicative of the abrupt variation of the first adiabatic exponent $\gamma_1$ induced by the ionization of helium. \begin{figure} \includegraphics[height=.42\textheight]{hg_fig1.eps} \caption{Top left: The symbols (with error bars obtained under the assumption that the raw frequency errors are independent) represent second differences, $\Delta_2\nu$, of low-degree solar frequencies from BiSON. Top right: The symbols are second differences $\Delta_2\nu$ of adiabatic pulsation eigenfrequencies of solar Model S. The solid curve in both panels is the diagnostic\, (\ref{eq:delnu}) -- (\ref{eq:secdiff}), whose eleven parameters have been adjusted to fit the data optimally. Bottom: The symbols denote contributions $\delta\nu$ to the frequencies produced by the acoustic glitches of the Sun (left panel) and Model S (right panel).} \end{figure} | 7 | 10 | 0710.0762 |
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0710 | 0710.2290_arXiv.txt | We study the possibility to extract the multipolar moments of an underlying distribution from a set of cosmic rays observed with non-uniform or even partial sky coverage. We show that if the degree is assumed to be upper bounded by $L$, each multipolar moment can be recovered whatever the coverage, but with a variance increasing exponentially with the bound $L$ if the coverage is zero somewhere. Despite this limitation, we show the possibility to test predictions of a model without any assumption on $L$ by building an estimate of the covariance matrix seen through the exposure function. | Anisotropy in the arrival directions of cosmic rays is a major observable to understand their origin. Magnetic fields bend their trajectories in such a way that transport of cosmic rays is mainly diffusive up to high energies: this makes their angular distribution isotropic. Nevertheless, above the so-called knee of cosmic rays up to the ankle, there are predictions for small but increasing anisotropies with energy, predictions which of course depend on the regular and the turbulent components of the assumed galactic magnetic field, as well as the assumed distribution of sources and composition of cosmic rays \cite{ptuskin,candia}. Further, at ultra-high energies, cosmic ray arrival directions are expected to be less and less smeared out by galactic and extragalactic magnetic fields, leading to a possible extraction of informations about the position of the sources \cite{isola,sigl,eric,dolag,demarco}. Hence, it is clear that any evidence for an anisotropy, or any limit on anisotropies in the cosmic ray locations observed by experiments are among the most important constraints upon models. The multipole expansion up to a given order $L$ is a powerful tool to study the structures standing out the noise down to an angular scale $\approx \pi/L$, whatever the shape of the underlying celestial pattern. In practice, the number of significant coefficients is limited by the angular resolution of the detector and, in the other hand, by the available statistics of observation. However, ground based experiments cover a limited range in declination, so that it is impossible to apply off the shelf the formalism of multipole moments: anyone of the coefficients may be modified in an unpredictable way by the unseen part of the sky. Methods have been developped to study the CMB with an incomplete coverage \cite{gorski,wright,tegmark,mortlock}, but here we are faced to a different problem: we cannot suppose a priori that the distribution of cosmic rays is described by a power spectrum, because we want to detect possible non-isotropic structures, a priori unknown. In other terms, the information carried by the $a_{\ell m}$ cannot be reduced to the only knowledge of the $C_{\ell}$. One purpose of this paper is to study the possibility of estimating the multipole moments of a distribution of points over a sphere in case of a non-uniform or even a partial coverage of the sky, together with the limitations of such an approach. The estimation of dipoles and quadrupoles was studied in \cite{sommers,julien,silvia}. Here, we use the moments of the observed distribution on a set of orthogonal functions: either the spherical harmonics themselves, or a set of functions tailored on the coverage function. With these two different methods, we show that the interference between the modes induced by the the non-uniformity or the hole of the coverage can be removed assuming a bounded expansion in the conjugate space, allowing to recover the underlying multipole moments. However, in accordance with the simple intuition that it is impossible to describe the unseen part of the sky, we point out that the uncertainty on the recovered coefficients increases with the assumed bound $L$ of the expansion. We show that the larger the hole in the coverage of the sky, the faster the increase of uncertainty with $L$. After some general considerations about the description of point processes on a sphere in Section 2, Sections 3 and 4 are dedicated to these methods whereas Section 5 illustrates them with some examples. Because of the incomplete knowledge of the distribution of cosmic ray sources, and the stochastic nature of the propagation through magnetic fields, the anisotropies we want to characterize are not reducible to explicit models: they may be interpreted as a particular realization of a random process. This means that some model predictions are better expressed as average values of the coefficients, with their covariance matrix. This matrix is not necessarily diagonal to describe the physics we are interested in, contrary to the case of a power spectrum. We show in Sect.6 that under reasonable assumptions, an estimate can be performed with a partial sky coverage, evading the problem of setting a bound to the expansion. | To cope with a partial sky coverage, a formalism using the computation of moments on orthogonal functions was developed to recover the angular distribution of the incident flux from a sample of $N$ observed points. If the multipolar expansion is assumed to be upper bounded by $\ell\leq L$, the coefficients $a_{\ell m}$ may be estimated with a variance proportional to $1/N$ as usually, with a penalty factor increasing exponentially with $L$ if there is a hole in the coverage (but stabilizing rapidly if the coverage is nowhere vanishing, even highly non-uniform). Two methods were tested, giving similar results, and practically the same variances. Statistical tests based on likelihood ratios may be built to check an hypothesis on the distribution, for example a given bound $\ell \leq L$. In any case, it is possible to express predictions of a model in terms of coefficients which can be computed without any assumption on $L$, and tested against the moments found with a sample of observed points. The methods presented in this paper may be applied any cosmic ray dataset, provided that the arrival directions and the coverage of the sky are known within a reasonable precision. | 7 | 10 | 0710.2290 |
0710 | 0710.0540_arXiv.txt | We present the discovery of a silicate disc at the centre of the planetary nebula Mz3 (the Ant). The nebula was observed with MIDI on the Very Large Telescope Interferometer (VLTI). The visibilities obtained at different orientations clearly indicate the presence of a dusty, nearly edge-on disc in the heart of the nebula. An amorphous silicate absorption feature is clearly seen in our mid-IR spectrum and visibility curves. We used radiative transfer Monte Carlo simulations to constrain the geometrical and physical parameters of the disc. We derive an inner radius of 9 AU ($\sim$6mas assuming D=1.4kpc). This disc is perpendicular to, but a factor of $10^{3}$ smaller than the optical bipolar outflow. | \label{sec:1} Complex phenomena perturb the mass-ejection in the late stages of stellar evolution: stellar magnetic fields, rotation or binarity are often invoked. Those mechanisms can lead to the creation of a circumstellar, dusty disc. The dust is created in the outer parts of the former stellar envelope. The Ant planetary nebula, or Mz3 is a bipolar nebula that has one of the most ``pinched'' waists. Its central object is suspected to be a binary but there are no stringent constraints. The temperature of the ionization source is about 30,000K and its distance is between $1-2$ kpc. Different expansion phases have shaped the current form of the nebula as seen in the Hubble Space Telescope images and long-slit spectroscopy \cite{gue,san}. Former studies \cite{smi1,smi2} have proposed the existence of a circumstellar disc, but due to its small size it could only be observed with high-angular resolution interferometry and only in the infrared. | 7 | 10 | 0710.0540 |
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0710 | 0710.2635_arXiv.txt | We have mapped the proto-binary source IRAS 16293--2422 in CO 2--1, $^{13}$CO 2--1, and CO 3--2 with the Submillimeter Array (SMA). The maps with resolution of 1\farcs5--5\arcsec\ reveal a single small scale ($\sim$3000 AU) bipolar molecular outflow along the east-west direction. We found that the blueshifted emission of this small scale outflow mainly extends to the east and the redshifted emission to the west from the position of IRAS 16293A. A comparison with the morphology of the large scale outflows previously observed by single-dish telescopes at millimeter wavelengths suggests that the small scale outflow may be the inner part of the large scale ($\sim$15000 AU) E--W outflow. On the other hand, there is no clear counterpart of the large scale NE--SW outflow in our SMA maps. Comparing analytical models to the data suggests that the morphology and kinematics of the small scale outflow can be explained by a wide-angle wind with an inclination angle of $\sim$30\degr--40\degr~with respect to the plane of the sky. The high resolution CO maps show that there are two compact, bright spots in the blueshifted velocity range. An LVG analysis shows that the one located 1\arcsec\ to the east of source A is extremely dense, n(H$_2$)~$\sim$10$^{7}$ cm$^{-3}$, and warm, T$_{\rm kin} >$~55 K. The other one located 1\arcsec\ southeast of source B has a higher temperature of T$_{\rm kin} >$~65 K but slightly lower density of n(H$_2$)~$\sim$10$^{6}$ cm$^{-3}$. It is likely that these bright spots are associated with the hot core-like emission observed toward IRAS 16293. Since both two bright spots are blueshifted from the systemic velocity and are offset from the protostellar positions, they are likely formed by shocks. | Outflow phenomena have been recognized as an important phase in the star formation processes. It is generally accepted that when stars are formed by gravitational infall, outflows transfer excess angular momentum out of the system \citep[e.g.,][]{shu1987}. Outflows are observed at various wavelengths from the radio to the ultraviolet, among which the millimeter/submillimeter bands probe the molecular outflows. Molecular outflows, recognized as high-velocity wings in CO and other molecular lines, are considered to be the ambient gas entrained or pushed by protostellar winds. By observing molecular outflows, we can determine their physical, kinematic and chemical properties, trace the mass-loss history of protostars, and probe the early phases of the star formation process (\citealt{arce2007}, and references therein). IRAS 16293--2422 (hereafter referred to as I16293) is a well-studied proto-binary system located in the nearby molecular cloud complex L 1689. The distance to L 1689 is often assumed to be 160 pc \citep{whittet1974}, however some recent studies suggested a distance of $\sim$120 pc \citep{degeus1989, knude1998}. In this paper we adopt 120 pc as the distance. The projected separation of the two protostars is $\sim$5\arcsec~\citep{mundy1992}, which then corresponds to 600 AU. I16293 is classified as a Class 0 young stellar object (YSO) since this system is undetectable at wavelengths shorter than 10 $\mu$m and shows a spectral energy distribution (SED) with very low temperature \citep{andre1994}. The bolometric luminosity of I16293 is estimated to be $\sim$36 $L_\odot$ \citep{correia2004}. I16293 is known to have a quadrupolar outflow \citep{walker1988, mizuno1990}; one bipolar pair extends along the east (blue) and west (red) (hereafter the E--W pair) directions, while the other one has an orientation from northeast (red) to southwest (blue) (hereafter the NE--SW pair) with an inclination angle of 30$^\circ$--45$^\circ$ \citep{hirano2001} or 65$^\circ$ \citep{stark2004}. The NE--SW pair is more collimated, and its axis crosses the position close to the SE continuum source IRAS 16293A (hereafter source A). Therefore this pair has been thought to be powered by source A \citep[e.g.,][]{walker1988, mundy1990}. On the other hand, the E--W pair is not as well collimated and its axis, though not well defined, seems to pass to the north of source A, where the continuum source IRAS 16293B (hereafter source B) is the only possible driving YSO. Hence source B has often been assumed as the driving source of the E--W pair. \citet{stark2004} showed that source B has very narrow line widths, a low luminosity, and is not evidently associated with high-velocity gas. They interpreted source B as a T--Tauri star that drove the E--W pair and that the E--W pair is now a fossil flow. However, these interpretations are based on observations at low angular resolution (i.e., $>$5\arcsec), that are insufficient to resolve the two protostars and reveal the relation between the quadrupolar outflow and the binary. In fact, using higher angular resolution observations \citet{chandler2005} proposed an alternative interpretation. They resolved source A into four components, two at 1~mm and another two at 7~mm, and suggested that two of them are responsible for driving the NE--SW and the E--W pair, respectively. In addition, \citet{chandler2005} proposed that source B is actually much younger than previously interpreted and may have not yet begun the phase of mass loss. Their proposition was, however, made based on (1) the dust continuum morphology and proper motions of source A, and (2) the velocity structures of emission of H$_2$CO, H$_2$S, and SO, rather than CO, which is a much better tracer of the overall outflow structure free from chemical peculiarities and shocks. In order to study the molecular outflows in the vicinity of I16293 in more detail, we have carried out CO 2--1, $^{13}$CO 2--1, and CO 3--2 observations at higher angular resolutions using the Submillimeter Array\footnote{The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} (SMA). Our SMA observations provide sufficient angular and spectral resolution in the CO lines to resolve the kinematics and the structure of the outflow in the central region of I16293. In this paper, we summarize the observational details in \S2; present our SMA CO 2--1, $^{13}$CO 2--1, and CO 3--2 results in \S3; discuss the outflow driving mechanism, physical parameters of the outflow, and the kinematics in the vicinity of I16293 in \S4. | We have carried out CO 2--1 $^{13}$CO 2--1, and CO 3--2 observations of the proto-binary system IRAS 16293-2422 with a resolution of 1\farcs5--5\arcsec~using the Submillimeter Array (SMA), which are about a factor of 5 better than previous observations in these lines. Our observations reveal the detailed structures of molecular outflows close to the binary system. Our main results are summarized as follows: 1. The high resolution images of CO 2--1 and CO 3--2 show a compact bipolar outflow along the east-west direction in the vicinity of the binary. The center of this small scale outflow is close to source A, suggesting that source A is the driving source. If we interpret this small scale outflow as the inner part of the large scale E--W pair of the quadrupolar outflow, the E--W pair is a currently active outflow rather than a fossil flow. On the other hand, there is no clear counterpart of the large scale NE--SW pair in our interferometric maps. 2. The shape and velocity structure of the western redshifted lobe are well described by an analytical model suggesting that the small scale E--W pair could be driven by a wide-angle wind with an inclination angle of $\sim$30$^\circ$--40$^\circ$. 3. There are two compact and bright spots at blueshifted velocities. One is located 1\arcsec\ east of source A (labeled b1) and the other is 1\arcsec\ southeast of source B (labeled b2). An LVG analysis shows that b1 is extremely dense, n(H$_2$)~$\sim$10$^{7}$ cm$^{-3}$ and warm T$_{\rm kin} >$~55 K, while component b2 has a higher temperature of T$_{\rm kin} >$~65 K but slightly lower density of n(H$_2$)~$\sim$10$^{6}$ cm$^{-3}$. It is likely that these bright spots are produced by means of shocks and are associated with hot core-like activity observed toward both source A and source B. | 7 | 10 | 0710.2635 |
0710 | 0710.0630_arXiv.txt | Some aspects of disk-halo interactions for models of in and out of equilibrium disk galaxies are reviewed. Specifically, we focus on disk-halo resonant interaction without and in the presence of a gas component. Another issue is the disk growth within an assembling triaxial dark matter halo. We argue that while the triaxiality is the result of the merger process and the radial orbit instability, it is the developing chaos that damps the first generation of bars and washes out the halo prolateness. This chaos is triggered by the gravitational quadrupole interaction(s) in the system and supported by a number of other processes which are characteristic of baryons. | The current paradigm of galaxy formation necessitates that galactic disks grow and evolve within dark matter (DM) halos. As such, the baryonic disks serve as a test bed for studying the DM properties, its dynamics and morphology, as well as mass, momentum and energy exchange within the disk-halo system. Interaction between the galactic disks and DM halos can in principle involve an exchange of these quantities. For example, observations of the low angular momentum H\,I gas, deep in the halo of NGC~891, require cold extragalactic gas influx (e.g., Fraternali et al. 2007), maybe through filaments (Dekel \& Birnboim 2006). On the other hand, the X-ray halos of starbursts and some normal galaxies can be explained by the supernovae-heated gas driven from the disk or by the shock-compressed gas in the halo (e.g., Strickland et al. 2004). The situation is much more straightforward with the angular momentum ($J$) flow in the system. Galactic disks are rotationally supported, while the DM halos have low $J$, and, therefore, a low spin parameter $\lambda$ (e.g., Barnes \& Efstathiou 1987; Frenk, White \& Efstathiou 1988). Most of the disk galaxies are barred, especially in the NIR (e.g., Knapen, Shlosman \& Peletier 2000), with bar properties which remain steady up to the redshift of $\sim 1$ at least (Jogee et al. 2004; Elmegreen, Elmegreen \& Hirst 2004), and except for the very early Hubble types, they exhibit a spiral structure. This prevailing {\it disk asymmetry} is extremely important for enhancing the $J$ transfer between different morphological components --- the alternative way can be achieved, e.g., by dynamical friction of baryons against the DM. The current understanding of galactic bar formation relies on the spontaneous breakup of the axial symmetry in the disk, the so-called classical bar instability (e.g., Hohl 1971), originally applied to isolated disks. Paradoxically, while Lynden-Bell \& Kalnajs (1972) have shown that bars and spirals facilitate the $J$ transfer, DM halos have been considered as the main stabilizers against the bar instability for years (Ostriker \& Peebles 1973), though they appear rather to be sinks of the disk momentum (e.g., Athanassoula \& Misiriotis 2002). Within the framework of the classical bar instability, a stellar bar can develop only if the unstable region loses its $J$. Gravitational torques serve as the mechanism that drives this process (Lynden-Bell \& Pringle 1974). Their action can be described in terms of a non-local viscosity (Lin \& Pringle 1987; Shlosman 1991) which shortens the characteristic timescale of the $J$ transfer dramatically, making this a dynamical rather then secular process. In principle, $J$ can flow across the corotation radius (CR) to the outer disk, to the DM halo, or to a flyby galaxy in case of a tidal galaxy interaction. The first option limits the $J$ flow because the mass of the outer disk is typically $\sim 20\%$ of the disk mass and its ability to absorb $J$ quickly saturates (Fig.~1, left). On the other hand, the halo is massive, its inner part has a comparable mass to the inner disk, and its overall $J$ is small --- the halo is not supported by rotation. Fig.~1 shows also that the outer halo, beyond the disk radius, is fully susceptible to the $J$ transfer, after the inner halo efficiency has decreased. \begin{figure}[!t] \plotfiddle{fig01a.ps}{3.5cm}{-90}{45}{45}{-170}{140} \plotfiddle{fig01b.ps}{3.5cm}{0}{31}{31}{-10}{-16.5} \caption{ \underline{\it Left:} Evolution of angular momentum, $J$, in the disk (upper panel) and the halo (lower panel). Note, $J$ saturates in the disk outside the bar CR (from Martinez-Valpuesta, Shlosman \& Heller 2006). \underline{\it Right:} The $m=2$ amplitude, $A_2$, of the stellar and DM ghost bars (upper panel) evolution. The fast dissolution of the ghost bar (lower) after its stellar counterpart is axisymmetrized. Note the very short time period for the lower frame (I. Berentzen \& I. Shlosman, unpublished)} \label{fig1} \end{figure} The relatively local density response in the DM halo to the growing bar in the disk is to generate a shadow `ghost' bar in the DM (Athanassoula 2006, 2007; Berentzen \& Shlosman 2006). The ghost has the pattern speed of the stellar bar but its mass distribution differs and is more centrally concentrated. The particle orbits in the DM ghost, therefore, are in resonance with the stellar orbits. The principal difference between these two bars is that the DM bar is not strictly speaking a bar at all --- it is not self-gravitating {\it per se} and represents a gravitational wake induced in the DM by its stellar counterpart. When the stellar bar is axisymmetrized, the DM bar dissolves in a fraction of a crossing time (Fig.~1, right). | We have summarized some aspects of disk-halo interactions in equilibrium and forming disk galaxies. Because the overall problem is of a high complexity, progress in this direction has been rather slow and dependent on numerical simulations. Nevertheless, certain trends have emerged. The baryonic disk and DM halo exchange mass, angular momentum and energy during the initial formation phase as well as during the subsequent quiescent evolution. A number of processes, like star formation as well as formation of the central SBH, require a much deeper understanding before they can be modeled beyond the phenomenological approach. On the other hand, nonlinear dynamics and the role of chaos in the evolution of disks and halos can be quantified already at present. | 7 | 10 | 0710.0630 |
0710 | 0710.4964_arXiv.txt | We write the averaged Einstein equations in a form suitable for use with Newtonian gauge linear perturbation theory and track the size of the modifications to standard Robertson-Walker evolution on the largest scales as a function of redshift for both Einstein de-Sitter and $\Lambda$CDM cosmologies. In both cases the effective energy density arising from linear perturbations is of the order of $10^{-5}$ times the matter density, as would be expected, with an effective equation of state $w_\mathrm{eff}\approx -1/19$. Employing a modified Halofit code to extend our results to quasilinear scales, we find that, while larger, the deviations from Robertson-Walker behaviour remain of the order of $10^{-5}$. | \label{intro} Observations of the cosmic microwave background (CMB) \cite{WMAP,WMAP(2),WMAP3}, large-scale structure (LSS) (e.g. \cite{2df,SDSS}) and supernovae type Ia data \cite{RiessEtAl98,PerlmutterEtAl98,RiessEtAl06,WoodVaseyEtAl07} consistently indicate the presence in the universe of significant ``dark'' components. Data from the LSS and from nucleosynthesis indicates that the density of baryonic matter should not exceed about 5\% of the critical density, while data from the CMB suggests that the universe should be flat, implying that around 95\% of the matter-energy content of the universe is unobserved. In combination with numerical studies (e.g. \cite{SpringelEtAl05}) and observations of the baryonic acoustic oscillations in the LSS \cite{2df,SDSS} this has lead to the ``concordance model'' in which roughly 5\% of the universe is in the form of baryonic matter, 20\% of the universe is in the form of some ``dark matter'' that interacts only gravitationally, and about 75\% is in the form of a ``dark energy'' fluid with a negative equation of state today $w\approx-1$. The most popular alternatives to a cosmological constant involve exotic fluids that violate the strong energy condition in the late universe, of which scalar-field models such as quintessence are the most popular (e.g. \cite{Wetterich88,RatraPeebles88,CaldwellEtAl97} for pioneering works and \cite{CopelandEtAl06,Linder07} for recent reviews.) There are also many other modifications of both the matter and gravitational sectors. Another approach is to consider the impact of local inhomogeneities on the luminosity distances in the local universe (see for example \cite{EnqvistMattson07,Kasai07,Marra07} for some recent work on this topic). The modifications that the local inhomogeneities introduce can, it has been suggested, account for many of the features exhibited by dark energy. However, the assumptions used to build up our standard model include an implicit averaging procedure, assuming that the averaged Einstein tensor is equivalent to the Einstein tensor built from an average metric -- and such is not the case \cite{Ellis84,Futamase89}. The Einstein equations are non-linear and local; the correct approach is to average the local equations across some domain rather than to assume ab initio ``averaged'' equations. Studies into averaged cosmologies include \cite{Kasai93,Futamase96,BuchertEhlers95,Boersma97,RussEtAl97} and the impact of inhomogeneities were applied to account for a dark energy in \cite{Buchert99}. In the years since, this ``backreaction'' and the impact of inhomogeneities on observables has been relatively well studied in a variety of models and remains an active field \cite{Buchert01,Wetterich01,BuchertCarfora02,Rasanen03,Rasanen04,KolbEtAl04,EllisBuchert05,AlnesEtAl05,Rasanen06,IshibashiWald06,BuchertEtAl06,LarenaBuchert06,LarenaEtAl06,BiswasEtAl06,KasaiEtAl06,VanderveldEtAl06,TanakaEtAl06,Brouzakis06,LiSchwarz07,Enqvist07,Wiltshire07,Mattsson07,Ishak07,VanderveldEtAl07,Wiltshire07_2,KhosraviEtAl07,Hossain07}; see also \cite{BuchertReview07} for a recent status report. The attraction of this approach is that it can recover quantities that -- in principle -- could act as a dark energy, without the need for exotic matter components or modifications to general relativity. Moreover, one can intuitively state that as the universe grows increasingly non-linear, the deviations on averaging from a truly homogeneous and isotropic model should increase, potentially providing an appealing solution to the coincidence problem. The current universe is significantly inhomogeneous at scales of up to 100-300Mpc/$h$ \cite{PercivalEtAl06}. While we might still expect to recover an FLRW-like metric when averaging across scales larger than 100Mpc/$h$ (see \cite{Lu07,McClure07} for recent studies) there is no automatic guarantee that it will obey the familiar Friedmann and Raychaudhuri equations. It has been shown in many studies that structure -- both at the linear order and for highly inhomogeneous models -- can introduce modifications to the large scale effective Friedmann equations (though this is still somewhat debated; e.g. \cite{AlnesEtAl05_2,KasaiEtAl06,TanakaEtAl06,EnqvistMattson07}). Even should the effect be relatively minor, with current and future experiments a divergence from Robertson-Walker behaviour on the order of $10^{-3}$ could be significant. In this paper we calculate numerically as a function of redshift the size of the deviations from standard Robertson-Walker behaviour arising from linear and quasilinear perturbations. Such a calculation complements the recent studies by \cite{VanderveldEtAl07}, in which the authors reconstruct the impact of backreaction effects from the observational data, and \cite{KhosraviEtAl07} wherein the authors evaluate the size of the effective density of the backreaction as a function of redshift in a structured Robertson-Walker model. Much of the current literature concerns exact inhomogeneous rather than perturbative models, not least because the impact in a perturbative approach is not expected to be large. However, the magnitude and nature of the backreaction and other corrections associated with averaging that arise naturally within a standard Einstein de-Sitter or $\Lambda$CDM cosmology remain open questions and this should be addressed. We work in conformal Newtonian gauge, in contrast to much of the literature which concerns itself with synchronous gauge, for two main reasons. Firstly Newtonian gauge, unlike synchronous gauge, is well-defined and should yield unambiguous conclusions; and, moreover, the use of Newtonian gauge enables us to employ the easily extensible cmbeasy \cite{Doran03} code to track the size of the perturbations. We find that linear perturbations for a model Einstein de-Sitter universe introduce an effective energy density $\rho_\mathrm{eff}\approx (4\times 10^{-5})\bkr_m$ with an equation of state $p_\mathrm{eff}/\rho_\mathrm{eff}=w_\mathrm{eff}\approx -1/19$, while those for the $\Lambda$CDM concordance model introduce $\rho_\mathrm{eff}\approx (1.3\times 10^{-5})\bkr_m$ with a marginally lower equation of state. Employing a modified Halofit \cite{SmithEtAl02} code to include quasilinear scales increases the effective energy densities to $\rho_\mathrm{eff}\approx (5.6\times 10^{-5})\bkr_m$ and $\rho_\mathrm{eff}\approx(1.6\times 10^{-5})\bkr_m$ for the EdS and $\Lambda$CDM cases respectively, with unaltered equations of state. These are in good agreement with other evaluations of $\rho_\mathrm{eff}$ (e.g. \cite{Wetterich01,Rasanen04,VanderveldEtAl07,KhosraviEtAl07}) which typically remain of the same order. We begin with a brief overview of the 3+1 formalism of general relativity and in \S\ref{Averaging} briefly discuss averages in cosmology before applying a simple average in \S\ref{Averages}. In \S\ref{NewtGauge} we calculate the forms of the backreaction terms for conformal Newtonian gauge and in \S\ref{Backfast} we present the deviations from an FLRW model from linear perturbations employing a modified Boltzmann code for both Einstein de-Sitter and $\Lambda$CDM universes, making it more general than the similar studies in \cite{Rasanen04,KasaiEtAl06}, who did not consider $\Lambda$CDM universes, and \cite{TanakaEtAl06}, who worked in a rather narrow range of domain sizes. We follow by estimating the impact from quasilinear scales in \S\ref{Halofit}. | We have presented quantitive estimates of the corrections in Newtonian gauge to a standard perturbative FLRW model from an explicit averaging procedure, which can be separated into four distinct quantities: the kinematic backreaction, which remains strictly negligible across the scales of interest, the dynamic backreaction, the impact of spatial curvature and a correction to the FLRW energy density. In line with expectation, the impact from linear perturbations is insignificant and of the order of $10^{-5}$ for both $\Lambda$CDM and an Einstein de-Sitter model. Specifically, the effective energy density arising at the background level from inhomogeneities is $\rho_\mathrm{eff}\approx (4\times 10^{-5})\bkr_m$ and $\rho_\mathrm{eff}\approx(1.3\times 10^{-5})\bkr_m$ for the EdS and $\Lambda$CDM models respectively. Moreover, we have presented estimates arising from halo model corrections on quasilinear scales which are not much larger than those from the linear perturbations, the largest contribution arising from the quasilinear modes in an EdS universe remaining below $10^{-4}$; we find $\rho_\mathrm{eff}\approx(5.6\times 10^{-5})\bkr_m$ for the EdS model and $\rho_\mathrm{eff}\approx(1.6\times 10^{-5})\bkr_m$ for $\Lambda$CDM. This is in broad agreement with recent calculations from Vanderveld \emph{et. al.} estimating the total possible backreaction from observations of Type Ia supernovae and agrees with a growing consensus that the impact of backreaction across all scales remains at the level of $10^{-5}-10^{-4}$. Additionally, the total effective equation of state arising from the different modifications does not act as a dark energy and instead as a dark matter with a slowly-varying equation of state $w_\mathrm{eff}\approx -1/19$ for $z\in(0,10)$. The impact is to decelerate the universe and, contrary to other studies, the corrections impact positively on the Friedmann equation implying a positive effective pressure. (The discrepancy arises due to our definitions of the effective energy density and pressure of the corrections taking into account the differences in definition of energy density in an FLRW frame and with respect to our foliation.) This does not, however, necessarily imply that there are no observational consequences arising from inhomogeneities, particularly on small scales, as our analysis is limited to the impacts in very large volumes and breaks down when mode-coupling and virialisation become significant. In particular, one can readily imagine situations in which the local Hubble rate is significantly larger than the global average (the so-called ``Hubble bubble''). The impact of local inhomogeneities directly on luminosity distances has also long been well studied and remains an active area of interest. Moreover, the issue of averaging in cosmology and general relativity in particular remains very much an active field. The average we have employed in this study is certainly not perfect; even were we to extend our approach to second-order perturbations it cannot cope with the vector and tensor perturbations that necessarily become significant.% Improving the average and applying it to nonlinear situations is an immediate aim. The prospect for future study remains large. Our model here is necessarily simple and there are many avenues both analytical and numerical that can be pursued. Much interest is currently focusing on the exact swiss-cheese models and particularly on the impact on direct observables from these models. The prospect of directly calculating both the direct backreaction in different domain sizes and the impact on null geodesics (and hence luminosity distances) in different matter configurations from numerical simulations is also an exciting one. More immediately, one can imagine employing second-order perturbation theory in which the ambiguity between averages taken across the assumed zeroth-order manifold and those across the linearly-perturbed manifold may become significant; depending on the surface one defines a vanishing $\langle\phi_{(2)}\rangle$ across, one would expect a non-vanishing contribution on the other. In particular, if one defined second-order perturbations to vanish on average on the background then there would be a residual contribution to the linear perturbations even before any other backreaction is taken into account. This is an interesting issue that deserves more study. More ambitiously, the fully non-linear approaches of Vernizzi and Langlois \cite{LangloisVernizzi05_1,LangloisVernizzi05_2,LangloisVernizzi06_1} and Enqvist \emph{et. al.} \cite{EnqvistEtAl06} might provide an interesting tool for studying the impacts of averaged nonlinear perturbations. The direct backreaction from more complicated systems than our linearly-perturbed FLRW universe might also be considerable; an immediate example, first considered in \cite{Wetterich01}, would be the backreaction from inhomogeneities in a cosmological scalar field which were shown to potentially have an impact of order unity. Such systems might also be considered in a fully non-linear manner \cite{LangloisVernizzi06_2}. In summary, one may comment that backreaction and the impact of inhomogeneities are direct physical effects and should certainly be taken into account in any full treatment of the observables. However, the direct impact when evaluated for the concordance cosmology on very large scales is too small to provide an alternative to the dark energy. | 7 | 10 | 0710.4964 |
0710 | 0710.1770_arXiv.txt | While masers in the 1720~MHz transition of OH are detected toward many supernova remnants (SNRs), no other OH transition is seen as a maser in SNRs. We present a search for masers at 6049~MHz, which has recently been predicted to produce masers by pure collisional excitation at conditions similar to that required for 1720~MHz masing. The Effelsberg 100 m telescope was used to observe the excited-state 6016, 6030, 6035, and 6049~MHz lines of OH toward selected SNRs, most of which have previously-detected bright 1720~MHz masers. No excited-state masers are found toward SNRs, consistent with previous observations of the 6049~MHz and other excited-state transitions. We do not see clear evidence of absorption toward SNR target positions, although we do see evidence of absorption in the molecular cloud at $+50$~km\,s$^{-1}$ near Sgr~A East. Weak absorption is detected at 6016~MHz toward W3(OH), while stronger, narrower emission is seen at 6049~MHz, suggesting that the 6049~MHz emission is a low-gain maser. We conclude that conditions in SNRs are not conducive to excited-state maser emission, especially in excited-state satellite lines. | About two dozen supernova remnants (SNRs) are known to host 1720~MHz OH masers \citep[][and references therein]{green06}. The excitation mechanism for 1720~MHz masers in SNRs is commonly believed to be collisional excitation from a C-type shock \citep{lockett99,wardle99}. Other ground-state OH transitions are sometimes seen in absorption \citep[e.g.,][]{goss68}, which can be helpful for modelling the physical conditions in the 1720~MHz masing region \citep{hewitt06}. However, to date no other OH transition, nor a transition of any other molecule, has been detected as a maser associated with an SNR. The detection of a second maser transition in SNRs would place strong constraints on physical conditions (density, temperature, OH fraction, ortho-to-para H$_2$ ratio, etc.) in the masing region, especially if spatial coincidence with a 1720~MHz maser was observed. Excited-state transitions, which occur at higher frequencies, would be especially useful if detected toward the Galactic center, where angular scattering is large \citep[see \S~1 of][]{pihlstrom07}. Another motivation for finding a second masing transition in SNRs is to confirm the Zeeman interpretation of splitting seen between left and right circular (LCP, RCP) components at 1720~MHz. Differences in LCP and RCP velocities are used to compute magnetic field strengths in SNRs \citep[e.g.,][]{brogan00}, important for quantifying magnetic pressures. However, \citet{elitzur96,elitzur98} shows that circular polarization can be generated by non-Zeeman mechanisms when the maser is saturated and splitting is small compared to a line width, as is the case in SNR 1720~MHz masers. While there are theoretical reasons to believe the Zeeman interpretation of SNR 1720~MHz maser splitting \citep[see \S~4.1 of][]{brogan00}, direct confirmation of the magnetic field strengths from a second OH transition would lend greater confidence in the derived magnetic field strengths. Theoretical modelling predicts that collisions alone can excite several transitions of OH. In the low-density ($n \sim 10^5$~cm$^{-3}$) regime, only the 1720~MHz transition is inverted, but at higher number and/or column densities the 6049~MHz transition also becomes inverted, with some overlap in parameter space allowing simultaneous inversion of the 1720 and 6049~MHz lines \citep{pavlakis00,wardle07,pihlstrom07}. At still higher densities ($n \sim 5~\times~10^6$~cm$^{-3}$), the 4765 and 1612~MHz transitions become inverted \citep{pavlakis96,pihlstrom07}. The 6035~MHz line may also become inverted at high temperatures \citep[$T \sim 200$~K;][]{pavlakis00}. Regardless of the details of the pump mechanism, weak inversion (emission) is predicted in the 6049~MHz and weak anti-inversion (enhanced absorption) in the 6016~MHz line from the structure of the OH level diagram alone when infrared trapping becomes important \citep{litvak69,elitzur77}. Few observations have sought excited-state OH masers in SNRs. \citet{mcdonnell07} looked for 6030, 6035, and 6049~MHz OH maser emission toward southern SNRs but did not detect any 6049~MHz emission. Although they did find three new 6030 and 6035~MHz maser sources, \citet{mcdonnell07} believe that these masers are associated with \ion{H}{2} regions and therefore represent emission from star-forming regions (SFRs), not SNRs. In states of higher excitation, \citet{pihlstrom07} searched for emission in the 4.7~GHz $\Lambda$-doubling triplet and the 7.8 and 8.2~GHz quadruplets toward four SNR complexes (as well as the 23.8~GHz quadruplet toward Sgr~A~East) but did not detect any emission. \citet{pihlstrom07} also cross-correlated positions of single-peaked and irregular-spectrum 1612~MHz masers from blind surveys in the literature with positions of known SNRs but found no probable associations. For small velocity differences along the amplification path, the 6049~MHz transition is theoretically the most promising transition in which to find new OH SNR masers, given that the collisional excitation conditions require only slightly higher densities than those that produce 1720~MHz masers. We report on targeted observations of the four 6.0~GHz lines mostly toward SNRs with bright 1720~MHz masers. | Though our original interest was on the 6.0~GHz OH transitions in SNRs, we start off the discussion with our findings on our reference source, the SFR region W3(OH). \subsection{Satellite Lines in W3(OH)} \label{satellite} The 6049~MHz ($F = 3 \rightarrow 2$) emission in W3(OH) is stronger than the 6016~MHz ($F = 2\rightarrow 3$) absorption, as noted in the calculations by \citet{guilloteau84}, who did not detect the 6016~MHz absorption at all. This effect, in which 6016~MHz absorption is enhanced relative to the expected LTE value (from 6030 and 6035~MHz absorption) and the 6049~MHz line is seen in weak absorption or even emission, was seen in other SFRs by \citet{whiteoak76} and \citet{gardner83}. This behavior was predicted by \citet{elitzur77}, who showed that anti-inversion at 6016~MHz and inversion at 6049~MHz should occur whenever the 120~$\mu$m radiation linking the first-excited $J = 5/2$ state with the ground $J = 3/2$ state in the $^2\Pi_{3/2}$ ladder is optically thick, regardless of the details of the pump mechanism. Interestingly, a similar phenomenon is seen in the next highest state in the $^2\Pi_{3/2}$ ladder. Absorption in the 13442~MHz ($F = 3 \rightarrow 4$) transition is always stronger than at 13433~MHz ($F = 4 \rightarrow 3$) and always enhanced from LTE values compared to the main lines \citep{fish05}. The \citet{elitzur77} result is generalizable: whenever intra-ladder decays dominate and are optically thick, the $\Delta F = +1$ line will be seen in enhanced absorption and the $\Delta F = -1$ line will be seen in emission or reduced absorption. For the 13.4~GHz transitions, infrared overlaps in the 84~$\mu$m transitions linking the $J = 7/2$ and $5/2$ states may also be important, as proposed by \citet{matthews86}. We find that a Gaussian fit to the 6049~MHz emission in W3(OH) is also \emph{narrower} than the fit to the 6016~MHz absorption, even though the full width at zero power of the two lines is the same. This is consistent with weak inversion and suggests that the 6049~MHz emission should be understood as a low-gain maser, as discussed by \citet{baudry97b}. Further sensitive observations at higher spectral resolution will be required to confirm the maser nature of the 6049~MHz emission. High angular resolution interferometric observations would be particularly useful to determine the size and brightness temperature of the 6049~MHz emission region, as well as to isolate the region of emission and determine whether the inversion exists over the entire source or merely in cluster C \citep[in the nomenclature of][]{fish07}, which must be a region of very high excitation based on the presence of highly-excited OH \citep{baudry93,baudry97}. In general, satellite-line masers are seen almost exclusively in the lowest states of the $^2\Pi_{3/2}$ and $^2\Pi_{1/2}$ ladders. In the $^2\Pi_{3/2}, J = 3/2$ rotational state, masers are seen in the 1612~MHz transition in evolved stars and some SFRs, and 1720~MHz masers are seen in SNRs and some SFRs. Masers in the $^2\Pi_{1/2}, J = 1/2$ state are seen only in SFRs, with 4765~MHz ($F = 1 \rightarrow 0$) masers much more common than 4660~MHz ($F = 0 \rightarrow 1$) masers \citep[e.g.,][]{gardner83,cohen95}, in agreement with \citet{elitzur77}. Even though the brightest main-line 6030 and 6035~MHz masers in SFRs can have flux densities of several hundred~Jy \citep[e.g.,][]{fish07}, the only other satellite-line maser found to date is the low-gain 6049~MHz maser in W3(OH). It thus appears to be a general result that it is extremely difficult to invert satellite lines except in the lowest rotational state in each ladder. \subsection{Sgr~A East} \label{ae-discussion} We see 6030 and 6035~MHz absorption near $+50$~km\,s$^{-1}$ in pointings on the eastern side of Sgr~A East. The absorption at 6035~MHz is deeper than at 6030~MHz, as would be expected for thermal absorption. This absorption is likely neither part of the interaction of the SNR with the molecular cloud nor located at the interaction region \citep[with most 1720~MHz masers in the slightly higher velocity range $V_\mathrm{LSR} = 53$--$68$~km\,s$^{-1}$;][]{pihlstrom06}, although interferometric confirmation, as with the Expanded Very Large Array (EVLA) will be required to establish the location of the absorption. The molecular material is denser to the east and partially obscured by the circumnuclear disk to the west \citep{mcgary01,herrnstein05}, consistent with our absorption detections as well as the ground-state absorption detected by \citet{karlsson03}. The detection of excited-state absorption toward Sgr~A East but not toward any other SNR may find its explanation in a chance alignment, the special conditions in the Galactic center region, or a combination thereof. The difference with the other SNRs may be due to the geometry of a dense cloud in the line of sight with a relatively intense radio continuum background in the case of Sgr~A East, resulting in a larger column density and absorption of OH. Or, the difference may be due to a difference in heating of the cloud: the +50~km\,s$^{-1}$ cloud toward Sgr~A East may be heated more, not by the impact of the SNR into the cloud \citep{herrnstein05} but, e.g., due to a strong local radiation field absorbed by high-metallicity material, dissipation of kinetic energy or collisions of local clumps in a steep gravitational potential, etc., providing a larger column density of OH in excited states. As this absorption is likely not due to the SNR, further studies are required to investigate its origin. Interferometric observations would be very useful in order to understand the distribution of excited-state OH in the molecular cloud near Sgr~A East. Higher angular resolution than that provided by Effelsberg will be required to understand the origin of the 6.0~GHz absorption in Sgr~A East. This absorption is an obvious target for reobservation with the EVLA when a sufficient number of antennas are equipped with the new C-band (4--8~GHz) receivers. \subsection{Other Sources} No emission was seen in any 6.0~GHz transition toward SNRs. To date, a total of 23 $\Lambda$-doublet transitions have been observed toward SNRs: the quadruplets at 1.6, 6.0, 7.7, 8.1, and 23.8~GHz and the triplet at 4.7~GHz. Only the 1720~MHz transition produces a detectable maser. The 6049~MHz transition is the next one after 1720~MHz predicted to go into inversion as densities are increased \citep{pihlstrom07,wardle07}, yet searches for 6049~MHz masers have not uncovered any detections \citep[this work, as well as][]{mcdonnell07}. Searches for SNR OH masers at 4765 and 1612~MHz, the next two transitions expected to produce masers via collisional excitation, have also produced only non-detections \citep{pihlstrom07}. It is probable that no OH masers, other than in the 1720~MHz transition, will be detected toward SNRs with current instrumentation. Absorption at 6016~MHz is not seen toward any SNR, despite predictions that it should be anti-inverted. We do not see absorption toward SNRs with the exception of 6030 and 6035~MHz absorption toward Sgr~A East. Much more sensitive observations will be required to test whether excited-state $\Delta F = +1$ transitions are in enhanced absorption in SNRs. No emission was seen toward the evolved stars AU Gem and NML Cyg. Excited-state OH maser emission has previously been reported in the 4750~MHz transition toward AU Gem \citep{claussen81} and the 6035 (and possibly 6030) MHz transition toward NML Cyg \citep{zuckerman72}, although subsequent observations have failed to confirm these detections \citep[e.g.,][]{sjouwerman07}. It is likely that excited-state OH maser emission from evolved stars is temporary and extremely rare \citep{sjouwerman07}. Our continued non-detection of excited-state emission toward these stars supports this hypothesis. | 7 | 10 | 0710.1770 |
0710 | 0710.1293_arXiv.txt | {We investigate, by means of numerical simulations, the possibility of forming counter-rotating old stellar components by major mergers between an elliptical and a spiral galaxy. We show that counter-rotation can appear both in dissipative and dissipationless retrograde mergers, and it is mostly associated to the presence of a disk component, which preserves part of its initial spin. In turn, the external regions of the two interacting galaxies acquire part of the orbital angular momentum, due to the action of tidal forces exerted on each galaxy by the companion. } | Simulations show that stellar counter-rotation in galaxies could emerge thanks to two different processes: dissipative and dissipationless mergers.\\ Confirming the suggestion of \citet{korm84}, \citet{balquin90} showed that unequal mass mergers of elliptical galaxies % can produce counter-rotation if the orbit of the encounter is retrograde with respect to the spin of the primary. They also pointed out that the rotation seen in the counter-rotating component is a tracer of the orbital angular momentum and that both primary and secondary stars counter-rotate at the core.\\ \citet{hb91} showed that counter-rotating central gas disks can form as a result of retrograde mergers between two gas rich spiral galaxies, discussing the possibility that star formation in such disks could produce components with decoupled kinematics, as in the core of some elliptical galaxies.\\ Some years later, \citet{balgon98} showed that kinematically peculiar cores may be generated also in retrograde stellar spiral-spiral mergers: in this picture, the central bulges transport orbital angular momentum inward to the center of the remnant, while the outer parts keep the spin signature of the precursor disks. Also \citet{bendob00} put in evidence the possibility to form counter-rotation at large radii, simulating mergers between equal mass disk galaxies. Finally, \citet{jes07}, studying the 2D kinematics of a sample of simulated disk merger remnants, showed that counter rotating cores, made of old stellar populations, are almost exclusively formed in equal mass mergers where a dissipative component is included. Evidently, both mechanisms (the dissipative and dissipationless one) can occur in real systems, producing a variety of kinematically decoupled components, of different ages and physical extensions \citep{mc06}. In this paper, we want to present a new scenario, according to which stellar counter-rotation can form both in dissipative and dissipationless retrograde\footnote{i.e. the spin of the spiral galaxy is initially anti-parallel to the orbital angular momentum.} major mergers of elliptical-spirals galaxies\footnote{Note that all the previous numerical works refer to mergers between two ellipticals or two disk galaxies. So far, the possibility of forming counter-rotating cores by mergers of galaxies of different morphologies has never been exploited.}. In this scenario, stars in the external regions of the galaxies involved in the encounter acquire orbital angular momentum, at first pericenter passage and, mostly, in the last phases of the merging process, while those in the most inner regions can maintain some of the initial spin (anti-parallel to the orbital one), producing decoupled counter-rotating central components. In the case of dissipative mergers, the central decoupled core could be composed of two different populations: the old stellar population, which has preserved its initial spin, and a new stellar population, born \emph{in situ} from the kinematically decoupled gas component. The redistribution of the angular momentum between the different galactic components (gas, stars and dark matter) is discussed. | We have presented a new scenario to form counter-rotating central components in early-type galaxies, by dissipative and dissipationless mergers of elliptical-spiral systems in retrograde orbits. In the case of dissipative mergers, the central decoupled core could be composed of two distinct populations: the old stellar population, which has preserved part of its initial spin, and a new stellar population, born \emph{in situ} from the kinematically decoupled gas component. Counter-rotating cores are also found in the remnants of non-coplanar mergers, at least for inclinations $i \le 20^0$ (being $i$ the angle between the spiral disk and the orbital plane). We plan to realise a wider and systematic study of the role played by orbital parameters and morphology of the interacting systems in the near future, varying the morphological parameters of the interacting galaxies (bulge-to-disk ratio of the spiral galaxy as well as different initial models for the elliptical), and the orbital initial conditions. % | 7 | 10 | 0710.1293 |
0710 | 0710.5897_arXiv.txt | \vspace{0.5cm} Modern cosmology has now emerged as a testing ground for theories beyond the standard model of particle physics. In this paper, we consider quantum fluctuations of the inflaton scalar field on certain noncommutative spacetimes and look for noncommutative corrections in the cosmic microwave background (CMB) radiation. Inhomogeneities in the distribution of large scale structure and anisotropies in the CMB radiation can carry traces of noncommutativity of the early universe. We show that its power spectrum becomes direction-dependent when spacetime is noncommutative. (The effects due to noncommutativity can be observed experimentally in the distribution of large scale structure of matter as well.) Furthermore, we have shown that the probability distribution determining the temperature fluctuations is not Gaussian for noncommutative spacetimes. | \label{sec:intro} The CMB radiation shows how the universe was like when it was only $400, 000$ years old. If photons and baryons were in equilibrium before they decoupled from each other, then the CMB radiation we observe today should have a black body spectrum indicating a smooth early universe. But in 1992, the Cosmic Background Explorer (COBE) satellite detected anisotropies in the CMB radiation, which led to the conclusion that the early universe was not smooth: There were small perturbations in the photon-baryon fluid. The theory of inflation was introduced \cite{guth, Linde, Albrecht} to resolve the fine tuning problems associated with the standard Big Bang cosmology. An important property of inflation is that it can generate irregularities in the universe, which may lead to the formation of structure. Inflation is assumed to be driven by a classical scalar field that accelerates the observed universe towards a perfect homogeneous state. But we live in a quantum world where perfect homogeneity is never attained. The classical scalar field has quantum fluctuations around it and these fluctuations act as seeds for the primordial perturbations over the smooth universe. Thus according to these ideas, the early universe had inhomogeneities and we observe them today in the distribution of large scale structure and anisotropies in the CMB radiation. Physics at Planck scale could be radically different. It is the regime of string theory and quantum gravity. Inflation stretches a region of Planck size into cosmological scales. So, at the end of inflation, physics at Planck region should leave its signature on the cosmological scales too. There are indications both from quantum gravity and string theory that spacetime is noncommutative with a length scale of the order of Planck length. In this paper we explore the consequences of such noncommutativity for CMB radiation in the light of recent developments in the field of noncommutative quantum field theories relating to deformed Poincar\'e symmetry. The early universe and CMB in the noncommutative framework have been addressed in many places \cite{Greene, Lizzi, Brandenberger1, Huang, Brandenberger2, BalBB, Fatollahi2, Fatollahi1}. In \cite{Greene}, the noncommutative parameter $\theta_{\mu \nu} = -\theta_{\nu \mu} =\textrm{constants}$ with $\theta_{0i} =0$, ($\mu, \nu = 0, 1, 2, 3$, with $0$ denoting time direction), characterizing the Moyal plane is scale dependent, while \cite{Brandenberger1, Brandenberger2, Huang} have considered noncommutativity based on stringy space-time uncertainty relations. Our approach differs from these authors since our quantum fields obey twisted statistics, as implied by the deformed Poincar\'e symmetry in quantum theories. We organize the paper as follows: In section II, we discuss how noncommutativity breaks the usual Lorentz invariance and indicate how this breaking can be interpreted as invariance under a deformed Poincar\'e symmetry. In section III, we write down an expression for a scalar quantum field in the noncommutative framework and show how its two-point function is modified. We review the theory of cosmological perturbations and (direction-independent) power spectrum for $\theta_{\mu \nu}=0$ in section IV. In section V, we derive the power spectrum for the noncommutative Groenewold-Moyal plane ${\cal A}_{\theta}$ and show that it is direction-dependent and breaks statistical isotropy. In section VI, we compute the angular correlations using this power spectrum and show that there are nontrivial ${\cal O}( \theta^{2})$ corrections to the CMB temperature fluctuations. Next, in section VII, we discuss the modifications of the $n$-point functions for any $n$ brought about by a non-zero $\theta^{\mu \nu}$ and show in particular that the underlying probability distribution is not Gaussian. The paper concludes with section VIII. | In this paper, we have shown that the introduction of spacetime noncommutativity gives rise to nontrivial contributions to the CMB temperature fluctuations. The two-point correlation function in momentum space, called the power spectrum, becomes direction-dependent. Thus spacetime noncommutativity breaks the rotational invariance of the CMB spectrum. That is, CMB radiation becomes statistically anisotropic. This can be measured experimentally to set bounds on the noncommutative parameter. Currently, we \cite{numerical} are making numerical fits to the available CMB data to put bounds on $\theta$. We have also shown that the probability distribution governing correlations of fields on the Groenewold-Moyal algebra ${\cal A}_{\theta}$ are non-Gaussian. This affects the correlation functions of temperature fluctuations. By measuring the amount of non-Gaussianity from the four-point correlation function data for temperature fluctuations, we can thus set further limits on $\theta$. We have also discussed the signals of non-causality of non-commutative field theories in the temperature fluctuations of the CMB spectrum. It will be very interesting to test the data for such signals. | 7 | 10 | 0710.5897 |
0710 | 0710.2872.txt | This contribution intends to give a pedagogical introduction to the topic of dark energy (the mysterious agent supposed to drive the observed late time acceleration of the Universe) and to various observational tests which require only assumptions on the geometry of the Universe. Those tests are the supernovae luminosity, the CMB shift, the direct Hubble data, and the baryon acoustic oscillations test. An historical overview of Cosmology is followed by some generalities on FRW spacetimes (the best large-scale description of the Universe), and then the test themselves are discussed. A convenient section on statistical inference is included as well. | Cosmology is a branch of physics which is experiencing a tremendously fast development lately triggered by the arrival of many new observational data of ever more exquisite precision. These findings have been crucial for the improvement in our understanding of the Universe, but the vast amount of knowledge on our Cosmos which we have nowadays would never have been possible without the concerted effort by experimentalists and theoreticians. One of the results of this fantastic intellectual pursue is the puzzling discovery that there is in the Universe a manifestation of the repulsive side of gravity. This is the topic to which this lectures are devoted, and more specifically I wish to address some methods which the community believes are useful for building a deeper understanding of why, how and when our universe began to accelerate. Put in more modest words, these lectures will dissert about how one can take advantage of various observational datasets to describe some basic features of the geometry of the Universe with the hope they will reveal us something on the nature of the agent causing the observed accelerated expansion. As this is quite an advanced topic in Physics (Astronomy), this contribution will build tougher as we proceed, so I hope the readers will enjoy to start off from a little historic stroll in the science of surveying the skies. For wider historical overviews than the one presented here, the two main sources I recommend to you, among the so many available, are \cite{conv} and \cite{wiki}. Historical records tells us that Astronomy was born basically because of the need by agrarian societies to predict seasons and other yearly events and by the esoteric need to place humanity in the Universe. This discipline developed in many ancient cultures (Egyptian, Chinese, Babylonian, Mayan), but it was only Greek people who cared understanding their observations and spreading their knowledge unlike in other cultures. Their influence was crucial as the modern scientific attitude of relying in empiricism (Aristotlean school) and translating physical phenomena into the mathematical language (Pithagorean school) built on their way of approaching science. Unfortunately, Aristotle was far so influential that his erroneous geocentric view of the Universe was not questioned aloud until the XVIth century. At the beginning of that century Copernicus started to spread quietly his heliocentric cosmological model, and even though he did not make much propaganda about his ideas, they deeply influenced other figures, such as Galileo, who introduced telescopes into astronomy and found solid evidence against the geocentric model. Another very important contribution to the subject was made by Kepler, who gave accurate characterization of the motions of the planets around the Sun. These findings were later on synthetically explained by Newton in his theory of gravitation, which built on Galileo's developments on dynamics (the law of inertia basically). The XVIIIth and XIXth century brought advances in the understanding of the Universe as a whole made of many parts which lie not necessarily in the Solar system. Stars began to be regarded as far-away objects in motion, and other objects like nebulae were discovered, so the idea of the existence of complicated structures outside the Solar system gained solidity. A discovery very much related to the main topic of this lectures was the realization that the combination of the apparent brightness of a star and its distance get combined to give its intrinsic brightness; basically the amount of energy reaching us in the form of light from a distant source decreases is inversely proportional to the square of the distance between the source an us (in an expanding the universe definition of distance is not the same as in a static one but this broad way of speaking applies all the same) This finding let observers realize that stars were objects very much like the Sun, or if you prefer they realized the Sun was nothing but yet another star. The most important next breakthrough was Einstein's theory of special relativity, which generalized Galileo's relativity to introduce light. On the conceptual realm, this was a very revolutionary theory at is warps the notions of space and time, and so it set the foundations for the best description of gravity so far: Einstein's theory of General Relativity. This second theory was the fruition of Einstein endeavors to unify the interactions known to him. In this theoretical framework matter/energy modifies the geometry, and in turn geometry tells matter/energy how to move/propagate (paraphrasing Wheeler's renowned quotation). Among other predictions this theory made a couple which are cornerstones of modern astronomy: gravitational redshift (light gets redder as it moves away from massive objects), and gravitational lensing (light gets bent as it passes close to massive objects). The next important advance come from the side of observations. In 1929 Hubble, and after having collected data carefully for almost a decade, presented the surprising conclusions that galaxies on average move away from us. This effect is encoded in the scale-invariant relation known as Hubble's law: $ v=H d$, where $v$ is the galaxy's velocity and $d$ its distance from us. The positiveness of the quantity $H$ as measured by Hubble is precisely what told him the Universe is expanding, this being a discovery which gets accommodated nicely in Einstein's theory of general relativity. Interestingly, on learning about this finding, Einstein discarded his idea of the necessity of some exotic fluid with negative pressure to counteract the attractive effect of usual matter (which would make the universe contract). It cannot look but funny from today's perspective that Einstein's idea has come to life again, at the end of the day the exotic fluid he imagined is one of the possible flavors of what the community calls dark energy \cite{turner} these days. At the risk of not giving everyone the credit they merit, I will just say that recognition for the concept of a expanding Universe is due to both theoreticians (de Sitter, Friedmann, Lema\^itre, Gamow,...) and experimentalists (Slipher, Hubble,..). However, the trampolin for this idea to jump into orthodoxy was put by Penzias and Wilson \cite{penzias} who first detected the cosmic microwave background, and by Dicke, Peebles, Roll and Wilkinson \cite{peebles} who were responsible for the not the less important interpretation of those observations. The existence of this radiation and its characteristic black body spectrum are a prediction of the Big-Bang theory, it is fair to say that if one combines it with other sources of evidence, it is almost impossible to refute it. The CMB in an invaluable source of cosmological information (visit \cite{hu} for a higly recommended site on the subject). It has a temperature of of around $2.73\,K$, with tiny temperature differences of about $10^{-5}$ between different patches of the sky. These anisotropies inform us that the photons forming the CMB where subject to an underlying gravitational potential which had fluctuations, and this is just and indication of density irregularities that seeded the structure one observes today. On the other hand, the image of the features of the CMB in different angular scales strongly favors a universe with no spatial curvature (flat or Euclidean on constant time slices) that is, it supports the theory of inflation. In addition, the CMB has a saying about the fraction of the different fluids filling the Universe. Fortunately, the mine of cosmic surprises was far from being exhausted, and it stored a diamond of many carats in the form of a discovery which has changed greatly the mainstream view of the Universe. Up to the late 1990, no repulsive manifestation of gravity had been spotted in Nature. In 1998 astronomical measurements provided evidence that gravity can not only push, but pull as well, and the repulsive side of gravity got unveiled. Very refined observations of the brightness of distant supernovae \cite{Riess:1998cb,Perlmutter:1998np} seemed to hint the presence of a negative pressure component in the Universe which would make it accelerate, and then a new revolution started. One may wonder at this stage what is the relation of supernovae luminosity and cosmic speed up. Objects in the Universe with a well calibrated intrinsic luminosity (supernovae, for instance) can be used to determine distances on cosmological scales. Supernovae are very bright objects, and so they result particularly attractive for this purpose (they can be $10^9$ times more luminous than the Sun so there is hope they will be visible from up to perhaps $1000$ Mpc \cite{physto}). As we anticipated in the last paragraph, in 1998 two independent teams reported evidence that some distant supernovae were fainted than expected. This involved tracing the expansion history of the Universe by combining measurements of the recession velocity, apparent brightness and distance estimations. The most compelling explanation was (and keeps being as far I am concerned) that their light had traveled greater distances than assumed. The orthodox view up to then was that the expansion pace of the Universe was barely constant, but supernovae seemed to contradict this. This unexpected and exciting discovery obliged researchers to broaden their mind and accept the Universe is undergoing accelerated expansion. I should have been able to have convinced the readers by now that these are exciting times to be working on Cosmology, as cosmic speed up is such and intriguing phenomenon with major open questions such as whether dark energy evolves with time, how much of it is there, and if is rather not a manifestation of extradimensional physics. The answer to these questions requires cannot be dissociated from the response to a perhaps more fundamental question: what is the Universe made of? The combination of various astronomical observations tell us our universe is basically made of thee major components (see for instance \cite{wmap}). The most abundant one is dark energy \cite{turner}, so this makes it even more interesting to find out whatever we can about it. At the other end the by far least abundant component is baryonic matter, and in the middle (as abundance is concerned) we have dark matter, in a proportion comparable to that of dark energy. There are various sources of astrophysical giving evidence in favor of it. Hints of it is existence are provided by the motion of stars, galaxies and clusters, but it is known it also played a crucial role in the amplification of the primordial density fluctuations which seeded the large scales structure we observe today and dark matter imprints can be found in the CMB as well. Dark matter represents quite a challenge as its nature remains a mystery; nevertheless if it were baryonic we know from big bang nucleosynthesis helium-4 would get converted to deuterium much easily and CMB calculations indicate indicate in addition anisotropies would be much larger so the odds are most of its is not baryonic. Up to here we have made a very broad introduction to our topic with a little bit of history and a little bit of physics, but we must not forget Maths are key to Cosmology \cite{Tegmark}. We only know how to study the Universe using numbers and equations, but of course progress in this direction is done with as many reasonable simplifications as possible (if they do not compromise rigor, of course). Einstein equations relate geometry and matter/energy content of the Universe. Cosmologists are concerned by this relation on a large scales picture so as to understand the expansion of the Universe. Those equations are non-linear, so studying them is painstaking unless one exploits the observational evidence the regularity of the Universe. Non-linearity of those equations makes their study a really hard task so simplifications are a must. The two basic ones are that galaxies are homogeneously distributed on galaxies larger than $50$ {\rm Mpc} \cite{Yadav}, and that the Universe is isotropic around us on angular scales larger than about $10$ degrees \cite{Souradeep}. But this is not enough, those two simplifications, which come from observational evidence must be completed with the assumption we occupy no special place in the Universe. This puts on the track of what we could call the parameterized Universe, cosmologists work with a greatly a simplified geometric description of the Universe which emerges from the latter assumptions. The models for the sources are of reduced complexity too, the most common being perfect fluids and fields with known dynamics. This we have on the side of theory, but on the side of observations we have to make our own life easier too. When doing observations oriented Cosmology, one uses as a variable the redshift $z$ of the electromagnetic radiation received, as it encodes information of how much the Universe has expanded between emission and reception. Observational tables of geometrical quantities can be given, and then one can test different theoretical values of the same quantities corresponding to models of interest. The theoretical predictions will depend on the sources assumed, and ultimately it will be possible to estimate the suitability of a given dark energy model to observations, but which are the available probes of dark energy? Basically dark energy can be scrutinized observationally from two main perspectives \cite{trotta}: one possility is doing it through its effect of the growth of structures, another one is through its impact on geometrical quantities. We will concentrate on geometrical constraints/tests in these lectures as in a way they are those which can perhaps be applied with less difficulty, although their simplicity does not mean they are the least interesting, for instance, the supernovae test is the only test giving a direct indication of the need of a repulsive component in the Universe, whereas the baryon acoustic oscillations test is thought to have much information in store. Discussion on these two tests will be given later on in these lectures. Finally, there is one more direction in which these topic is related to Maths apart from the geometrical side of it, statistics plays an important role too. Physics is attractive because of its ability to ``tame'' natural phenomena in the sense that laws of physics bring order to the apparent chaos of Nature, Astrophysics is even more attractive because it evidences how laws of physics apply outside Earth, which is certainly surprising given the manifest differences between Earth and every other locations in the Universe of which we are aware. However, since these phenomena occur in places far, far away, and sometimes they can only be observed indirectly, there is typically an important degree of uncertainty. Thus, research is Astrophysics requires understanding not only Physics, but also inference. Inferential statistics cares for the identification of patterns in the data taking in account the randomness and uncertainties in the observations. In contrast, descriptive statistics is concerned with giving a summary of the data either numerically or graphically. Both will be needed toward two goals: we need to know optimal ways to extract information from the astronomical data, but we also need to know rigorous approaches to compare theoretical predictions to observations. Our approach will be that of Bayesian inference, and we will justify jut below my preference, but it must be admitted there is an old vivid controversy on the definition of probability between the two main schools: frequentists and Bayesians. Frequentists use a definition based on the possibility of repeating the experiment, but this does not apply to the Universe, and this sounds like the reason why the number of Bayesian astronomers grows every day, but you should not care for this battle right now; just stick to the idea that observations related Cosmology needs to resort to Statistics and that Inference is vital for constraining dark energy cosmological models. By now you should be able to guess what to expect in the next sections. I will present you some basics of FRW (Friedmann-Robertson-Walker) cosmologies, then I will devote a great deal of this text to details of geometrical tests, and leave for the last but one section a convenient primer on statistics. | In this contribution I have tried to review some of the background on geometrical tests of dark energy models. A historical review of the development of Cosmology has been followed by an account of the importance of the discovery of the late-time acceleration, which is commonly attributed to the existence of an exotic component in the cosmic budget. Other possible explanations have been attempted, which required being open minded enough to admit the existence of extra dimensions. Examples of both conventional (if that adjective can be used) and extradimensional models have been mentioned here with respect to their geometrical features. All these models can be subject to different observational tests which only require postulating a parametrization of the Hubble factor of the Universe or quantities derived from it. Here I have been concerned with four tests only: luminosity of supernovae, direct $H(z)$ measurements, the CMB shift and baryon acoustic oscillations. I cannot deny this selection is biased in the sense these are tests I have made research on, but the first and the last one are definitely very important and much effort is doing by research groups all over the world in finding the data their application requires. This contribution has also tried to maintain a pedagogical tone as it has been prepared for the Advanced Summer School 2007 organized by Cinvestav in Mexico DF. I just hope it will be useful to any reader which happens to come across it. | 7 | 10 | 0710.2872 |
0710 | 0710.5918_arXiv.txt | {} {We present the results of near-infrared, follow-up imaging and spectroscopic observations at VLT, aimed at characterizing the long-period companions of the exoplanet host stars HD\,196885, HD\,1237 and HD\,27442. The three companions were previously discovered in the course of our CFHT and VLT coronographic imaging survey dedicated to the search for faint companions of exoplanet host stars.} {We used the NACO near-infrared adaptive optics instrument to obtain astrometric follow-up observations of HD\,196885\,A and B. The long-slit spectroscopic mode of NACO and the integral field spectrograph SINFONI were used to carry out a low-resolution spectral characterization of the three companions HD\,196885\,B, HD\,1237\,B and HD\,27442\,B between 1.4 and 2.5~$\mu$m.} {We can now confirm that the companion HD\,196885\,B is comoving with its primary exoplanet host star, as previously shown for HD\,1237\,B and HD\,27442\,B. We find that both companions HD\,196885\,B and HD\,1237\,B are low-mass stars of spectral type M$1\pm1$V and M$4\pm1$V respectively. HD\,196885\,AB is one of the closer ($\sim23$~AU) resolved binaries known to host an exoplanet. This system is then ideal for carrying out a combined radial velocity and astrometric investigation of the possible impact of the binary companion on the planetary system formation and evolution. Finally, we confirm via spectroscopy that HD\,27442\,B is a white dwarf companion, the third one to be discovered orbiting an exoplanet host star, following HD\,147513 and Gliese~86. The detection of the broad Br$\gamma$ line of hydrogen indicates a white dwarf atmosphere dominated by hydrogen.} {} | The radial velocity (RV) technique is without contest nowadays the most successful method for detecting and characterizing the properties of exo-planetary systems. Since the discovery of 51 Peg (Mayor \& Queloz 1995), more than 200 exo-planets have been identified featuring a broad range of physical (mass) and orbital (P, $e$) characteristics (Udry et al. 2003). Major progress has been made in improving detection performance and data analysis and has enabled us to explore the mass regime down to neptunian and telluric masses (Santos et al. 2004; McArthur et al. 2004). Hitherto, most surveys have been focused on solar-type stars because these stars show more and thinner lines than their more massive counterparts and less activity than their less massive ones, which ensures in both cases comparatively higher RV precision. Only recently have been planet-search programs devoted to late-type stars with the use of high precision spectrograph and repeated measurements (Delfosse et al. 1998; Endl et al. 2003; Wright et al. 2004; Bonfils et al. 2005) as well as to early-type stars with the development of new methods for radial velocity measurements (Chelli 2000; Galland et al. 2005). Despite the success of this technique, the time span explored limits the study to the close ($\le4-5$~AU) circumstellar environment. To understand the way exo-planetary systems form and evolve, it is therefore clearly worthwhile using complementary techniques such as pulsar timing, micro-lensing, photometric transit or direct imaging to fill out our knowledge. For solar analogs, the current deep-imaging capability is limited to the detection of massive brown dwarf (BD) companions. Typical separations larger than $50-100$~mas (i.e. $\geq 3-5$~AU from a star at 50~pc) can be explored in the stellar regime and $0.5-1~\!''$ (i.e. $\geq 30-50$~AU from a star at 50~pc) in the substellar regime. The recent discovery of a T7 dwarf companion at 480~AU from the exoplanet host star HD\,3651 (Mugrauer et al. 2006) illustrates the scope of deep, near-infrared imaging for resolving ultra-cool companions at large separations. Recent efforts have been also devoted to systematic search for stellar companions to nearby stars with and without planets, aimed at studying the impact of stellar duplicity on planet formation and evolution (Eggenberger et al. 2007). Since 2003, we have conducted a deep-coronographic imaging survey of 26 exoplanet host stars, using PUEO-KIR at CFHT, and NACO at VLT (Chauvin et al. 2006). Three probable companions were detected around the stars HD\,196885, HD\,1237 and HD\,27442 (see Tables~\ref{tab:tab1} and \ref{tab:tab2}). Follow-up observations were obtained for HD\,1237\,B and HD\,27442\,B, confirming their companionship (Chauvin et al. 2006; Raghavan et al. 2006). Based on their photometry, HD\,196885\.B and HD\,1237\,B are likely to be low-mass stars and HD\,27442\,B is probably the third white dwarf companion known to date orbiting an exoplanet host star. In this paper, we report new imaging and spectroscopic observations of these three companions, using NACO and SINFONI at VLT. In section~2, the instrument set-up and the data analysis are described. In section~3, we discuss the astrometric and spectroscopic results that confirm and refine their previous suggested nature. \begin{table}[t] \caption{Characteristics of the observed exoplanet host stars HD\,196885\,A, HD\,1237\,A and HD\,27442\,A.} \label{tab:tab1} \centering \begin{tabular*}{\columnwidth}{@{\excs}lllllll} % \hline\hline Name & SpT & d & Age$^a$ & $\rm{M}_{2}\rm{sin}i$$^b$ & $P\,^b$ & $e\,^b$ \\ & & (pc) & (Gyr) & $(M_{\rm{Jup}})$ & (days) & \\ \\ \hline HD\,196885\,A & F8IV & 33.0 & 0.5 & 1.84 & 386.0 & 0.30 \\ HD\,1237\,A & G6V & 17.6 & 0.8 & 1.94 & 311.29 & 0.24\\ HD\,27442\,A & K2IV & 18.2 & 10 & 1.28 & 423.84 & 0.07 \\ \hline \end{tabular*} \begin{list}{}{} \item[\scriptsize{($^a$) AGE REFERENCES:}] \scriptsize{Naef et al. 2001; Randich et al. 1999; Lambert et al. 2004} \item[\scriptsize{($^b$) RADIAL VELOCITY REFERENCES:}]\scriptsize{Naef et al. 2001; Butler et al. 2001; http://exoplanets.org/esp/hd196885/hd196885.shtml} \end{list} \end{table} | We present the follow-up imaging and spectroscopic characterization of three long-period companions to the exoplanet host stars HD\,196885, HD\,1237 and HD\,27442. The three objects were discovered during a previous deep-imaging survey carried out at CFHT and at VLT. Our new observations confirm their companionship unamiguously as well as their nature, which had previously been inferred from their photometry. Whereas HD\,196885 and HD\,1237 are two stellar companions of spectral type M$1\pm1$V and M$4\pm1$V respectively, HD\,27442 is the second confirmed white dwarf companion of an exoplanet host star. The detection of the broad Br$\gamma$ hydrogen line indicates a white dwarf atmosphere dominated by hydrogen. HD\,196885\,AB is one of the closer resolved binaries known to host an exoplanet. The presence of this long-period companion may have played a key-role in the formation and the evolution of the inner planetary system, but how? The two main mechanisms of planetary formation, core accretion and disk instability, do not lead to the same predictions while the impact of a binary companion, depending on the separation and the mass ratio, is still debated. A complete dynamic characterization of nearby, tight binaries with planets and a dedicated imaging survey to study the multiplicity among stars with and without planets are clearly required to throw new light on the mechanisms of planetary formation and evolution. | 7 | 10 | 0710.5918 |
0710 | 0710.5129_arXiv.txt | { Unparticle $\U$ with scaling dimension $d_\U$ has peculiar thermal properties due to its unique phase space structure. We find that the equation of state parameter $\omega_\U$, the ratio of pressure to energy density, is given by $1/(2d_\U +1)$ providing a new form of energy in our universe. In an expanding universe, the unparticle energy density $\rho_\U(T)$ evolves dramatically differently from that for photons. For $d_\U >1$, even if $\rho_\U(T_{\rm D})$ at a high decoupling temperature $T_{\rm D}$ is very small, it is possible to have a large relic density $\rho_\U(T^0_\gamma)$ at present photon temperature $T^0_\gamma$, large enough to play the role of dark matter. We calculate $T_{\rm D}$ and $\rho_\U(T^0_\gamma)$ using photon-unparticle interactions for illustration. \PACS{ {98.80.Cq}{Particle-theory and field-theory models of the early Universe} \and {11.15.Tk}{Other nonperturbative techniques} \and {11.25.Hf}{Conformal field theory, algebraic structures} \and {14.80.-j}{Other particles} } % } % | 7 | 10 | 0710.5129 |
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0710 | 0710.1961_arXiv.txt | In this paper, we present a detailed hydrodynamical study of the properties of the flow produced by the collision of a pulsar wind with the surrounding in a binary system. This work is the first attempt to simulate interaction of the ultrarelativistic flow (pulsar wind) with the nonrelativistic stellar wind. Obtained results show that the wind collision could result in the formation of an "unclosed" (at spatial scales comparable to the binary system size) pulsar wind termination shock even when the stellar wind ram pressure exceeds significantly the pulsar wind kinetical pressure. Moreover, the post-shock flow propagates in a rather narrow region, with very high bulk Lorentz factor ($\gamma\sim100$). This flow acceleration is related to adiabatical losses, which are purely hydrodynamical effects. Interestingly, in this particular case, no magnetic field is required for formation of the ultrarelativistic bulk outflow. The obtained results provide a new interpretation for the orbital variability of radio, X-ray and gamma-ray signals detected from binary pulsar system PSR~1259-63/SS2883. | Pulsars lose their rotation energy through relativistic winds, the collision of which with the Interstellar Medium results in the formation of the Pulsar Wind Nebulae (regions of nonthermal synchrotron radiation of ultrarelativistic electrons accelerated at the termination of pulsar winds \citep{rees74,kc}). The Crab Nebula is the most famous example of such an object. The recent X-ray and TeV gamma-ray observations \citep{GenSlane} show that this is a common phenomenon. A very interesting situation arises when a pulsar is located in a binary system. In this case the pulsar wind interacts with the wind from the companion star. This case, in particular, is realized in the binary system PSR1259-63/SS2883 which consists of a $\sim 48 \rm ms$ pulsar in an elliptic orbit around a massive B2e optical companion \citep{johnston1}. The density and velocity of the stellar wind depend on the separation distance between two stars. Thus, the processes related to the interaction of two winds, in particular particle acceleration and radiation proceed under essentially different physical conditions depending on the orbital phase. This makes this object a unique laboratory for the study of nonthermal processes in ``on-line'' regime, due to the short acceleration and cooling time-scales characterizing this system, especially close to the periastron \citep{khangoulian_psr}. Observations show that this system is indeed a strong source of nonthermal time-dependent emission extending from radio to TeV gamma-rays (see e.g. \cite{neronov}). Two variable TeV galactic gamma-ray sources, LS 5039 and LSI~61~303 (see e.g. \cite{paredes_rev}), are discussed as possible, although less evident, candidates representing ``binary pulsar'' source population \citep{dubus,mirabel,neronov2}. LS 5039 consists of O6.5V star and an unidentified compact object in a 3.9 day orbit. This object has been detected as the source of gamma rays by EGRET (source 3EG J1824-1514) \citep{ls5039EGRET} and by HESS \citep{ls5039}. LSI~61~303 is a binary system with a B0Ve star in a 26.5 day orbit. This source presumably associates with a low energy gamma-ray source 2CG 135+01/3EG J0241+6103 \citep{lsiaspulsar,tavani_lsi}. Recently, it has been detected in TeV gamma-rays \citep{lsi61303}. The nature of the compact objects (black hole or a neutron star) in both sources is not yet firmly established. The collision of supersonic winds from two stars located in a binary system results in the formation of two terminating shock fronts and a tangential discontinuity separating relativistic and nonrelativistic parts of the shocked flow. It is believed that the shocked flow should propagate into a limited solid angle with rather high velocity. However the properties of the flow downstream the shock is unknown. The collision of the pulsar wind with the supersonic flow of the nonrelativistic plasma has been recently numerically modeled by several groups \citep{bucciantini,swaluw,vigelius,romero} using different nonrelativistic versions of hydrodynamical codes. However, the dynamics of the relativistic plasma is rather specific, and the use of nonrelativistic codes cannot be \textit{a priori} justified. The modeling of collision of a relativistic pulsar wind with a nonrelativistic one demands an adequate treatment of distinct features of relativistic outflows. Moreover, in many previous studies the stellar wind has been approximated as plane parallel, while both winds initially expand radially. However, this purely geometric difference in the formulation of the problem, leads, in fact, to significantly different results and conclusions. In this paper we present the results of our studies conducted in the hydrodynamical limit, i.e. the role of the magnetic field in dynamics has been ignored. The impact of the magnetic field will be published elswhere. | In this paper, we have conducted a detailed numerical study of the interaction of the relativistic and nonrelativistic winds assuming isotropic, radially expanding winds and ignoring the role of the magnetic field. In fact, the cold ultrarelativistic pulsar winds are generally believed to be highly anisotropic \citep{bogovalov99,lyubarsky}. The anisotropy of the energy flux in the wind can result in a nonaxisymmetric form of the termination shock wave at the vicinity of the symmetry axis, as it follows from the studies of \cite{vigelius}. Whether the anisotropy has a noticeable impact on the result and conclusions of this paper is a subject of further investigations and will be discussed elsewhere. This concerns also the role of the magnetic field. Generally, as it follows from calculations of the interaction of the pulsar wind with the interstellar medium \citep{crab,lk,bucciantini2}, the role of the magnetic field is rather important in the post shock region almost independent of the strength of the field. This is explained by fast amplification of the magnetic field in the post shock region due to deceleration of the flow \citep{Khangoulian}. However, as shown in this paper, the flow is not decelerated. Just the opposite - it can be accelerated to large bulk motion Lorentz factor. Thus although one should expect an impact of the magnetic field on the post shocked flow, this affect most likely will not be so strong as in the case of the interaction of the pulsar wind with interstellar medium. Detailed MHD calculations are needed to to clarify the role of the magnetic field. The effect of reacceleration of the post shock flow to relativistic bulk motion Lorentz factors has direct implication to the interpretation of observations of high energy $\gamma$- and X-rays from binary pulsar systems like PSR 1259-63/SS2883. This effect strongly modifies the relationship between the synchrotron X-ray and inverse Compton gamma-ray fluxes produced by the same population of relativistic electrons. It is well known that in the pulsar wind nebulae (plerions) which are formed, by the interaction of pulsar winds with the interstellar medium, the ratio between the X-ray and VHE $\gamma$-ray fluxes are defined by the ratio of the energy density of the magnetic field to the energy density of soft radiation field, provided that Compton scattering takes place in the Thompson regime \citep{atoyan}. In binary systems inverse Compton scattering proceeds in the Klein-Nishina regime which changes the relationship between the X-ray and VHE gamma-ray fluxes \citep{khangulyan2}. Due to our calculations it becomes clear that a significant deviation from the standard relations should be expected also from hydrodynamics. Let us consider these processes in more deatils assuming that the magnetic field is present in the wind. Relativistic particles moving in the magnetic field usually produce synchrotron radiation. However, if the wind is cold (in this case the velocity of particles coincides with the bulk motion velocity) these particles do not produce synchrotron radiation. However, they can produce gamma-radiation with a specific sharp spectral feature through the IC scattering \citep{bogovalov00}. An example of such a system is a cold pulsar wind which does not produce synchrotron radiation in the pre-shock region. In such a system the ``standart" relation between synchrotron and IC radiation components is violated. Remarkably, the possibility for the formation of relativistic flows in the post-shock region, as revealed in this paper, shows that in binary pulsar systems we should expect a ``non-standard'' relation between synchrotron and IC radiation components in post-shock region as well. Moreover, due to the large bulk motion Lorentz factor we should expect strong modulation of the observed nonthermal radiation of electrons. Indeed in the case of binary-pulsar systems the direction of the post shock flow varies with the motion of the pulsar along the orbit around the star. This implies significant changes of the Doppler factor, $\delta$, given the large value of the Lorentz factor. Namely, $\delta \ll 1$ for large viewing angles $\phi$ (e.g. close to $90^\circ$) and $\delta \gg 1$ for small viewing angles. Correspondingly this will have a strong impact on the lightcurve of nonthermal radiation of electrons, $F_\gamma \propto \delta^{n}$ where typically $n \geq 3$. This effect, in particular, can naturally explain the interesting feature of the nonthermal emission of PSR~1259/SS2883, the both synchrotron and inverse Compton components of which disappear during the periastron passage of the pulsar (see e.g. \cite{neronov}). This interesting issue will be discussed elsewhere. | 7 | 10 | 0710.1961 |
0710 | 0710.3363_arXiv.txt | {% Very high energy (VHE) $\gamma$-ray observations have proven to be very successful in localizing Galactic acceleration sites of VHE particles. Observations of shell-type supernova remnants have confirmed that particles are accelerated to VHE energies in supernova blast waves; the interpretation of the $\gamma$-ray data in terms of hadronic or leptonic particle components in these objects relies nevertheless strongly on input from X-ray observations. The largest identified Galactic VHE source class consists of pulsar wind nebulae, as detected in X-rays. Many of the remaining VHE sources remain however unidentified until now. With X-ray observations of these enigmatic ``dark'' objects one hopes to solve the following questions: What is the astrophysical nature of these sources? Are they predominantly electron or hadron accelerators? And what is their contribution to the overall cosmic ray energy budget? The paper aims to provide an overview over the identification status of the Galactic VHE source population.} | Ground-based Cherenkov telescopes detect cosmic $\gamma$-rays in the {\em Very High Energy} (VHE, 100\,GeV - 100\,TeV) domain. In this frequency range, sources are visible in which charged particles are accelerated to TeV energies and beyond; those particles give rise to the detected $\gamma$-ray emission. The energetic particles can be confined inside the objects, close to the actual acceleration site, like in young \linebreak% shell-type supernova remnants (SNRs). In other cases, the particles diffuse out into the surrounding medium after their acceleration to high energies; in these cases, the size of the emitting region itself defines the extent of the ``source'' or object, like for example in a pulsar wind nebula (PWN). The truly diffuse component of the high energy cosmic ray particles is predominantly being traced in the MeV-GeV domain, accessible to $\gamma$-ray satellites. Lower energy diffuse {\em electrons} are also traced through synchro\-tron emission in the radio band. At VHE energies, localized sources dominate because of their intrinsically harder particle spectra compared to the diffuse component. A successful VHE survey of the Galactic plane became possible with the high sensitivity and large field of view (FoV) of the H.E.S.S.\ (High Energy Stereoscopic System) Cherenkov telescope system, with $F_{\mathrm{min}}(>100\,\mathrm{GeV}) \sim 4 \times 10^{-12}\mathrm{erg\,cm^{-2}s^{-1}}$ for a $5\,\sigma$ point-source detection in 25 hrs, and a FoV of 3$^{\circ}$ FWHM (Aharonian et al. 2006c). The H.E.S.S. telescope system became operational end of 2003, and provides through its location in Namibia an excellent view of the Galactic center region. The majority of the currently over 50 known Galactic VHE sources was discovered in the H.E.S.S. survey of the Galactic plane (Aharonian et al. 2005h, 2006g; Hoppe et al. 2007; Kosack et al. 2007), or in the H.E.S.S. FoV of observations of other targets in the Galactic plane (e.g. Aharonian et al. 2005c, 2007b). Perhaps somewhat unexpectedly, many of the new \linebreak% VHE-emitting sources can not readily be identified with \linebreak% known astrophysical objects. Therefore, Galactic VHE astronomy does not only deal with the identification of particle acceleration mechanisms and efficiencies in well-known astrophysical objects, but also with the identification of \linebreak% sources which are so far ``only'' defined through their $\gamma$-ray emission. Follow-up observations of those sources with highly sensitive X-ray instruments such as {\it XMM-Newton} provide the very promising possibility to identifiy those \linebreak% enigmatic sources of high energy particles in our Galaxy. | 7 | 10 | 0710.3363 |
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0710 | 0710.3649_arXiv.txt | We examine the use of the CMB's TE cross correlation power spectrum as a complementary test to detect primordial gravitational waves (PGWs). The first method used is based on the determination of the lowest multipole, $\ell_0$, where the TE power spectrum, $C_{\ell}^{TE}$, first changes sign. The second method uses Wiener filtering on the CMB TE data to remove the density perturbations contribution to the TE power spectrum. In principle this leaves only the contribution of PGWs. We examine two toy experiments (one ideal and another more realistic) to see their ability to constrain PGWs using the TE power spectrum alone. We found that an ideal experiment, one limited only by cosmic variance, can detect PGWs with a ratio of tensor to scalar metric perturbation power spectra $r=0.3$ at $99.9\%$ confidence level using only the TE correlation. This value is comparable with current constraints obtained by WMAP based on the $2\sigma$ upper limits to the B-mode amplitude. We demonstrate that to measure PGWs by their contribution to the TE cross correlation power spectrum in a realistic ground based experiment when real instrumental noise is taken into account, the tensor-to-scalar ratio, $r$, should be approximately three times larger. | Primordial gravitational waves (PGWs) polarize the cosmic microwave background (CMB) (see for example \cite{basko, Polnarev85, Crittenden93, Frewin94, colefrewinpolnarev95, kamionkowski97, Seljak97Pol, selzal97, Kamionkowski98, baskaran06, Keating2006}). Current experiments are using the polarization of the CMB to search for this PGW background (\cite{CloverIntro, BowdenOpt, BicepPaper}. This polarization can be used as a direct test of inflation. An alternative probe of the inflationary epoch which does not use the PGW background was studied by (\cite{sperzal97}). This probe was used in recent analyses by the WMAP team (\cite{peiris03}) to provide plausibility for the inflationary paradigm. This paper presents a test similar in spirit to that of \cite{sperzal97}. CMB polarization can be separated into two independent components: E-mode (grad) polarization and B-mode (curl) polarization. B-mode polarization can only be generated by PGWs (see for example \cite{Seljak97Pol, selzal97, Kamionkowski98}), therefore most CMB polarization experiments which are searching for evidence of PGWs focus on measuring the BB power spectrum. However the TE cross correlation power spectrum offers another method to detect PGWs (\cite{crittenden94}). The TE power spectrum is two orders of magnitude larger than the BB power spectrum and it was suggested that it may therefore be easier to detect gravitational waves in the TE power spectrum (\cite{baskaran06, grishchuk07}). In this paper we first discuss the method of detection of PGWs by measuring the TE power spectrum for low $\ell$. This method, originally proposed in \cite{baskaran06}, is based on a measurement of $\ell_0$, the multipole where the TE power spectrum first changes sign. Hereafter we will call this method ``the zero multipole method''. The TE power spectrum due to density perturbations is positive on large scales, corresponding to $\ell < \ell_0$, changes sign at $\ell = \ell_0$, and then oscillates for $\ell > \ell_0$, while for PGWs the TE power spectrum must be negative for small $\ell$ and then also oscillates for larger $\ell$. The current best set of cosmological parameters, obtained in \cite{Spergel2007WMAP}, gives, in the absence of PGWs, $\ell_0 = 53$. Therefore, the measurement of the difference between the multipole number, $\ell_0$, where the TE power spectrum changes sign, and $\ell=53$ is the way to detect PGWs. We will then consider an alternative method based on Wiener filtering, removing the contribution to the TE power spectrum due to density perturbations. Since the TE power spectrum due to PGWs is megative on large scales a test of negativity of the resulting TE power spectrum is a test of PGWs. In this paper, we present an analysis of both of these methods, based on Monte Carlo simulations. At the present time, the main priority and the main challenge in CMB polarization observations is the detection of the PGW background via the BB power spectrum. In connection with BB experiments, the methods based on the TE cross correlation can be considered as very useful auxiliary measurements of PGWs because systematic effects in TE measurements are not degenerate with those in BB measurements. For example, T/B leakage or even E/B leakage could swamp a detection of BB, whereas T/E leakage would be small and well controlled (see \cite{meir07}). These BB systematics could falsely indicate a detection of PGWs, but measurements of the TE power spectrum provide insurance against such a spurious detection. Additionally, galactic foreground contamination affects BB and TE in different ways, which enables us to perform powerful cross-checks and subtraction of foregrounds in BB measurements. Another advantage of TE measurements for experiments which measure a small fraction of the sky, is related to the fact that a significant contaminant to the B modes is caused by E/B mixing. This limits the power spectrum of PGWs that can be detected (\cite{challinor05}). The E-modes are practicly unaffected by E/B mixing so, in contrast to the BB measurements, the TE power spectrum should be nearly the same for both full and partial sky measurements. The plan of this paper is the following. In Section \ref{physmath}, we introduce the primordial power spectra of scalar (density) and tensor (PGW) perturbations (\ref{primspectra}). Then following \cite{crittenden94} and \cite{baskaran06}, we explain why the sign of the TE power spectra for scalar and tensor perturbations is opposite for large scales (\ref{TECross}). In Section \ref{ell0det}, we describe in more detail the zero multipole method for the detection of PGWs. In Section \ref{wienfilt}, we describe the method for detection of PGWs based on Wiener filtering along with the statistical tests used and a comparison of the tests. In Section \ref{resul}, we present results of numerical Monte Carlo simulations for two toy experiments. In the first toy experiment we neglect instrumental noise and the uncertainties are limited only by cosmic variance (\ref{toy1res}). In the second toy experiment, along with cosmic variance, we take into account instrumental noise which is comparable to real noise in current ground experiments (\ref{toy2res}). For comparison, we also present results of simulations for the two satellite experiments, WMAP (\ref{wmapres}) and Planck (\ref{planckres}). In Section \ref{comptebb}, we compare the the signal-to-noise ratio of the TE measurements with those of BB measurements. | The measurement of where the TE cross correlation first changes sign can be used to detect or put constraints on PGWs. Such constraints are not as strong as the ones given by measurements of the BB power spectrum, however it is useful to have a supplementary method to detect PGWs. We have shown how well the TE mode can constrain the amount of PGWs from just a measurement of the angular scale where it first changes sign for two different toy experiments and two real satellite experiments. The absolute best limit with which we can measure $\ell_0$ only gives us less than a $3\sigma$ detection of the PGW component if $r=0.3$. The current confidence limits gives us $r < 0.3$ at $95\%$ confidence level. Current and future experiments are optimized to measure the BB power spectrum if $r \le 0.1$ even in the presence of foregrounds, which are not taken into account in this paper. Future satellite experiments should be able to detect $r<0.01$ which is $10$ times better than the sensitivity to $r$ than the result of the ideal experiment. If one neglects even cosmic variance, the discreteness of $\ell$ limits the calculation of $\ell_0$, and the sensitivity to $r$, to values considerably larger than $0.01$. The cosmic variance is largest at low $\ell$ and is proportional to the total power spectrum. Since the TE cross correlation has contributions from density perturbations the errors in the measured TE power spectrum make detecting deviations of $\ell_0$ from $53$ difficult, though they also provide insurance against a false detection or imperfect subtraction of instrumental and foreground systematic effects. The other method described in this paper is one in which we filter out the signal due to density perturbations, leaving only the contribution to the TE power spectrum due to PGWs. We then test the resulting TE power spectrum to see if it is negative. Three different statistical tests were used to see if there was a significant detection of PGWs. The $S/N$ test can give a value for $r$ using a comparison with Monte Carlo simulations, while the Wilcoxon rank sum test can only give an allowable range for $r$. The sign test will only tell us if $r \not= 0$. Using the Wiener filtering method, we are unable to make as significant of a detection as using the zero multipole method. The best result was for the $S/N$ test which would give a $2.3\sigma$ detection of $r=0.3$. To detect PGWs on the level of $3\sigma$, the tensor-to-scalae ratio $r$ should be $r \ge 0.4$. The sign test would give $2\sigma$ detection for $r=0.3$ and a $3\sigma$ detection for $r=0.45$. The Wilcoxon ranked sum test gives only a $1.2\sigma$ detection for $r=0.3$ and a $3\sigma$ detection for $r=0.7$. Similar results were gotten for the other three experiments tested. Thus in the sense of potential to detect PGWs, the zero multipole method is the best, next best is the $S/N$ test, then the sign test, and the worst is the Wilcoxon ranked sum test. \cite{baskaran06} present illustrative examples in which high $r$ is consistent with measured TT, EE, and TE correlations. The value of $r$ is so high in these examples that if PGWs with such $r$ really existed, current BB experiments would already detect PGWs. All models predict that the TE cross correlation power spectrum change sign only once for $\ell < 100$. The fact WMAP cannot exclude several multipoles with $C_{\ell}^{TE} > 0$ in between multipoles of $C_{\ell}^{TE} < 0$ means that the TE cross correlation power spectrum either changes sign several times for $\ell < 100$ or there is some instrumental noise which causes some anticorrelation measurements. Using instrumental noise consistent with WMAP, our Monte Carlo simulations give $\Delta \ell_0 \approx 16$ and $\ell_0 > 40$, which means that there is no evidence of PGWs in the TE correlation power spectrum. | 7 | 10 | 0710.3649 |
0710 | 0710.4296_arXiv.txt | { Weakly interacting massive particles (WIMPs) are one of the leading candidates for Dark Matter. So far we can use direct Dark Matter detection to estimate the mass of halo WIMPs only by fitting predicted recoil spectra to future experimental data. Here we develop a model--independent method for determining the WIMP mass by using experimental data directly. This method is independent of the as yet unknown WIMP density near the Earth as well as of the WIMP--nuclear cross section and can be used to extract information about WIMP mass with ${\cal O}(50)$ events. \PACS{ {95.35.+d}{Dark Matter} \and {29.85.Fj}{data analysis} } % } % | \label{intro} Today astrophysicists have strong evidence \cite{evida}-\cite{bullet} to believe that a large fraction (more than 80\%) of the matter in the Universe is dark (i.e., interacts at most very weakly with electromagnetic radiation). The dominant component of this cosmological Dark Matter (DM) must be due to some yet to be discovered, non--baryonic particles. Weakly interacting massive particles (WIMPs) $\chi$ with masses between 10 GeV and a few TeV are one of the leading candidates for DM \cite{susydm}. Currently, the most promising method to detect different WIMP candidates is the direct detection of the recoil energy deposited in a low--background laboratory detector by elastic scattering of ambient WIMPs off the target nuclei \cite{detaa}, \cite{detab}. The differential rate for elastic WIMP--nucleus scattering is given by \cite{susydm}: \beq \label{eqn2101} \dRdQ = \calA \FQ \intvmin \bfrac{f_1(v)}{v} dv\, . \eeq Here $R$ is the direct detection event rate, i.e., the number of events per unit time and unit mass of detector material, $Q$ is the energy deposited in the detector, $F(Q)$ is the elastic nuclear form factor, and $f_1(v)$ is the one--dimensional velocity distribution function of the WIMPs impinging on the detector. The constant coefficient $\calA$ is defined as \beq \label{eqn2102} \calA \equiv \frac{\rho_0 \sigma_0}{2 \mchi m_{\rm r,N}^2}\, , \eeq where $\rho_0$ is the WIMP density near the Earth and $\sigma_0$ is the total cross section ignoring the form factor suppression. The reduced mass $m_{\rm r,N}$ is defined by \beq \label{eqn2103} m_{\rm r,N} \equiv \frac{\mchi \mN}{\mchi+\mN}\, , \eeq where $\mchi$ is the WIMP mass and $\mN$ that of the target nucleus. Finally, $\vmin$ is the minimal incoming velocity of incident WIMPs that can deposit the energy $Q$ in the detector: \beq \label{eqn2104} \vmin = \alpha \sqrt{Q}\, , \eeq where we define \beq \label{eqn2105} \alpha \equiv \sfrac{\mN}{2 m_{\rm r,N}^2}\, . \eeq So far most theoretical analyses of direct WIMP detection have predicted the detection rate for a given (class of) WIMP(s), based on a specific model of the galactic halo. This can be used to estimate the mass of halo WIMPs only by fitting the predicted recoil spectra to future experimental data, e.g., \cite{Green07}. The goal of our work is to develop methods which allow to extract information on halo WIMPs by using the experimental data directly. In our earlier work \cite{DMDD} we used a time--averaged recoil spectrum, assumed that no directional information exists, and derived an expression for estimating moments of the normalized one--dimensional velocity distribution function of halo WIMPs: \beqn \label{eqn2208} \expv{v^n} \eqnequiv \int_{v_{\rm min}(\Qthre)}^\infty v^n f_1(v) \~ dv \non\\ \= \alpha^n \bfrac{2 \Qthre^{(n+1)/2} \rthre+(n+1) I_n \FQthre}{2 \Qthre^{1/2} \rthre+I_0 \FQthre}\, , \non\\ \eeqn where $Q_{\rm thre}$ is the threshold energy of the detector, $\rthre \equiv (dR/dQ)_{Q = \Qthre}$ is an estimated value of the scattering spectrum at $Q = \Qthre$. $I_n$ can be either determined from a given expression (e.g., a fit to data) for the recoil spectrum: \beq \label{eqn2110} I_n = \int_{\Qthre}^{\infty} \frac{Q^{(n-1)/2}}{\FQ} \adRdQ \~ dQ\, , \eeq or estimated directly from the measured recoil energy: \beq \label{eqn2210} I_n = \sum_a \frac{Q_a^{(n-1)/2}}{F^2(Q_a)}\, , \eeq where the sum runs over all events in the data set. Note that all these expressions are independent of the as yet unknown WIMP density near the Earth as well as of the WIMP--nucleus cross section. (More details about the reconstruction of the velocity distribution function of halo WIMPs and the estimate of its moments as well as all formulae needed can be found in Ref.~\cite{DMDD}.) | In this paper we have presented a method which allows to extract information on the WIMP mass from the recoil energy measured in elastic WIMP--nucleus scattering experiments directly. In the long term this information can be used to constrain e.g., SUSY models in the elementary particle physics and compare with information from future collider experiments. Our method for determining the WIMP mass by combining two (or more) experiments with different detector materials is independent of the as yet unknown WIMP density near the Earth as well as of the WIMP--nucleus cross section. The only information which one needs is the measured recoil energy. Due to the maximal measuring energy of the detector, there will be a deviation of the reproduced WIMP mass from the true one. Nevertheless, for experiments with very few events and thus a quite large statistical error in the near future, this deviation should not affect the reproduced WIMP mass very significantly. Moreover, the numerical analysis shows also that, for WIMP masses $\le 100~{\rm GeV}/c^2$ some meaningful information on the WIMP mass can already be extracted from ${\cal O}(50)$ total (each experiment ${\cal O}(25)$) events. The analyses of this work are based on several simplified assumptions. First, all experimental systematic uncertainties, as well as the uncertainty on the measurement of the recoil energy have been ignored. Comparing with large statistical uncertainty this should be a quite good approximation. Second, the analysis treats each recorded event as signal, i.e., background has been ignored altogether. This may in fact not be unrealistic for modern detectors. Third, each detector consists of a single isotope. This is quite realistic for the % semiconductor % detectors. For detectors containing more than one nucleus, by simultaneously measuring two signals, one might be able to tell on an event--by--event basis which kind of nucleus has been struck. \subsubsection* | 7 | 10 | 0710.4296 |
0710 | 0710.0561_arXiv.txt | {% We review the most important findings on AGN physics and cosmological evolution as obtained by extragalactic X--ray surveys and associated multiwavelength observations. We briefly discuss the perspectives for future enterprises and in particular the scientific case for an extremely deep (2--3 Ms) XMM survey} | Since their launch in 1999, both XMM--{\it Newton} and {\it Chandra} have performed a large ($>$ 30) number of surveys covering a wide fraction of the area vs. depth plane (see fig~1 in Brandt \& Hasinger 2005). Thanks to vigorous programs of multiwavelength follow-up observations which, since a few years, have became customary, our understanding of AGN evolution has received a major boost. The high level of completeness in redshift determination for a large number of X--ray selected AGN (up to a few thousands) has made possible a robust determination of the luminosity function and evolution of unobscured and mildly obscured AGN (Ueda et al. 2003; Hasinger et al. 2005; La Franca et al. 2005), which turned out to be luminosity dependent: the space density of bright QSO ($L_X > 10^{44}$ erg s$^{-1}$) peaks at z$\sim$ 2--3, to be compared with the z$\sim$0.7--1 peak of lower luminosity Seyfert galaxies. The fraction of obscured AGN is also strongly dependent from the X--ray luminosity (Ueda et al. 2003; La Franca et al. 2005). Even though the shape and the normalization of the function describing the obscured fraction vs. luminosity is still matter of debate, it is clear that absorption is much more common at low X--ray luminosities. Such a trend has been observed also in the optical (Simpson 2005) and in the near infrared (Maiolino et al. 2007) and may be linked to the AGN radiative power which is able to ionize and expel gas and dust from the nuclear regions. More debated is the claim of an increase of the obscured fraction towards high redshifts. First suggested by La Franca et al. (2005), it has been confirmed by Treister \& Urry (2006), while it is not required in the models discussed by Ueda et al. (2003) and Gilli et al. (2007). A redshift dependence of the obscuring fraction would be naturally explained in the current framework of AGN formation and evolution (see, for example, Hopkins et al. 2006): the anti--hierarchical behaviour observed in AGN e\-vo\-lution (similar to that observed in normal galaxies), along with several other evidences, suggests that super\-massive bl\-ack holes (SMBH) and their host galaxies co--evolve and that their formation and evolution are most likely different aspects of the same astrophysical problem. At early times large quantities of cold gas were available to efficiently feed and obscure the growing black holes. Later on, the ionizing nuclear flux is able to "clean" its environment appearing as an unobscured QSO. However, the picture sket\-ched above is likely to be much more complicated, depending on many other parameters (such as the BH mass, the Eddington ratio, the QSO duty cicle) and may not necessarily result in an increasing fraction of obscuration towards high redshifts. It should also be remarked that sensitive X--ray observations are highly efficient to unveil weak and/or elusive accreting black holes which would be missed, or not classified as such, by surveys at longer wavelengths. Among them XBONG (Comastri et al. 2002), sources with high X--ray to optical flux ratio (Fiore et al. 2003), Extremely Red Objects (Brusa et al. 2005), Sub Millimeter Galaxies (Alexander et al. 2005). Although they are probably not representative of the sources of the X--ray background their study has allowed us to better understand the physics of accreting black holes. Finally, the coverage of several {\it Chandra} and XMM fie\-lds has allowed to uncover several redshifts spikes in the distribution of X--ray sources (Gilli et al. 2005) and demonstrated that AGN may be used as reliable tracers of the large scales structures. The underlying large scale structure may play an important role in triggering AGN activity. According to the analysis of a sample of X--ray selected AGN in the AEGIS survey, Georgakakis et al. (2006) suggest that $z \sim$ 1 AGN are more frequent in dense environments. | 7 | 10 | 0710.0561 |
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0710 | 0710.0757_arXiv.txt | Analysis of the hydrodynamic and helioseismic effects in the photosphere during the solar flare of July 23, 2002, observed by Michelson Doppler Imager (MDI) on SOHO, and high-energy images from RHESSI shows that these effects are closely associated with sources of the hard X-ray emission, and that there are no such effects in the centroid region of the flare gamma-ray emission. These results demonstrate that contrary to expectations the hydrodynamic and helioseismic responses (''sunquakes") are more likely to be caused by accelerated electrons than by high-energy protons. A series of multiple impulses of high-energy electrons forms a hydrodynamic source moving in the photosphere with a supersonic speed. The moving source plays a critical role in the formation of the anisotropic wave front of sunquakes. | ``Sunquakes", the helioseismic response to solar flares, are caused by strong localized hydrodynamic impacts in the photosphere during the flare impulsive phase. The helioseismic waves are observed directly as expanding circular-shaped ripples in SOHO/MDI Dopplergrams, which can be detected in Dopplergram movies and as a characteristic ridge in time-distance diagrams, \citep{Kosovichev1998, Kosovichev2006a}, or indirectly by calculating integrated acoustic emission \citep{Donea1999, Donea2005}. Solar flares are sources of high-temperature plasma and strong hydrodynamic motions in the solar atmosphere. Perhaps, in all flares such perturbations generate acoustic waves traveling through the interior. However, only in some flares the impact is sufficiently localized and strong to produce the seismic waves with the amplitude above the convection noise level. It has been established in the initial July 9, 1996, flare observations \citep{Kosovichev1998} that the hydrodynamic impact follows the hard X-ray flux impulse, and hence, the impact of high-energy electrons. Nevertheless, a common paradigm is that the sunquake events are caused by accelerated protons because protons carry more momentum and penetrate deeper into the solar atmosphere than electrons, which loose most of their energy in the upper chromosphere. This paradigm is not easy to test because the gamma-ray emission, which indicates the presence of high-energy protons, is rarely observed. In a large X17 flare of October 28, 2003, the gamma-ray emission observed by RHESSI was located close to the hard X-ray sources and two of the three places of the phototospheric impacts (sunquake sources) \citep{Kosovichev2006a}. Because of the close locations of the hard X-ray and gamma-ray sources these observations could not exclude the possibility of the proton or mixed electron-proton impacts \citep{Zharkova2007}. However, in one event, X4.8 flare of July 23, 2002, the hard X-ray and gamma-ray sources were significantly separated from each other. The centroid of the $\gamma$-ray 2.233 MeV neutron-capture emission was found to be displaced by $20"\pm 6"$ (with 5-sigma confidence) from that of the $0.3-0.5$ MeV X-ray emission implying a difference in acceleration and/or propagation between the accelerated electrons and ions \citep{Hurford2003}. Therefore, this flare provides a unique opportunity to investigate the photospheric and helioseismic responses separately for high-energy electrons and protons. In this Letter, I present results of the analysis of the relationship between the hard X-ray and gamma-ray emissions and the hydrodynamic and seismic signals in the photosphere, using data from RHESSI \citep{Lin2002} and MDI on SOHO \citep{Scherrer1995}. RHESSI provides X-ray/gamma-ray imaging spectroscopy from 3 keV to 17 MeV with angular resolution $2.3''-3'$ ($35''$ at gamma-ray energies) over the full Sun. MDI measures the Doppler velocity and the line-of-sight magnetic field of the photospheric plasma every minute with 2 arcsec/pixel resolution also over the full Sun. | The analysis of RHESSI X-ray and gamma-ray images and SOHO/MDI Dopplergrams of the July 23, 2002, X4.8 solar flare revealed that the hydrodynamic and seismic responses are closely associated the hard X-ray emission, both spatially and temporally, but showed no significant responses in the gamma-source centroid area. Because this flare was one of strongest gamma-flares, and the hard X-ray and gamma-ray sources were separated, these observations show that the accelerated protons are unlikely to be a source of the hydrodynamic response and sunquakes. Furthermore, the detailed analysis of the dynamics of sunquake sources in this Letter and in the paper by \citet{Kosovichev2006a} reveals their close association with expanding flare ribbons and rapid HXR source motion along the ribbons, and, thus, with the magnetic reconnection process. The fast motion of these sources results in strong anisotropy of the seismic waves, clearly observed in the MDI data. The general picture that comes from the analysis of MDI and RHESSI data is consistent with the previously developed hydrodynamic thick-target model, illustrated in Fig.~\ref{fig2} \citep{Kostiuk1975, Livshits1981, Fisher1985, Kosovichev1986}. In this model, high-energy electrons heat the upper chromosphere to high temperatures generating a high-pressure region, expansion of which causes evaporation of the chromospheric plasma and a high compression shock. The shock reaches the photosphere and excites the seismic waves. However, the new results show that it is important to include effects of the multiple impact and moving source in the thick-target and sunquake models. The photospheric and helioseismic effects observed during the impulsive phase of solar flares are closely related to the processes of acceleration and propagation of electrons and ions, and may provide new important information about these processes. \clearpage | 7 | 10 | 0710.0757 |
0710 | 0710.4518_arXiv.txt | We investigate the birth and evolution of isolated radio pulsars using a population synthesis method, modeling the birth properties of the pulsars, their time evolution, and their detection in the Parkes and Swinburne Multibeam (MB) surveys. Together, the Parkes and Swinburne MB surveys \citep[][]{2001MNRAS.328...17M, 2001MNRAS.326..358E} have detected nearly 2/3 of the known pulsars and provide a remarkably homogeneous sample to compare with simulations. New proper motion measurements \citep[][]{2002ApJ...571..906B, 2003AJ....126.3090B} and an improved model of the distribution of free electrons in the interstellar medium, NE2001 \citep[][]{NE2001}, also make revisiting these issues particularly worthwhile. We present a simple population model that reproduces the actual observations well, and consider others that fail. We conclude that: pulsars are born in the spiral arms, with the birthrate of $2.8\pm0.5$ pulsars/century peaking at a distance $\sim3$ kpc from the Galactic centre, and with mean initial speed of $380^{+40}_{-60}$ km s$^{-1}$; the birth spin period distribution extends to several hundred milliseconds, with no evidence of multimodality, implying that characteristic ages overestimate the true ages of the pulsars by a median factor >2 for true ages <30,000 yr; models in which the radio luminosities of the pulsars are random generically fail to reproduce the observed $P-\dot{P}$ diagram, suggesting a relation between intrinsic radio luminosity and $(P, \dot{P})$; radio luminosities $L \propto \sqrt{\dot{E}}$ provides a good match to the observed $P-\dot{P}$ diagram; for this favored radio luminosity model, we find no evidence for significant magnetic field decay over the lifetime of the pulsars as radio sources ($\sim100$ Myr). | From soon after the discovery of the pulsars \citep[][]{1968Natur.217..709H}, their Galactic population has been the focus of numerous studies \citep[][]{1970ApJ...160..979G, 1971IAUS...46..165L, 1977ApJ...215..885T, 1977MNRAS.179..635D, 1985MNRAS.213..613L, 1987A&A...171..152S, 1989ApJ...345..931E, 1992A&A...254..198B, 1993MNRAS.263..403L, 1997A&A...322..127H, 1997MNRAS.289..592L, 2002ApJ...568..289A, 2002ApJ...565..482G, 2004ApJ...604..775G}. Nevertheless, in spite of much progress, many outstanding questions remain. What are the birth positions, velocities, spin periods, and magnetic fields of the pulsars? How do these evolve in time and how are they related to the radio luminosities of the pulsars? One approach to answering these questions is to make use of the statistical power of the growing pulsar catalogue to study the pulsar population as a whole.\\ \\ Many recent advances in pulsar astronomy make it particularly worthwhile to revisit the above questions through population synthesis. The recently completed Parkes and Swinburne Multibeam (PMB and SMB; \cite{2001MNRAS.328...17M, 2001MNRAS.326..358E}) surveys have detected nearly 2/3 of the known pulsars and provide a large and remarkably homogeneous observed sample to compare with simulations. New proper motion measurements \citep[][]{2002ApJ...571..906B, 2003AJ....126.3090B} and an improved model of the interstellar medium \citep[][]{NE2001} also provide valuable new information.\\ \\ In this work, the details of which have been reported by Faucher-Gigu\`ere \& Kaspi (2006) \cite{2006ApJ...643..332F}, we investigate the birth properties of Galactic isolated radio pulsars and their time evolution. To do so, we generate an ensemble of mock galaxies populated by pulsars with prescribed birth properties (spatial locations, velocities, spin periods, radio luminosities, magnetic fields) and evolve the pulsars in time using physical models. We model the selection function of the PMB and SMB surveys using a modified version of the radiometer equation \citep[][]{1985ApJ...294L..25D} and apply it to our mock galaxies. We compare the observed histograms of Galactic longitudes, latitudes, dispersion measures, 1.4 GHz radio fluxes, pulse periods, magnetic fields, as well as the observed $P-\dot{P}$ diagrams, to judge how well each population model reproduces the actual observations. \begin{figure*} \includegraphics[height=.25\textheight]{ppdot_diagrams.eps} \caption{Comparison of the $P-\dot{P}$ diagram for the pulsars detected in the Parkes and Swinburne Multibeam surveys (left) with the corresponding diagrams in our best simulation with $L\propto \sqrt{\dot{E}}$ (middle) and with a simulation in which the radio luminosities of the pulsars are uncorrelated with their other characteristics (``random''; right). While the simulation with $L\propto \sqrt{\dot{E}}$ reproduces the observed diagram remarkably well, the model with random radio luminosities produces a pile-up of pulsars near the death line that is in clear disagreement with the observations.} \label{ppdot diagrams} \end{figure*} | 7 | 10 | 0710.4518 |
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0710 | 0710.0494_arXiv.txt | We study the color structure of disk galaxies in the Groth strip at redshifts $0.1<z<1.2$. Our aim is to test formation models in which bulges form before/after the disk. We find smooth color distributions with gentle outward blueing across the galaxy image: bulges are not distinctly redder than their disks; and bulge colors strongly correlate with global colors. The results suggest a roughly coeval evolution of bulges and disks. About 50\% of the nuclei of galaxies with central light excesses above the outer exponential profile hold passively evolving red populations. The remainder 50\% are galaxies with central blue colors similar to their disks. They may be bulges in formation, or the central parts of disks with non-exponential surface brightness profiles. | This work addresses the simple question: \textsl{are bulges older than their host disks?}. The hope is to be able to falsify one or more of the hypothesis for bulge formation: bulges before disks, via mergers or primordial collapse, \textit{vs.}\ bulges after disks, from disk instabilities or other secular processes. The plan is admitedly na\"ive. A prevalence of blue bulges would argue against an old formation age for bulges. But bulges redder than their disks would not rule out the disk instability model, as the younger disk age might come from subsequent star formation in the disk after the bulge is formed. However, the question is a useful guideline, especially when combined with structural and isophotal diagnostics. For early- to intermediate-type disk galaxies in the nearby Universe, \cite[Peletier \& Balcells (1996)]{Peletier96} concluded that the age differences between bulges and disks was in most cases consistent with zero. But age dating of stellar populations has large uncertainties. We have repeated the Peletier \& Balcells study at redshifts up to $z=1$, closer to the expected formation epoch of bulges. An extensive analysis is presented in \cite[Dom\'{\i}nguez-Palmero \& Balcells (2008)]{DominguezPalmero08}. | 7 | 10 | 0710.0494 |
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0710 | 0710.2491_arXiv.txt | We present optical, IR and millimeter observations of the solar-type star 13-277, also known as GM Cep, in the 4 Myr-old cluster Tr 37. GM Cep experiences rapid magnitude variations of more than 2 mag at optical wavelengths. We explore the causes of the variability, which seem to be dominated by strong increases in the accretion, being similar to EX-or episodes. The star shows high, variable accretion rates (up to $\sim$10$^{-6}$ M$_\odot$/yr), signs of powerful winds, and it is a very fast rotator (V$sini \sim$43 km/s). Its strong mid-IR excesses reveal a very flared disk and/or a remnant envelope, most likely out of hydrostatic equilibrium. The 1.3 millimeter fluxes suggest a relatively massive disk (M$_D \sim$0.1 M$_\odot$). Nevertheless, the millimeter mass is not enough to sustain increased accretion episodes over large timescales, unless the mass is underestimated due to significant grain growth. We finally explore the possibility of GM Cep having a binary companion, which could trigger disk instabilities producing the enhanced accretion episodes. | } The open cluster Tr 37 is one of the best studied intermediate-aged young star formation regions. Aged $\sim$ 4 Myr (Sicilia-Aguilar et al. 2004, 2005; hereafter Paper I, Paper II), and located at 900 pc distance (Contreras et al. 2002), it contains a rich population of T Tauri stars (TTS; $\sim$180 members with spectral types G to M2) of which about 48\% still show IR excesses consistent with protoplanetary disks at different evolutionary stages (Sicilia-Aguilar et al. 2006a, b; hereafter Paper III, Paper IV). While most of the stars in Tr 37 show important evidence of disk evolution, with lower accretion rates and near-IR excesses than in younger regions, a few objects still display characteristics of much younger systems. The most remarkable one is the solar-type star 13-277, also called GM Cep (Morgenroth 1939). Because of its spatial location within the cluster, it belongs most likely to the Tr 37 main population, with average ages $\sim$4 Myr ($\sim$85\% of the population is older than 2 Myr, and $\sim$95\% is older than 1 Myr), rather than to the young population associated to the Tr 37 globule, aged $\sim$1 Myr (Paper II). Its continuum spectrum and the strong and broad H$\alpha$ emission suggested an accretion rate \.{M}$\sim$3 10$^{-7}$ M$_\odot$/yr, about 2 orders of magnitude over the median accretion rate of TTS in Tr 37 (Paper IV). Its bolometric luminosity (L$\sim$26 L$_\odot$ in 2000) is about 1 order of magnitude higher than other stars with similar late G-early K spectral type. Finally, we notice an unusually high mid-IR flux for a solar-type star: its MIPS flux at 70 $\mu$m is comparable to that of the Tr 37 Herbig Be star MVA-426, being one of the only four cluster members detected at this wavelength (Paper III). All observations suggest that GM Cep is a variable star of the EX-or type, with an unstable disk and variable accretion rate, which is remarkable in an ``old'' cluster like Tr 37, where disk evolution seems ubiquitous. FU-ors and EX-or objects have been suggested to be either normal stages within the very early TTS evolution (see Hartmann \& Kenyon 1996 for a review) or a special type of young object, probably binary (Herbig 2003 and references therein). Here we present the results of our photometry monitoring campaign during 2006-2007\footnote{Based on observations collected at the German-Spanish Astronomical Center, Calar Alto, jointly operated by the Max-Planck-Institut f\"{u}r Astronomie Heidelberg and the Instituto de Astrof\'{\i}sica de Andaluc\'{\i}a (CSIC)}, lucky imaging, millimeter continuum, and $^{12}$CO data, together with optical spectra (high- and low-resolution), and a compilation of the data available from the literature. The multiwavelength data are described in Section \ref{observations}. In Section \ref{analysis} we explore the characteristics of the star: optical and IR variability, spectral type, accretion, winds, disk mass, and potential companions. Section \ref{outburst} discusses the possible causes of variability and the comparison with similar stars, and in Section \ref{conclu} we summarize our results. | } GM Cep is an extremely variable late G star in the 4 Myr-old cluster Tr 37, resembling younger EX-or objects. Its complex variability in color and magnitude is not consistent with a single typical variability mechanism in TTS (hot/cold spots, variable obscuration, or changes in the accretion rate). The amplitude of the variations, the high accretion rate, the luminous mid-IR disk, and the high stellar luminosity suggest variable accretion to be the stronger contributor, maybe mixed with variable extinction by circumstellar material (similar to UX-or variables), and some minor influence of cold spots and scattering. Increased accretion episodes are thought to produce the strong, irregular variations in young EX-or and FU-or objects, typically still surrounded by massive infalling envelopes. Nevertheless, GM Cep has a medium-mass disk and a small or inexistent envelope, and belongs to a cluster where strong disk evolution is ubiquitous. Large changes in the accretion rates can result in changes in the stellar and disk structure, and conversely, the disruption of the disk structure (for instance, by companions, or via gravitational instability if the disk is very massive and compact), can increase the accretion rate. The presence of companions (stellar, substellar or planetary) is a plausible mechanism to produce disk instabilities in a relatively old star like GM Cep. Simultaneous multiwavelength observations, including IRS spectra scheduled for 2007, providing coverage in the 5-35$\mu$m range, and detailed monitoring in the coming years will help us to reveal the nature of the variability in GM Cep. Sub-millimetre observations in the future should be used to determine the SED slope and constrain the size of the grains. Old but extremely accreting stars are thus a key to understand the processes of disk accretion and evolution and the consequences for planet formation. | 7 | 10 | 0710.2491 |
0710 | 0710.5087_arXiv.txt | We report on the results of spectroscopic mapping observations carried out in the nuclear region of Centaurus A (NGC5128) over the 5.2 - 15 and 20-36$\mu$m spectral regions using the Infrared Spectrograph on the {\it Spitzer Space Telescope}. We have detected and mapped S(0), S(2), S(3), and S(5) pure rotational transition lines of molecular hydrogen and emissions in the fine-structure transitions of [SiII], [SIII], [FeII], [FeIII], [ArII], [SIV], [NeII], [NeV] and [OIV]. The 500 pc bipolar dust shell discovered by Quillen et al.(2006) is even more clearly seen in the 11.3$\mu$m dust emission feature than previous broad band imaging. The pure rotational lines of molecular hydrogen other than the S(0) line are detected above the dusty disk and associated with the oval dust shell. The molecular hydrogen transitions indicate the presence of warm gas at temperatures 250--720K. The column density of the warm molecular hydrogen in the shell is $N(H_2) \sim 10^{20}$cm$^{-2}$ and similar to that estimated from the continuum dust shell surface brightness. The ratio of the dust emission features at 7.7$\mu$m and 11.3$\mu$m and the ratio of the [NeII](12.8$\mu$m) and 11.3$\mu$m dust emission feature are lower in the 500 pc dust shell than in the star forming disk. The clearer shell morphology at 11.3$\mu$m, warm molecular hydrogen emission in the shell, and variation in line ratios in the shell compared to those in the disk, confirm spectroscopically that this shell is a separate coherent entity and is unlikely to be a chance superposition of dust filaments. The physical conditions in the shell are most similar to Galactic supernova remnants where blast waves encounter molecular clouds. The lines requiring the highest level of ionization, [NeV]($24.318\mu$m) and [OIV]($25.890\mu$m), are detected 20--25\arcsec north-east and south-west of the nucleus and at position angles near the radio jet axis. Fine structure line ratios and limits from this region suggest that the medium is low density and illuminated by a hard radiation field at low ionization parameter. These higher S molecular hydrogen pure rotational transitions are also particularly bright in the same region as the [OIV] and [NeV] emission. This suggests that the gas associated with the dust shell has been excited near the jet axis and is part of an ionization cone. | The nearest of all the giant radio galaxies, Centaurus~A (NGC~5128) provides a unique opportunity to observe the dynamics and morphology of an active galaxy in detail across the electromagnetic spectrum. For a recent review on this remarkable object see \citet{israel}. In its central regions, NGC~5128 exhibits a well recognized and optically-dark band of absorption across its nucleus. Images from the {\it Spitzer Space Telescope} (SST) with the Infrared Array Camera (IRAC) and Multiple Imaging Photometer for Spitzer (MIPS) cameras in the mid-infrared reveal a 3\arcmin long parallelogram shape \citep{quillen_irwarp} that has been modeled as a series of folds in a dusty warped thin disk (e.g., \citealt{bland86,bland87,nicholson92,sparke96,quillen_irwarp}). The dusty disk is also the site of ongoing star formation at a rate of about $1M_\odot$~yr$^{-1}$, (based on the infrared luminosity estimated by \citealt{eckart90}). Centaurus~A hosts an active nucleus (e.g., \citealt{whysong04,mirabel99}) that has been recently studied using infrared spectra from the SST. \citep{weedman05,gorjian07}. Its nucleus exhibits a strong silicate absorption feature and emission from [NeV]. Here we do not study the nucleus but focus on structure exterior to it. The IRAC and MIPS images of Centaurus~A have revealed another surprise, a bipolar shell-like structure 500\,pc north and south of the nucleus \citep{quillen_supershell}. This shell, seen for the first time with Spitzer imaging, could be the first extragalactic nuclear shell discovered in the infrared. In this paper we present mapping of the central 500--600pc of Centaurus A done with the Infrared Spectrograph (IRS) on board the SST. The infrared spectral maps were obtained to test the possibility that the apparent dust shell is a separate coherent structure and not a chance superposition of dust filaments. They were also obtained to search for possible interactions between the AGN and the interstellar medium either by illumination from the AGN or caused by jets or outflows. Based on the discussion by \citet{israel}, we adopt a distance to Cen A of 3.4~Mpc. At this distance, 1~kpc corresponds to 1\arcmin on the sky. All positions reported in this manuscript are given with respect to Epoch 2000. | We have carried out a spectroscopic study of the central 2 arcminutes of Centaurus A using short low and long high spectral modules of the Infrared Spectrograph on board the {\it Spitzer Space Telescope}. Most of the emission lines detected in the spectral cubes (e.g., [SIII](33.5$\mu$m), [SiII](34.8), [FeII](26.0), [FeIII](23.9), [ArII](6.98) and H$_2$S(0)(28.2)) and dust emission features primarily trace regions of star formation in the warped disk. Our previous study based on IRAC and MIPS imaging suggested that Centaurus A hosts an oval or bipolar dustshell at a distance approximately 500 pc from the nucleus seen above and below the warped disk. This dust shell, if confirmed, would be the first extragalactic shell to be discovered in the infrared. Here we see the dust shell even more clearly and prominently in the 11.3$\mu$m dust emission feature than we saw previously in the broad band IRAC images. We have found variations in the dust emission feature 7.7$\mu$m/11.3$\mu$m ratio and dust 11.3$\mu$m/[NeII](12.8$\mu$) ratio, with the oval dust shell having the lowest ratios compared to the star forming disk. The clearer shell morphology at 11.3$\mu$m than previously seen in broad band images, the association of the molecular hydrogen emission in the shell, and the variation in line ratios in the shell compared to those in the disk, confirm spectroscopically that the shell discovered previously \citep{quillen_supershell} is a separate coherent entity and is unlikely to be a chance superposition of dust filaments. We find evidence for higher ionization species line emission in [NeV]($24.3\mu$m) and [OIV]($25.9\mu$m) near the jet axis. Emission in these two lines is seen both north-east and south-west of the nucleus along position angles $\sim 40^\circ$ and $\sim -120^\circ$. These angles are similar to but not exactly the same as the jet axis at 55$^\circ$ as seen at 5GHz. Outside the nucleus, the peak surface brightness in these lines is 25\arcsec or 400~pc south-west of the nucleus. Emission line ratios and limits at the location of the [OIV] peak suggest that the emitting region is at low ionization parameter, $U \lesssim 10^{-2}$, and has low electron density, $n_e \lesssim 10^2$cm$^{-3}$. We crudely estimate that the AGN can provide sufficient UV photons to account for the [OIV] luminosity 400~pc from the nucleus, as long as UV photons are not absorbed by intervening material as they travel from the AGN to the dust shell. A much more detailed photo-ionization study is required to understand the excitation of the [OIV] and [NeV] emission. Previous reports of an ionization cone in Cen A were based on near-infrared imaging of the central few arcseconds \citep{bryant99}. Subsequent studies interpreted line emission in terms of a disk rather than ionization cone \citep{schreier98,krajnovic06}. Unfortunately the morphology of the near-infrared images is strongly affected by extinction and the warp disk models have not been good enough to accurately predict the extinction in the central few arcsecond of the nucleus. Cen A might be the only active galaxy in which mid-infrared spectroscopy has found evidence for high ionization species such as NeV at hundred pc distances from the nucleus. As far as we know Cen A hosts the only ionization cone that has been resolved with observations from the Spitzer Space Telescope. We see evidence for warm molecular hydrogen coincident with the peak in [OIV] in the odd pure rotational odd transitions S(3) and S(5). The S(7) and S(2) transitions are also detected but at weaker levels. The S(3) and S(5) emission also lies in the vicinity of the dust shell that is most prominent in the 11.3$\mu$m dust (PAH) emission feature. A two temperature component model can fit the rotational line ratios and implies that there is warm molecular hydrogen with temperatures in the range 250--720K. The temperatures are warmer than seen in nucleus of non-active nearby galaxies, similar to those of LINERS and Seyferts but lower than exhibited by supernova remnants ULIRGS (as compared to studies by \citealt{roussel07}, \citealt{higdon06} and \citealt{neufeld07}). Near the jet axis, the column depth of warm molecular hydrogen is $N(H_2) \sim 10^{20}$cm$^{-2}$ similar to that estimated from the infrared continuum emission of the dust shell by \citet{quillen_supershell}. This suggests that gas associated with the dust shell has been heated near the jet axis. There is probably more than one excitation process as [OIV] and [NeV] emission are detected primarily along the jet axis whereas the higher S pure rotational molecular hydrogen lines are detected there and in the vicinity of the dust shell. Previous studies of the pure-rotational molecular hydrogen lines in extra-galactic objects (e.g., \citealt{higdon06,panuzzo07,roussel07,ogle07}) have not well resolved the emission. The association of the warm molecular hydrogen gas with a shell is most similar to phenomena exhibited by Galactic supernova remnants where the blast wave encounters molecular clouds \citep{neufeld07}. The physical conditions estimated from the molecular hydrogen observations are similar in properties (column depth, temperatures and fraction of gas in the two temperature components) to the parameters estimated by \citet{neufeld07} in 4 Galactic supernova remnants. This suggests that theory of interstellar shock waves be applied to interpreting observations of the shell. \citet{neufeld07} associates the warmer molecular hydrogen component responsible for the higher S pure rotational transitions with dissociative shocks. These require shock velocities $\gtrsim 70$km/s (e.g., \citealt{drainemckee93}). A physical scenario and model accounting for the shell's structure and energetics is currently lacking. Deep optical spectroscopic and radio studies are particularly needed to better constrain gas motions and physical conditions in this shell. \vskip 1.0truein We thank Dan Watson, Paul van der Werf, Ralph Kraft, Jacqueline van Gorkom, Martin Hardcastle, and Christine Jones-Forman for helpful discussions and correspondence. We thank Martin Hardcastle for providing us with images of Centaurus A in the radio. Support for this work was in part provided by by NASA through an award issued by JPL/Caltech, National Science Foundation grants AST-0406823 $\&$ PHY-0552695, the National Aeronautics and Space Administration under Grant No.$\sim$NNG04GM12G issued through the Origins of Solar Systems Program, and HST-AR-10972 to the Space Telescope Science Institute. JBH is funded by a Federation Fellowship from the Australian Research Council. \clearpage \begin{figure} \includegraphics[angle=0,width=3.5in]{cont25.eps} \includegraphics[angle=0,width=3.5in]{cont32.eps} \caption{ \label{fig:cont25} Continuum emission from the LH spectral cube from $0.4\mu$m wide bands. Contours are evenly spaced with the lowest contour and spacing at 0.01 MJy/SR. a) Continuum centered at $25.0\mu$m. b) Continuum centered at $32.0\mu$m. } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{PAH11.eps} \caption{ \label{fig:PAH11} Dust emission feature at 11.3$\mu$m. No continuum has been subtracted as the emission feature dominates the spectrum by a factor of 3--8. The minimum contour and spacing is approximately 10 MJy/SR in the peak of the line. The black contours show the oval dust shell previously described by \citet{quillen_supershell}. } \end{figure} \clearpage \begin{figure} \includegraphics[angle=0,width=3.5in]{SIIIc.eps} \includegraphics[angle=0,width=3.5in]{SiIIc.eps} \caption{ \label{fig:Ssi} Continuum subtracted line emission images in the LH spectral cube, showing emission in the folded star forming disk. Contours are evenly spaced. The lowest contours and spacing for the [SIII]($33.481\mu$m) and [SiII]($34.815$m) images are 1.0 and 1.5 $\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$, respectively. a) For [SIII]($33.481\mu$m). b) For [SiII]($34.815\mu$m). } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{FeIIc.eps} \includegraphics[angle=0,width=3.5in]{FeIIIc.eps} \caption{ \label{fig:Fe} Continuum subtracted line emission images in the LH spectral cube, showing emission in the folded star forming disk. Contours are evenly spaced. The lowest contours and spacings for [FeII](25.988$\mu$m) and [FeIII](22.925$\mu$m) images with lowest contour are 0.05 and 0.025 $\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$, respectively. a) For [FeII]($25.988\mu$m). b) for [FeIII](22.925$\mu$m). } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{H2S0c.eps} \caption{ \label{fig:H2S0} Continuum subtracted line emission images in the LH spectral cube, showing emission in the folded star forming disk. for the H$_2$S(0)($28.22\mu$m) line with lowest contour at and spacing at 0.01 $\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$. } \end{figure} \clearpage \begin{figure} \includegraphics[angle=0,width=3.5in]{OIVc.eps} \includegraphics[angle=0,width=3.5in]{NeVc.eps} \caption{ \label{fig:OIVNEV} Continuum subtracted line emission images in the LH spectral cube, showing emission near the jet axis. Contours are evenly spaced. The lowest contours and spacing for the [OIV]($25.890\mu$m) and [NeV]($24.318\mu$m) images are 0.5 and 0.1 $\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$, respectively. a) For [OIV]($25.890\mu$m). b) For [NeV]($24.318\mu$m). } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{radioo4.eps} \caption{ Radio emission at 5Ghz shown as grayscale with [OIV]($25.9\mu$m) contours. The radio map is a 6\arcsec resolution map of the inner radio lobes by \citet{hardcastle06}. The [OIV]$25.9\mu$ and [NeV]$24.3\mu$m line emission are oriented approximately but not exactly along the jet axis. \label{radioo4} } \end{figure} \clearpage \begin{figure} \includegraphics[angle=0,width=3.5in]{H2S3.eps} \includegraphics[angle=0,width=3.5in]{H2S5.eps} \caption{ \label{fig:H2S3} Continuum subtracted line emission images from the SL spectral cube, showing the H$_2$S(3)(9.665$\mu$m) and H$_2$S(5)(6.909$\mu$m) lines with lowest contours at 1 $\times 10^{-7}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$. The higher S rotational lines from H$_2$ exhibit different morphology than the H$_2$S(0)28$\mu$m line that was seen primarily in the folded star forming disk. Emission in the higher S lines is seen above the disk. Contours are evenly spaced. Lowest contour and spacing is $\times 10^{-7}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$ for both lines. For a) H$_2$S(3)(9.665$\mu$m). For b) H$_2$S(5)(6.909$\mu$m). } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{h2s3_o4.eps} \includegraphics[angle=0,width=3.5in]{h2s3_pah11.eps} \caption{ \label{fig:h2s3_ov} a) Emission in H$_2$S(3)(9.665$\mu$m) shown as gray scale overlayed with [OIV]($25.890\mu$) contours. Contours are evenly spaced with lowest contours at are 0.5 $\times 10^{-8}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$ and a spacing 4 times this. The gray scale range for the H$_2$S(3) image is 0(white) to 7(black) $\times 10^{-7}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$. b) Same as a) except the H$_2$S(3) emission is overlayed with contours of the 11.3$\mu$m PAH dust emission feature (as shown in figure \ref{fig:PAH11}). Contour spacing and lowest level is 20 MJy/SR. } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{PAH7.eps} \includegraphics[angle=0,width=3.5in]{PAH8.eps} \caption{ \label{fig:PAH} Flux at 7.7 and 8.6$\mu$m showing dust emission features. Contours are evenly spaced. No continuum subtraction has been done. The star forming disk is evident as the parallelogram shaped feature corresponding to the folded disk. The dust shell is seen above and below the parallelogram feature. The minimum contour and spacing is approximately 10 MJy/SR in the peak of the line. a) The 7.7$\mu$m dust emission feature. b) The 8.6$\mu$m dust emission feature. } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{NeII.eps} \caption{ \label{fig:NeII} Line emission in [NeII]($12.81\mu$m). No continuum has been subtracted as the [NeII] line dominates the continuum by a factor greater than 2 everywhere. Contours are evenly spaced with lowest contour at 4 $\times 10^{-7}$erg cm$^{-2}$ s$^{-1}$ SR$^{-1}$. } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{PAH7_11.eps} \includegraphics[angle=0,width=3.5in]{NeII_PAH11.eps} \caption{ \label{fig:ratio} a) The dust emission feature at 7.7$\mu$m divided by that at 11.3$\mu$m. The lowest and highest contours are shown at a ratios of 0.5 (in shell) and 1.0 (in parallelogram), with spacing of 0.1. Black is 1.0, white is 0.5. The 7.7 to 11.0$\mu$m dust feature ratio varies with the dust shell having the lowest ratio. b) The [NeII](12.8$\mu$m) line divided by the $11\mu$m dust emission feature. The lowest and highest contours are shown at ratios of 1.9 (in shell) and 3.0 (in parallelogram), with spacing of 0.1. Black is 3.0, white is 1.9. The strength of the NeII](12.8$\mu$m) compared to the 11.3$\mu$m dust feature also varies with the dust shell having the lowest ratios. } \end{figure} \clearpage \begin{figure} \includegraphics[angle=0,width=3.5in]{sfrSL.eps} \includegraphics[angle=0,width=3.5in]{jetSL.eps} \caption{ \label{fig:SLspec} a) Spectrum from the star forming parallelogram. D refers to a dust emission feature. b) Spectrum from the jet region. The pure rotational molecular hydrogen lines S(2)--S(7) are labeled as S2-S7. Molecular hydrogen lines are more prominent in the jet region than in the star forming parallelogram. In the star forming disk the [ArII](6.985$\mu$m) line is brighter than the nearby pure rotational molecular hydrogen S(2) line at $6.909\mu$m whereas the opposite is true in the jet region. There is a change in the ratio of the dust emission features in these two spectra and in the ratio of the [NeII]/11.3 dust emission feature. } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{ansi.eps} \includegraphics[angle=0,width=3.5in]{ap1.eps} \caption{ \label{fig:LHspec} a) Spectra from the star forming warped disk (seen as a parallelogram in continuum) are shown at the 4 peaks in the parallelogram and at one additional point also in the parallelogram. Each spectrum is offset by $+$0.1 MJy/SR from the other. From bottom to top: The northern peak west of the nucleus, the southern peak west of the nucleus, the southern peak east of the nucleus, the northern peak east of the nucleus, a point between the southern and eastern peak and the nucleus. Positions are given in the text. In the star forming disk or parallelogram [NeV]($24.318\mu$m) is not detected and [OIV]($25.890\mu$) is weaker than [FeII]($25.988\mu$m). b) Spectrum from a region near the jet axis south-west of nucleus. Near the jet axis [NeV]($24.318\mu$m) is detected and [OIV]($25.890\mu$) is 3-5 times brighter than [FeII]($25.988\mu$m). The presence of [NeV] implies suggests that the radiation field is hard near the jet axis. } \end{figure} \begin{figure} \includegraphics[angle=0,width=3.5in]{h2plot.eps} \caption{ \label{fig:h2plot} Excitation diagram of the H$_2$ pure rotational lines. Points are based on those measured in the jet region with fluxes listed in Table \ref{tab:hjet}. Solid squares refer to measurements and open triangles to upper limits. The solid line is for a temperature of $T_2=720K$ whereas the dotted one is $T_1 = 260K$. } \end{figure} \newpage | 7 | 10 | 0710.5087 |
0710 | 0710.5278_arXiv.txt | {By re-processing the data of the second season of the OGLE survey for planetary transits and adding new mesurements on the same fields gathered in subsequent years with the OGLE telescope, we have identified 23 new transit candidates, recorded as OGLE-TR-178 to OGLE-TR-200. We studied the nature of these objects with the FLAMES/UVES multi-fiber spectrograph on the VLT. One of the candidates, OGLE-TR-182, was confirmed as a transiting gas giant planet on a 4-day orbit. We characterised it with further observations using the FORS1 camera and UVES spectrograph on the VLT. OGLE-TR-182b is a typical ``hot Jupiter'' with an orbital period of 3.98 days, a mass of 1.01 $\pm 0.15$ M$_{Jup}$ and a radius of 1.13 $^{+0.24}_{-0.08}$ R$_{Jup}$. Confirming this transiting planet required a large investment in telescope time with the best instruments available, and we comment on the difficulty of the confirmation process for transiting planets in the OGLE survey. We delienate the zone were confirmation is difficult or impossible, and discuss the implications for the Corot space mission in its quest for transiting telluric planets. | Transiting extrasolar planets are essential to our understanding of planetary structure, formation and evolution outside the Solar System. The observation of transits and secondary eclipses gives access to such quantities as a planet's true mass, radius, density, surface temperature and atmospheric spectrum. The first transiting exoplanet was identified in 1999 around HD 209458. In the past three years, transiting exoplanet have been found in rapidly increasing number, both by radial velocity planet searches and by photometric surveys\footnote{For an updated list see obswww.unige.ch/$\sim$pont/TRANSITS.htm}. The OGLE search for transiting planets and low-mass stellar companions \citep{ud02a} has been the first photometric transit survey to yield results. The first three seasons of photometric observations have revealed 137 transit candidates \citep{ud02a,ud02c,ud02b,ud03}, among which 5 planets were found \citep{ko03,bo04,po04a,kon04,bo05,ko05}, as well as two planet-sized low-mass stars \citep{po05b,po06b}. Three further seasons of the OGLE transit survey have now been completed and await publication (Minniti et al., Udalski et al., in prep.). The spectroscopic follow-up of most of the 137 first OGLE transit candidates, presented in \citet{bo05} and \citet{po05a}, has shown that the vast majority of the transit candidates were eclipsing binaries. A rate of one transiting planet for 10-20 eclipsing binaries is typical. A higher rate of planets can be found among candidates near the detection threshold. Two of the five planets from the OGLE survey, OGLE-TR-56 and OGLE-TR-132, were identified as candidates only after the application of a more sensitive transit detection algorithm \citep{ko02}. However, lowering the detection threshold comes at the price of including some false positives of the detection procedure. The objective of the present study is to explore the regime near the detection threshold, the zone where the ratio of planets to eclipsing binaries will be more favourable than for deeper transit signals, but where the reality of the signal itself is not beyond doubt. The exploration of this zone is relevant not only to identify new transiting planets in the OGLE survey, but also because other transit surveys will face similar issues, notably the {\it CoRoT} and {\it Kepler} space-based transit searches. | \subsection{OGLE-TR-182b as a transiting planet} The companion of OGLE-TR-182 has parameters typical of the planets detected by photometric transit surveys in all respects: it orbits a high-metallicity dwarf star, it has a mass comparable to that of Jupiter and a slightly larger size. Its period is close to an integer number of days, reflecting the strong selection bias due to the window function \citep[see e.g.][about OGLE-TR-111$b$, another $P\sim 4$ days transiting planet]{po04a}. At present, the constraint on the planetary radius is not sufficient to determine whether its radius corresponds to model expectations or whether is belongs to the set of anomalously large transiting hot Jupiters. Its position in the mass-period diagram is also similar to other known transiting planets, and reinforces the link between mass and period for close-in gas giants first pointed out by \citet{ma05}. \subsection{The ``Twilight Zone'' of transit surveys} \begin{figure*}[ht] \resizebox{16cm}{!}{\includegraphics[angle=-90]{figMR.ps}} \caption{''Observational'' mass-radius plot for known transiting exoplanets. The horizontal axis is the planet mass, scaled to $M^{2/3}P^{1/3}$, proportional to the radial-velocity semi-amplitude $K$. The vertical axis is the planet radius scaled to the host star radius, related to the depth of the photometric transit signal. "TR-" labels refer to OGLE candidates (closed symbols). Open symbols mark the position of other known transiting planets. The star symbol marks the best-fit location of another unsolved planet candidate from the third OGLE season. The gray bands show the near-threshold zones for the photometric detection (horizontal) and the spectroscopic confirmation (vertical) in the case of the OGLE survey. The dashed area is the ``twilight zone'' defined in the text, where confirmation is problematic.} \label{mr} \end{figure*} The confirmation follow-up process for OGLE-TR-182 necessitated more than ten hours of FLAMES/VLT time for the radial velocity orbit, plus a comparable amount of FORS/VLT time for the transit lightcurve. In addition, several unsuccessful attempts were made to recover the transit timing in 2007 with the OGLE telescope, and 7 hours of UVES/VLT were devoted to measuring the spectroscopic parameters of the primary. This represents a very large amount of observational resources, and can be considered near the upper limit of what can be reasonably invested to identify a transiting planet. Therefore, OGLE-TR-182 is a useful object to quantify the zone were neither the photometric signal nor the radial velocity signal are clear beyond doubt, the ``twilight zone'' of planetary transit candidate confirmation. In these cases, confirming the nature of the system is very difficult and time consuming. When the photometric signal is a possible false positive, a clear radial velocity orbit at the same period is an essential confirmation | as was for instance the case for OGLE-TR-132 \citep{bo04}. On the other hand, when the radial velocity signal is marginal, a clear transit signal allows the phase and period of the orbit to be determined with confidence, reducing the radial-velocity orbit fit to a two-parameter problem ($V_0$, the systemic velocity, and $K$, the orbital semi-amplitude) | as was the case for OGLE-TR-10 \citep{bo05,ko05}. However, when both the photometric and spectroscopic signals are marginal, many more observations are necessary until reasonable certainty can be achieved about the presence of a planetary companion. The uncertainties on the lightcurve make it difficult to phase the radial velocity data. The high radial velocity uncertainties hinder the identification of an orbital motion with the correct period, and the elimination of eclipsing binary blend scenarios The OGLE survey is the first to explore this ``twilight zone'' in real conditions, since other ground-based surveys target brighter stars, for which very precise radial velocities can be obtained, so that the significance of the radial velocity signal can be established relatively easily \citep[e.g.][]{ca07}. Based on the cases of OGLE-TR-10, OGLE-TR-132 and OGLE-TR-182, and on the discussion in \citet{po06}, we define the limits of the follow-up twilight zone as follows: \noindent -- photometric transit detection with $8< S_r < 12$, where $S_r$ is the transit significance in the presence of red noise \citep[see definition in][]{po06} \noindent -- radial velocity orbital semi-amplitude 1-2 times the radial velocity uncertainties for 1-hour exposures with the facilities available: $\sigma_{vr} < K < 2\cdot \sigma_{vr} $ These limits can be translated, for a circular orbit and a central transit, into limits on the radius and mass of the planet. Figure~\ref{mr} shows the ``observational'' mass-radius diagram for the known transiting exoplanets. The horizontal axis is the planet mass divided by $M_*^{2/3}P^{1/3} $, to make it proportional to the observed radial velocity semi-amplitude ("$K$"). The vertical axis is the planet radius divided by the radius of the star, to make it proportional to the squareroot of the transit depth (at wavelengths where stellar limb darkening can be neglected). The units are such that a Jupiter-sized planet transiting a solar-sized star on a 4-day orbit will be placed at (1;1). Objects at similar positions in this plot will present similar challenges for confirmation. The \mbox{ ``TR-''} labels refer to OGLE candidates, the other unmarked points are transiting planets from other surveys. The gray horizontal band is the zone where the transit detection is near the detection threshold, the vertical band is the zone where the radial velocity orbital signal is near the threshold. The intersection of the two, delimited by the dashed lines, represents the ``twilight zone'' for the OGLE survey. We use a red noise level of $\sigma_r = 3$ mmag \citep{po06}, a radial velocity uncertainty of 60-70 m/s (photon noise plus systematics), and assume that 5-10 transits are observed by the photometry. We find $ 0.85 < R_{pl}/R_* < 1.20$ and $ 0.45 < M_{pl}/{M_*}^{-2/3} P^{-1/3} < 1.05$ [$M_J {M_\odot}^{-2/3} ({4\ {\rm days}})^{-1/3}$] for the zone boundaries. Candidates in this zone will be very difficult or impossible to confirm. On the left of the zone, radial velocity confirmation is out of reach, and below, the transit signal is below the photometric detection threshold. The twilight zone for the OGLE transit survey encompasses the region in the mass-radius diagram corresponding to a normal hot Jupiter around a solar-type star. Hence the difficulty of the OGLE survey to detect transiting gas giants unless they are exceptionally heavy with a very short period like OGLE-TR-56, OGLE-TR-113 and OGLE-TR-132, or have an exceptionally high radius ratio like OGLE-TR-111. The planetary transit system OGLE-TR-10 is also located within the zone, and indeed confirmation of its planetary nature required large investments in follow-up means both with the VLT \citep{bo05} and Keck \citep{ko05} 8-10 meter telescopes. As a further illustration of the extent of the zone, the plot shows the position of another candidate from the OGLE survey, yet unsolved despite extensive measurements with FLAMES and FORS. The radial velocity data is compatible with a planetary orbit, but many more observations would be required to confirm it securely. The best-fit planetary solution for this object is plotted on Fig.~\ref{mr}. In the wider context of transit searches in general, it is interesting to find where the ``twilight zone'' is located for different surveys, especially the space-based transit searches {\it CoRoT} and {\it Kepler}. For wide-field, small-camera surveys like HAT, WASP, XO and TrES, the twilight zone is not an important issue. Because the candidates are brighter, standard planet-search spectrographs can be used for the radial velocity follow-up, and the zone moves to the left of the mass-radius diagram, in a region were no planets are expected (low-mass, Jupiter-sized planets). In other words, if a planet is large enough to be detected by these surveys, it produces a radial velocity signal that is easily picked up by Doppler spectrographs. Deeper surveys like SWEEPS \citep{sa06} and planet searches in star clusters also have no twilight zone problem, for the opposite reason: their candidates are too faint to be confirmed in radial velocity for Jupiter-like planetary masses. In the case of the {\it CoRoT} space transit search, the zone will be located in a key position. In planet radius, it is expected to cover the 2-4 $R_\oplus$ range \citep[see][]{mo05} for the brightest targets. In planet mass, using the HARPS spectrograph, it will be in the 5-20 $M_\oplus $ range, for short periods. This is a zone were planets are thought to be numerous, the domain of the ``hot Neptunes'' and ``super-Earths''. Indeed, the detection of this type of planets constitutes the main objective of the {\it CoRoT} planetary transit search. From our experience with the OGLE follow-up, we therefore conclude that the {\it CoRoT} mission will face similarly difficult cases in the confirmation process of transiting planets. The telescope time necessary for the follow-up of these candidates should be adequately evaluated. The OGLE follow-up process can provide some useful guidelines. | 7 | 10 | 0710.5278 |
0710 | 0710.1201_arXiv.txt | We show that: (i) the long-term X-ray outburst light curve of the transient AXP XTE J1810--197 can be accounted for by a fallback disk that is evolving towards quiescence through a disk instability after having been heated by a soft gamma-ray burst, (ii) the spin-frequency evolution of this source in the same period can also be explained by the disk torque acting on the magnetosphere of the neutron star, (iii) most significantly, recently observed pulsed-radio emission from this source coincides with the epoch of minimum X-ray luminosity. This is natural in terms of a fallback disk model, as the accretion power becomes so low that it is not sufficient to suppress the beamed radio emission from XTE J1810--197. | Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are systems of young neutron stars identified mainly by X-ray luminosities ($10^{34}-10^{36}$ erg s$^{-1}$) much higher than their rotational powers indicated by their observed spin period and period derivatives (Woods \& Thompson 2005, Kaspi 2006). Brief soft gamma-ray bursts (\la 1 s) with super-Eddington luminosities observed from SGRs and some AXPs led to the identification of these sources as magnetars (Duncan \& Thompson 1992). While the burst energetics are very likely to have magnetic origin, bursts themselves, being local events, do not require that the extended dipole component of the magnetic field has magnetar strength. Strong magnetic fields possibly in higher multipoles which decrease rapidly beyond the neutron star surface is likely to be the source of these bursts. Fallback disk models (Chatterjee et al. 2000; Alpar 2001) present a self-consistent explanation for the period clustering, soft X-ray, optical and infrared (IR) light curves of AXPs and SGRs in their persistent and enhanced phases (Ek\c si \& Alpar 2003; Ertan \& Alpar 2003; Ertan \& Cheng 2004; Ertan et al. 2006; Ertan \& \c{C}al{\i}\c{s}kan 2006, Ertan et al. 2007). These explanations entail active, viscously dissipating disks around these systems. Recent observations of AXP 4U 0142+61 in the mid-IR bands with {\it Spitzer} Observatory (Wang et al. 2006) provided the first instance of a disk around an AXP. Wang et al. (2006) attributed the optical and mid-IR data of this source to magnetospheric emission as proposed by magnetar models, and showed that the mid-IR data can be fitted with an irradiated passive disk. Detailed analysis of the overall data set including earlier detections in optical and near-IR bands strongly indicate the presence of an active disk around this source (Ertan et al. 2007). The difference between a passive and an active disk is that there is no viscous dissipation and therefore no mass inflow along a passive disk, while in an active disk, viscous dissipation works and leads to mass flow towards the inner disk and to the emission of X-rays through accretion onto the neutron star. Both models expect a similar luminosity in the mid-IR bands which is emitted from the outer disk regions at which irradiation dominates dissipation. Nevertheless, estimates of these models differ in near IR and especially optical luminosities which, in the active disk model, are expected from the inner disk regions where the intrinsic dissipation becomes important. In a passive-disk model, the inner disk must be cut at a sufficiently large radius having a critical temperature below which temperatures are assumed to be not sufficient to sustain magnetorotational instability (MRI)(Balbus $\&$ Hawley 1991), as proposed by Wang et al.(2006) for the disk around AXP 4U 0142+61. On the other hand, the minimum temperature required for MRI which leads to viscous dissipation is uncertain. Recent work by Inutsuka $\&$ Sano (2005) shows that MRI can be sustained in most regions of a protoplanetary disk even at temperatures as low as 300 K. Recently, Ertan et al. 2007 showed, by means of numerical fits to data, that the overall optical and IR dataset together with the X-ray luminosity of AXP 4U 0142+61 can be fitted with a single active, irradiated-disk model. The upper limit for the dipole field strength estimated from these model fits is about $10^{13}$ G on the neutron star's surface (see Ertan et al. 2007 for details). The X-ray enhancement of SGR 1900+14 following its giant flare can be explained by the evolution of a disk after it has been pushed back and piled up at larger radii by the burst (Ertan \& Alpar 2003). This pushed-back disk model can also reproduce the long lasting, contemporaneous X-ray and IR enhancement light curve of AXP 1E 2259+586 for which the initial burst assumed by the model is likely to have remained below the detection limits (Ertan et al. 2006). In these models, reasonable fits to data are obtained by employing a single $\a$ viscosity (Shakura \& Sunyaev 1973) without invoking a disk instability. It might be concluded that in the observed X-ray luminosity, $\Lx$, range of these and other persistent AXPs and SGRs ($10^{34} - 10^{36}$ erg s$^{-1})$, fallback disks do not experience global disk instabilities. Are these persistent sources already in their quiescent states, or do they have a lower viscosity state which will finally take them into quiescence with a sharp decrease in $\Lx$ after the mass-flow rate, $\Mdot$, at the inner disk has decreased to below some critical level? In this paper, we suggest that the known persistent AXPs and SGRs are not in their quiescent states, and that the transient AXP XTE J1810$-$197 is a good example for an AXP that is now, following an outburst, making a transition back into quiescence by means of a viscous disk instability. The observations of pulsed radio emission from this source fall in the epoch following the last X-ray observation by Gotthelf and Halpern (2007) (Fig 1). Our model X-ray light curve predicts that in this epoch the $\Lx$ remains even below the pre-outburst quiescent level. We propose that the pulsed-radio emission is a natural outcome of accretion rate decreasing below a critical value such that it is no longer sufficient to hinder pulsed-radio emission. In this picture, all the other persistent AXPs and SGRs which are expected to be in the hot viscosity state at present would finally evolve into quiescence as their luminosity and thus the irradiation strength decrease below a critical level. The persistent and quiescent states would then correspond to two different viscosity states, and the AXPs and SGRs could make transitions between these cold and hot states if the system undergoes a global instability triggered, for instance, by an energetic soft gamma-ray burst. Furthermore, persistent AXPs and SGRs are not expected to show pulsed radio emission, because their accretion rates are much higher than the critical values below which pulsed-radio emission is allowed. Three of the AXPs (of which one is a candidate AXP) and one SGR were observed to show transient behavior in their X-ray luminosity (Kaspi 2006, Mereghetti et al. 2006, Israel et al. 2007, Gotthelf and Vasisht 1998, Torii et al. 1998). Of these transient sources, XTE J1810$-$197 is a confirmed AXP with measured spin period and period derivative (Ibrahim et al. 2004). This source was discovered during the early decay phase of its X-ray outburst (Ibrahim et al. 2004). The $\Lx$ of XTE J1810$-$197 has still been decaying three years after its first detection, and now is about two orders of magnitude less than the observed maximum $\Lx$. This source was not observed around the onset of the X-ray outburst. Earlier archival data show that the $L_x$ prior to the outburst is also about two orders of magnitude fainter than the observed maximum luminosity (Gotthelf \& Halpern 2007). During the first year of the decay phase of XTE J1810$-$197, $\Lx$ remained in the range characteristic of persistent AXPs. In this phase, the inner disk seems to be in a hot viscosity state, with temperatures above the critical level, likely to be the same as that operating in the persistent AXPs. In the present work, we show that for this first year, the X-ray light curve of XTE J1810--197 can be simulated by a pure viscous evolution of a disk pushed back by a burst, but subsequently as $\Lx$ decreases below a critical level around $10^{34}$ erg s$^{-1}$ the light curve deviates from the characteristic viscous relaxation model curve. This transition can be explained by a viscous disk instability leading the system into quiescence. The disk instability model for hydrogen disks is summarized and the possibilities for the instabilities of the fallback disks are discussed in \S\ 2. In \S\ 3, we show that the overall decay curve of the transient AXP XTE J1810--197 can be reproduced by an irradiated active disk that undergoes a transition into quiescence below some critical X-ray luminosity that is predictable from the model fits. In \S\ 4 and \S\ 5, we discuss the X-ray light curve and its connection to pulsed radio emission respectively. In \S\ 6, we investigate the spin-frequency evolution of XTE J1810--197 during the post-burst fading of the X-ray flux. Our results are summarized in \S\ 7. | We have firstly shown that the X-ray outburst light curve of the transient AXP XTE J1810$-$197 lasting for about three years can be reproduced by the evolution of an irradiated disk that undergoes a transition into quiescence by means of a viscous disk instability as the X-ray luminosity decreases below a critical limit less than around $10^{34}$ erg s$^{-1}$ which is near the minimum of the observed luminosity range of AXPs and SGRs. We have assumed that the inner disk is initially pushed back by a missed soft gamma-ray burst that took a large portion of the inner disk into a hot viscosity state which is likely the state prevailing in the disks of persistent AXPs and SGRs. The mass-flow rate, $\Mdot$, provided by the extended disk of XTE J1810$-$197 with surface densities likely to be lower than those of persistent sources is not sufficient for the inner disk to be kept in the hot state for longer periods by fuelling the X-ray irradiation through accretion. After the X-ray luminosity has decreased below a critical level, the system evolves into quiescence, by means of a viscous disk instability. Observations indicate that the source was also in the quiescent state prior to the X-ray outburst. Unlike XTE J1810$-$197, persistent sources AXP 1E 2259+586 and SGR 1900+14 were already in the hot state with $\Lx ~\ga ~10^{34}$ erg s$^{-1}$ before their enhancements and their X-ray luminosities remained above the critical level during the outburst decay as well. The X-ray luminosity of these and other persistent AXPs and SGRs are also expected to decrease slowly as the disk surface densities decrease gradually through expansion and accretion in the persistent phase. They will also finally evolve towards the quiescent regime by means of a disk instability as their accretion rate falls below critical values. Quiescent sources, on the other hand, could make an upward transition to the hot state when a viscous instability is triggered by an energetic gamma-ray burst. Secondly, we have shown that the magnetospheric torque applied by the disk on the star is in agreement with the spin-frequency evolution of this transient AXP (see Fig. 2). Along the X-ray outburst decay, XTE J1810$-$197 is in the accretion regime with spin-down. In this phase, the system is slowly receding from the rotational equilibrium between the disk and the magnetosphere, while most of the mass flowing towards the disk is being accreted onto the neutron star. Our torque model is almost independent of the mass-inflow rate $\Mdot$ to the inner disk in this period. We have obtained the model curve for spin evolution directly from the evolution of the mass-inflow rate corresponding to the model X-ray light curve presented in Fig. 1 and by adjusting the field strength accordingly. The model curve given in Fig.2 is obtained with a dipolar magnetic field of $B_{\ast }\simeq 2\times 10^{12}$ G. Below some critical value of $\Mdot$, most of the inflowing matter is propelled out of the system, and the functional form of the torque is expected to change. In this regime the torque would probably be less efficient than the model torque we employ here. The dependence of the torque on $\Mdot$ also changes as the system approaches rotational equilibrium (see \S\ 6), and the spin-up torque due to accretion is no longer negligible in comparison with the spin-down torque. Finally, we have proposed that the pulsed-radio emission from XTE J1810--197 could be accounted for by the decreasing accretion rate below a critical level which switches on the beamed radio emission from the polar caps. Our model predicts a dip-like structure at the end of the decay phase such that the X-ray luminosity remains even below the pre-outburst quiescent level for about one year. This period matches the epoch of observations of pulsed-radio emission from this transient AXP (Camilo et al. 2007). This interpretation of the pulsed radio emission together with the model predictions for the X-ray light curve could be tested by future radio and X-ray observations of XTE J1810--197 and other transient AXPs and SGRs. Transient AXPs can show up as radio pulsars when their disks make a transition to the cold state and the X-ray luminosity drops sufficiently such that their accretion power can no longer suppress the beamed radio emission. A detailed analysis on the torque evolution and the beamed radio emission from AXPs and SGRs will be the subject of our future work. | 7 | 10 | 0710.1201 |
0710 | 0710.3518_arXiv.txt | {The results from H.E.S.S. observations towards Westerlund~2 are presented. The detection of very-high-energy gamma-ray emission towards the young stellar cluster Westerlund 2 in the HII complex RCW49 by H.E.S.S. provides ample evidence that particle acceleration to extreme energies is associated with this region. A variety of possible emission scenarios is mentioned, ranging from high-energy gamma-ray production in the colliding wind zone of the massive Wolf-Rayet binary WR~20a, collective wind scenarios, diffusive shock acceleration at the boundaries of wind-blown bubbles in the stellar cluster, and outbreak phenomena from hot stellar winds into the interstellar medium. These scenarios are briefly compared to the characteristics of the associated new VHE gamma-ray source HESS~J1023--575, and conclusions on the validity of the respective emission scenarios for high-energy gamma-ray production in the Westerlund~2 system are drawn.} \begin{document} | 7 | 10 | 0710.3518 |
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0710 | 0710.3803_arXiv.txt | We present interferometric observations of $^{12}$CO and $^{13}$CO $J$=2$-$1 emission from the butterfly-shaped, young planetary nebula NGC 6302. The high angular resolution and high sensitivity achieved in our observations allow us to resolve the nebula into two distinct kinematic components: (1) a massive expanding torus seen almost edge-on and oriented in the North-South direction, roughly perpendicular to the optical nebula axis. The torus exhibits very complex and fragmentated structure; (2) high velocity molecular knots moving at high velocity, higher than 20 \kms, and located in the optical bipolar lobes. These knots show a linear position-velocity gradient (Hubble-like flow), which is characteristic of fast molecular outflow in young planetary nebulae. From the low but variable $^{12}$CO/$^{13}$CO $J$=2$-$1 line intensity ratio we conclude that the $^{12}$CO $J$=2$-$1 emission is optically thick over much of the nebula. Using the optically thinner line $^{13}$CO $J$=2$-$1 we estimate a total molecular gas mass of $\sim$ 0.1 M$_\odot$, comparable to the ionized gas mass; the total gas mass of the NGC 6302 nebula, including the massive ionized gas from photon dominated region, is found to be $\sim$ 0.5 M$_\odot$. From radiative transfer modelling we infer that the torus is seen at inclination angle of 75$^\circ$ with respect to the plane of the sky and expanding at velocity of 15 \kms. Comparison with recent observations of molecular gas in NGC 6302 is also discussed. | Low and intermediate-mass stars (1 M$_\odot$ $\le$ M$_*$ $\le$ 8 M$_\odot$) evolve through the Asymptotic Giant Branch (AGB) phase, which is characterized by copious mass loss in the form of dusty slow wind, before emerging as planetary nebulae (PNe). The mass-loss process is commonly assumed to be isotropic, resulting in the formation of a spherically symmetric envelope around the central AGB star. However, a significant fraction of the circumstellar envelope around post-AGB stars and planetary nebulae possess bipolar and even multipolar morphology. The mechanisms responsible for such departure from spherical symmetry are still unknown (e.g.\ Balick \& Frank 2002). The interaction between collimated fast outflows and the surrounding envelope or the influence of a binary companion are often cited as possible shaping mechanisms. Large expanding or rotating disks/tori has been frequently inferred to be present in bipolar nebulae such as the Egg Nebula (Sahai et al. 1998) and the Red Rectangle (Bujarrabal et al. 2005). These disks or tori might confine and channel the wind from the central star into bipolar directions. If the disks/tori are dense enough, the strong interaction (i.e oblique shock) between the wind and the disks/tori could focus and collimate the outflow (Frank et al. 1996). Thus, more detailed studies of the structure and kinematics of such disk/torus could provide better understanding on the formation of the nebulae. NGC\,6302 is a young planetary nebula and belongs to the class of the highest excitation PNe (Pottasch et al.\ 1996). From the presence of numerous emission lines from highly ionized species, Pottasch et al.\ (1996) concluded that the central star is a white dwarf or approaching a white dwarf and its excitation temperature is very high, $\sim$380,000 K. In optical images, NGC\,6302 appears as a butterfly shaped nebula with two huge bipolar lobes separated by a dark equatorial lane (Matsuura et al. 2005), which is presumably the location of a massive disk or torus. An expanding H\2\ region is detected at the center of the nebula in the radio continuum (G\'{o}mez et al. 1989). The presence of a massive molecular envelope is known through the detection of strong CO rotational lines in NGC\,6302 (Huggins et al.\ 1996, Hasegawa \& Kwok 2003). The CO lines have very peculiar and complex shapes, with a double-peaked profile and high velocity wings, extending up to $\sim$40 \kms\ from the systemic velocity. Such line shape suggests that the molecular envelope is likely non-spherical and could contain a fast molecular outflow, similar to fast bipolar outflows observed in some proto-planetary nebulae and young planetary nebulae. Detailed observations with the VLT and HST of Matsuura et al.\ (2005) show the presence of a large warped disk oriented perpendicular to the bipolar lobes seen prominently in the optical images. The disk is massive and has a large extinction. Strong emission from crystalline silicates together with PAH bands have also been detected in NGC\, 6302 (Kemper et al.\ 2002). We note that, interestingly, OH maser emission, which is usually associated with oxygen rich circumstellar envelopes, has also been detected in NGC\,6302 (Payne et al. 1988). Thus, NGC\,6302 is chemically peculiar, containing a mixture of carbon-rich and oxygen-rich material. The distance to NGC\,6302 is uncertain, with estimates ranging from 0.15 to 2.4 kpc. From the expansion proper motion of the central H\2\ region, Gom\'{e}z et al. (1989) estimates a distance of 2.2 $\pm$ 1.1 kpc. However, Matsuura et al. (2005) note that a distance much larger than 1 kpc would lead to very high luminosity and shell mass, which would exceed the highest luminosity of post-AGB stars and PNe predicted by stellar evolution models. with even the highest core mass. More recently, Meaburn et al. (2005) detect directly the proper motion of the prominent optical lobes in NGC\,6302 nebula and infer a distance of 1 kpc. We will adopt a distance of 1 kpc for NGC\, 6302, similar to Matsuura et al. (2005) and Kemper et al. (2002). In this paper we present high angular resolution observations of the $J$=2$-$1 line of $^{12}$CO and its isotope $^{13}$CO, in order to study the spatial distribution and kinematics of the molecular gas in the envelope of NGC\,6302, especially the massive equatorial disk and the gas moving at higher velocities. After the submission of our paper for publication, we learned of a recent paper by Peretto et al. (2007), which also presents maps of molecular gas in NGC\,6302. Although the spatial distribution of CO emission is similar in both papers, our higher sensitivity (by a factor of 2) observations allow us to produce maps for more velocity channels and better understand the spatial kinematics of the circumstellar envelope. We are able, in particular, to image and study the fast molecular outflows, which are barely detected in lower sensivity data of Peretto et al. (2007). Where appropriate, we will compare our results with those obtained by Peretto et al. (2007) | We have imaged at high sensitivity and high angular resolution the envelope around the young planetary nebula NGC\,6302 in $^{12}$CO $J$=2--1 and $^{13}$CO $J$=2--1 lines. Continuum at 1.3mm wavelegth is also imaged and seems to come from the inner H\2\ region. We find that the very complex molecule-rich nebula is well resolved and can be separated into two components: a massive low-velocity torus seen nearly edge-on and high-velocity knots. The image of the dense torus is very accurately coincident with the dark lane that separates the conspicuous two lobes in the optical image of NGC\,6302. The fast knots are located within the optical lobes and show a linear velocity gradient, which is characteristic of fast molecular gas in young planetary nebulae. We find that the $^{12}$CO/$^{13}$CO $J$=2--1 brigthness ratio is low and varies between 2 and 5, decreasing with increasing line intensity. We conclude that the $^{12}$CO $J$=2--1 emission is optically thick over much of the nebula, but that $^{13}$CO $J$=2--1 is optically thin in most velocities and lines of sight. Using our observations of this line, we estimate masses for the different components, yielding a total molecular gas mass of $\sim$ 0.1 M$_\odot$. We discuss the value of the total mass of the NGC\,6302 nebula, including the ionized gas, whose mass is comparable to that of the molecular gas, and the massive PDR, which is probably the dominant component. We find a total mass of $\sim$ 0.5 M$_\odot$; but we recall that the very uncertain contribution from very outer layers could be significant. Using a radiative transfer model we infer that the torus is seen at an inclination angle of 75$^\circ$ with respect to the plane of the sky and expanding at a velocity of 15 \kms. The mass loss rate is found to be very high, $\sim$ 1.5$\times$10$^{-4}$ \ms\ yr$^{-1}$, resulting in a dense (average gas density $\sim$ 2$\times$10$^4$ cm$^{-3}$) and a massive torus. \bigskip We are grateful to SMA staff for carrying out the observations. We thank an anonymous referee for constructive criticisms that helped to improve our paper significantly. V.B.\ acknowledges support from the \emph{Spanish Ministry of Education \& Science}, project numbers AYA2003-7584 and ESP2003-04957. Help from Sebastien Muller is gratefully acknowledged. This research has made use of NASA's Astrophysics Data System Bibliographic Services and the SIMBAD database, operated at CDS, Strasbourg, France. | 7 | 10 | 0710.3803 |
0710 | 0710.4081_arXiv.txt | Accurate alignment of the radio and optical celestial reference frames requires detailed understanding of physical factors that may cause offsets between the positions of the same object measured in different spectral bands. Opacity in compact extragalactic jets (due to synchrotron self-absorption and external free-free absorption) is one of the key physical phenomena producing such an offset, and this effect is well-known in radio astronomy (``core shift''). We have measured the core shifts in a sample of 29 bright compact extragalactic radio sources observed using very long baseline interferometry (VLBI) at 2.3 and 8.6 GHz. We report the results of these measurements and estimate that the average shift between radio and optical positions of distant quasars would be of the order of 0.1--0.2 mas. This shift exceeds positional accuracy of GAIA and SIM. We suggest two possible approaches to carefully investigate and correct for this effect in order to align accurately the radio and optical positions. Both approaches involve determining a Primary Reference Sample of objects to be used for tying the radio and optical reference frames together. | } Extragalactic relativistic jets are formed in the immediate vicinity of the central black holes in galaxies, at distances of the order of 100 gravitational radii, and they become visible in the radio at distances of about 1000 gravitational radii \citep{LZ2007}. This apparent origin of the radio jets is commonly called the ``core''. In radio images of extragalactic jets, the core is located in the region with an optical depth $\tau_s\approx 1$. This causes the absolute position of the core, $r_\mathrm{core}$, to vary with the observing frequency, $\nu$, since the optical depth profile along the jet depends on $\nu$: $r_\mathrm{core} \propto \nu^{-1/k_\mathrm{r}}$ \citep{BlandfordKonigl79}. Variations of the optical depth along the jet can result from synchrotron self-absorption \citep{Koenigl81}, pressure and density gradients in the jet and free-free absorption in the ambient medium most likely associated with the broad-line region (BLR) \citep{L98}. The core shift is expected to introduce systematic offsets between the radio and optical positions of reference sources, affecting strongly the accuracy of the radio-optical matching of the astrometric catalogues. The magnitude of the core shift can exceed the inflated errors of the radio and optical positional measurements by a large factor. This makes it necessary to perform systematic studies of the core shift in the astrometric samples in order to understand and remove the contribution of the core shift to the errors of the radio-optical position alignment. Measurements of the core shift have been done so far only in a small number of objects \citep[e.g.,][]{MES94,Lara_etal94,PR97,L96,L98,RL2001,Kadler_etal04,SokolovCawthorne2007}. In this paper, we present results for 29 compact extragalactic radio sources used in VLBI astrometric studies and discuss the core shift effect on the task of the radio-optical reference frame alignment. | } Measurements of the frequency-dependent shift of the parsec-scale jet cores in AGN are reported for 29 bright extragalactic radio sources. It is shown that the shift can be as high as 1.4~mas between 2.3 and 8.6~GHz. We have shown that core shifts are likely to pose problems for connecting radio and optical reference frames. We have estimated from theory an average shift between the radio (4~cm) and optical (6000~\AA) bands to be of an order of 0.1~mas for a complete sample of radio selected AGN. The estimated radio-optical core shift exceeds the positional accuracy of GAIA and SIM. It implies that the core shift effect should be carefully investigated, and corrected for, in order to align accurately the radio and optical positions. Based on our investigation, we suggest two possible approaches, both involving determining a Primary Reference Sample of objects to be used for tying the radio and optical reference frames together. 1) In the first approach, multi-frequency VLBI measurements can be used for calculating the projected optical positions, assuming that the radio and optical emission regions are both dominated by a spatially compact component marginally resolved with VLBI and SIM and point-like for GAIA. The discrepancies between the measured optical and radio positions can then be corrected for the predicted shifts, and the subsequent alignment of the radio and optical reference frames can be done using standard procedures. 2) A more conservative approach may also be applied, by employing the VLBI observations to identify and including in the Primary Reference Sample only those quasars in which no significant core shift has been detected in multi-epoch experiments. Either of the two approaches should lead to substantial improvements of the accuracy of the radio-optical position alignment. | 7 | 10 | 0710.4081 |
0710 | 0710.3160_arXiv.txt | The first stars form in dark matter halos of masses $\sim$$10^6 \Ms$ as suggested by an increasing number of numerical simulations. Radiation feedback from these stars expels most of the gas from their shallow potential well of their surrounding dark matter halos. We use cosmological adaptive mesh refinement simulations that include self-consistent Population III star formation and feedback to examine the properties of assembling early dwarf galaxies. Accurate radiative transport is modelled with adaptive ray tracing. We include supernova explosions and follow the metal enrichment of the intergalactic medium. The calculations focus on the formation of several dwarf galaxies and their progenitors. In these halos, baryon fractions in 10$^8$ \Ms~halos decrease by a factor of 2 with stellar feedback and by a factor of 3 with supernova explosions. We find that radiation feedback and supernova explosions increase gaseous spin parameters up to a factor of 4 and vary with time. Stellar feedback, supernova explosions, and \hh~cooling create a complex, multi-phase interstellar medium whose densities and temperatures can span up to 6 orders of magnitude at a given radius. The pair-instability supernovae of Population III stars alone enrich the halos with virial temperatures of 10$^4$ K to approximately 10$^{-3}$ of solar metallicity. We find that 40\% of the heavy elements resides in the intergalactic medium (IGM) at the end of our calculations. The highest metallicity gas exists in supernova remnants and very dilute regions of the IGM. | The majority of galaxies in the universe are low-luminosity, have masses of $\sim$$10^8$ solar masses, and are known as dwarf galaxies \citep{Schechter76, Ellis97, Mateo98}. Galaxies form hierarchically through numerous mergers of smaller counterparts \citep{Peebles68, White78}, whose properties will inevitably influence their parent galaxy. Dwarf galaxies are the smallest galactic building blocks, and this leads to the question on even smaller scales: how were dwarf galaxies influenced by their progenitors? A subset of dwarf galaxies, dwarf spheroidals (dSph), have the highest mass-to-light ratios \citep{deBlok97, Mateo98} and contain a population of metal-poor stars that are similar to Galactic halo stars \citep{Tolstoy04, Helmi06}. There is a metallicity floor exists of $10^{-3}$ and $10^{-4}$ of solar metallicity in dSph and halo stars, respectively \citep{Beers05, Helmi06}. Stellar metallicities increase with time as previous stars continually enrich the interstellar medium (ISM). Hence the lowest metallicity stars are some of the oldest stars in the system and can shed light on the initial formation of dwarf galaxies. This metallicity floor also suggests that metal enrichment was widespread in dark matter halos before low-mass stars could have formed \citep[e.g.][]{Ricotti02b}. Supernovae (SNe) from metal-free (Pop III) stars generate the first metals in the universe and may supply the necessary metallicity to form the most metal-poor stars observed \citep{Ferrara98, Madau01a, Norman04}. Dwarf galaxy formation can be further constrained with observations that probe reionization and semi-analytic models. Observations of luminous quasars powered by supermassive black holes (SMBH) of mass $\sim$$10^9 \Ms$ \citep{Becker01, Fan02, Fan06} and low-luminosity galaxies \citep{Hu02, Iye06, Kashikawa06, Bouwens06, Stark07} at and above redshift 6 indicate that active star and BH formation began long before this epoch. Semi-analytic models have argued that cosmological reionization was largely caused by low-luminosity dwarf galaxies \citep{Haiman97, Cen03, Somerville03, Wise05, Haiman06}. Some of the most relevant parameters in these models control star formation rates, ionizing photon escape fractions, metal enrichment, and the minimum mass of a star forming halos. They are usually constrained using (i) the cosmic microwave background (CMB) polarization observation from Wilkinson Microwave Anisotropy Probe (WMAP) that measures the optical depth of electron scattering to the CMB \citep{Page06}, (ii) Gunn-Peterson troughs in $z \sim 6$ quasars, and (iii) numerical simulations that examine negative and positive feedback of radiation backgrounds \citep{Machacek01, Machacek03, Yoshida03, Mesigner06}. Radiation hydrodynamical \textit{ab initio} simulations of the first stars \citep{Yoshida06a, Abel07} and galaxies can further constrain the parameters used in semi-analytic models by analyzing the impact of stellar feedback on star formation rates and the propagation of \ion{H}{2} regions in the early universe. Moreover, these simulations contain a wealth of information pertaining to the properties of Pop III star forming halos and early dwarf galaxies that can increase our understanding of the first stages of galaxy formation. First we need to consider Pop III stars, which form in the progenitor halos of the first galaxies, to capture the initial properties of dwarf galaxies. Cosmological numerical studies have shown that massive (30--300 \Ms) Pop III stars form in dark matter halos with masses $\sim$$10^6 \Ms$ \citep{Abel02, Bromm02, Yoshida06b, Gao07, OShea07}. Recently, \citeauthor{Yoshida06b} and \citet{Turk08} followed the gaseous collapse of a molecular cloud that will host a Pop III star to cosmologically high number densities of $10^{16}$ and $10^{21}$ \cubecm, respectively. The former group thoroughly analyzed the gas dynamics, cooling, and stability of this free-fall collapse. The latter group observed a protostellar core forming with 10 Jupiter masses that is bounded by a highly asymmetric protostellar shock. Both groups found no fragmentation in the fully molecular core that collapses into a single, massive $\sim$100\Ms~star. Furthermore, \citet{Omukai03} determined that accretion may halt at the same mass scale, using protostellar models even for different mass accretion histories. Pop III stars with stellar masses roughly between 140 and 260 \Ms~end their life in a pair-instability SN (PISN) that releases $10^{51} - 10^{53}$ ergs of energy and tens of solar masses of heavy elements into the ambient medium \citep{Barkat67, Bond84, Heger02}. These explosions are an order of magnitude larger than typical Type II SNe in both quantities \citep{Woosley86}, such explosions energies are larger than the binding energies in their low-mass hosts, e.g., $2.8 \times 10^{50}$ ergs for a $10^6 \Ms$ halo at redshift 20. Thus gas structures in the host halo are totally disrupted and expelled, effectively enriching the surrounding intergalactic medium (IGM) with the SN ejecta \citep{Bromm03, Kitayama05, Greif07}. The combination of the shallow potential well and large explosion energy suggests that these events are good candidates for enriching the first galaxies and IGM. Outside of the pair-instability mass range, Pop III stars die by directly collapsing into a BH \citep{Heger03}, possibly providing the seeds of high-redshift quasars in galaxies that are associated with the rarest density fluctuations \citep[e.g.][]{Madau01b, Volonteri05, Trenti07}. One-dimensional calculations \citep{Whalen04, Kitayama04, Kitayama05} and recent three-dimensional radiation hydrodynamical simulations \citep{Yoshida06a, Abel07} have investigated how the Pop III stellar feedback affects its host halo and nearby cosmic structure. In addition to SNe, \ion{H}{2} regions surrounding Pop III stars, which have luminosities $\sim$$10^6 L_\odot$ \citep{Schaerer02}, alone can dynamically affect gas at distances up to a few proper kpc. Ionization fronts and \ion{H}{2} regions \citep[see][for a review]{Yorke86} have been extensively studied in literature on star formation since \citet{Stroemgren39}. Stellar radiation generates an ionization front that begins as a R-type front and transforms into a D-type front when its speed slows to twice the sound speed of the ionized gas. Then a strong shock wave forms at the front and recedes from the star at $\sim$$30\kms$. The ionization front decouples from the shock wave and creates a final \ion{H}{2} region that is 1 -- 3 proper kpc in radius for massive Pop III stars residing in low-mass halos. The ionized gas is warm ($\sim$$3 \times 10^4$ K) and diffuse ($\sim$$1 \cubecm$). The shock wave continues to accumulate gas and advance after the star dies. Eventually it stalls in the IGM, but in the process, it reduces the baryon fraction of the halo below one percent \citep{Yoshida06a, Abel07}. Clearly the number of progenitors of a given galaxy as well as the star formation and feedback history of the progenitors will play a role in shaping all of its properties. But how much? If most stars of a galaxy are formed later, will the earliest episodes not be entirely negligible? To start addressing these questions, we have carried out a suite of simulations that include accurate three dimensional radiative transfer and the SN explosions of Pop III stars and have followed the buildup of several dwarf galaxies from those Pop III star hosting progenitors. The Pop III radiative and SN feedback dramatically alter the properties of high redshift dwarf galaxies, and we discuss some of the most striking differences here. We leave a more detailed exposition of star formation rates, star forming environments, and the beginning of cosmic reionization for a later paper. In the following section, we detail our cosmological, radiation hydrodynamics simulations and the star formation algorithm. Then we describe the global characteristics of dwarf galaxies that forms in our simulations in \S\ref{sec:reionResults}. There we also focus on metal enrichment of star forming halos and the IGM, arising from PISNe. In \S\ref{sec:reionDiscuss}, we discuss the implications of our findings on the paradigm of high-redshift galaxy formation by including \hh~chemistry and Pop III star formation and feedback. We summarize in the last section. | \label{sec:reionDiscuss} We find that the combination of Pop III stellar feedback and continued \hh~cooling in \tvir~$<$ 10$^4$ K halos alters the landscape of high-redshift galaxy formation. The most drastic changes are as follows: 1. \textit{Dynamic assembly of dwarf galaxies}--- A striking difference when we include Pop III radiative feedback are the outflows and gas inhomogeneities in the halos and surrounding IGM. The outflows enrich the IGM and reduce the baryon fraction of the $10^4$ K halo as low as 0.05, much lower than the cosmic fraction \Ob/\Om~= 0.17 \citep[cf.][]{Yoshida06a, Abel07}. This substantially differs from the current theories of galaxy formation where relaxed isothermal gas halos hierarchically assemble a dwarf galaxy. For instance in simulation A, there are remarkable filamentary structures and a clumpy ISM. Furthermore, Pop III feedback increases the total baryonic angular momentum of the system up to a factor of 3 without SNe and up to 5 with SNe. 2. \textit{Pop III sphere of influence}--- Pop III feedback is mainly a local phenomenon except its contribution to the UVB. How far its \ion{H}{2} region, outflows, and metal ejecta (if any) will predominately determine the characteristics of the next generation of stars. Highly biased (clustered) regions are significantly affected by Pop III feedback. The first galaxies will form in these biased regions and thus should be significantly influenced by its progenitors. 3. \textit{Dependence on star forming progenitors}--- Although our calculations with SNe only provided an upper limit of metal enrichment, it is clear that the metallicity, therefore metal-line and dust cooling and metal-enriched (Pop II) star formation, depends on the nature of the progenitors of the dwarf galaxy. If the galaxy was assembled by smaller halos that hosted a Pop III star that did not produce a SN, e.g., the galaxy would continue to have a top-heavy initial mass function (IMF). 4. \textit{Complex protogalactic ISM}--- The interplay between stellar and SNe feedback, cold inflows, and molecular cooling produce a truly multi-phase ISM that is reminiscent of local galaxies. The cool, warm, and hot phases are interspersed throughout the dwarf galaxy, whose temperatures and densities can span up to 6 orders of magnitude at a given radius. 5. \textit{Metallicity floor}--- When the halo is massive enough to host multiple sites of star formation, the metal ejecta does not significantly increase the mean metallicity of the host halo. There seems to be a balance between (a) galactic outflows produced from SNe, (b) inflowing metal-enriched gas, (c) inflowing pristine gas, and (d) SNe ejecta that is not blown out of the system. In our high yield models, the metallicity interestingly fluctuates around $10^{-3} Z_\odot$ in the most massive halo when this balance occurs at and above mass scales $\sim$10$^7$\Ms. Clearly the first and smallest galaxies are complex entities, contrary to their low mass and generally assumed simplicity. Our calculations reflect the important role of Pop III stellar feedback in early galaxy formation. These high-redshift galaxies have a $\sim$5--15\% chance of being undisturbed by mergers until the present day, being ``fossils'' of reionization \citep{Gnedin06}. Dwarf spheroidals (dSph) galaxies are some of the darkest galaxies in the universe, having high mass-to-light ratios up to 100 \citep{Mateo98}. Gas loss in dSph's close to the Milky Way or M31 can be explained by gas tidal stripping during orbital encounters \citep{Mayer07}. However there are some galaxies (e.g. Tucana, Cetus) removed from both the Milky Way and Andromeda galaxies and cannot be explained by tidal stripping. In addition to ultraviolet heating from reionization \citep{Bullock00, Susa04} and intrinsic star formation \citep{MacLow99}, perhaps stellar feedback from Pop III stars influenced the gas-poor nature of dSph's. Even at the onset of widespread star formation in the objects studied here, the baryon fraction can be three times lower than the cosmic mean, and the dwarf galaxy may never fully recover from the early mass loss. This initial deficit may play an important role in future star formation within these low-mass galaxies and could help explain the large mass to light ratio in isolated dSph's. With the radiative and chemical feedback from the progenitors of the early dwarf galaxies, we have an adequate set of cosmological ``initial conditions'' to study the transition from Pop III to Pop II stars. In this setup, the current metal tracer field can be used to include metal line cooling. Dust cooling may induce fragmentation of solar mass fragments at metallicities as low as $\sim$10$^{-6}$ at high densities \citep{Schneider06}. However, metal-line cooling might not be important at these low metallicities in diffuse gas. \citet{Jappsen07a} showed that metal-line cooling at metallicities below 10$^{-2} Z_\odot$ in low density gas does not significantly affect the dynamics of a collapsing halo. In a companion paper, \citet{Jappsen07b} found that the fragmentation of a metal-poor ($Z = 10^{-3.5} Z_\odot$) collapsing object may depend more on the conditions, e.g. turbulence and angular momentum distributions, created during the assembly of such a halo than some critical metallicity. Perhaps when the protogalactic gas cloud starts to host multiple sites of star formation, the associated SNe produce sufficient dust in order for a transition to Pop II. In lower mass halos, the SN ejecta is blown out of the halo, and future star formation cannot occur until additional gas is reincorporated into the halo. However in these halos with masses $\gsim$10$^7$\Ms, the SN does not totally disrupt the halo. A fraction of the SN ejecta is contained within halo and could contribute to subsequent sites of star formation. SN ejecta and associated dust could instigate the birth of the first Pop II stars. The metallicities of the first generation of metal-enriched stars could depend on the metal mixing timescales in these early dwarf galaxies. Fortunately the dispersion of heavy elements from SNe into the Galactic ISM is a rich field of study \citep[for a review, see][]{Scalo04}. Metallicity dispersions in local stellar clusters show fluctuations of 5--20\% around the mean, which suggest that the ISM is well-mixed \citep[e.g.][]{Edvardsson93, Garnett00, Reddy03}. In addition, very metal-poor ([Fe/H] $<$ --2.7) halo stars have very little scatter in elemental abundances, suggesting that the dispersal of the first metals that formed low mass stars originated from single starbursts instead of single SN explosions \citep{Cayrel04}. In the ISM, laminar flows and turbulence driven by SN explosions provides the impetus of metal mixing on large scales, which then cascades to smaller and smaller length scales eventually reaching a length scale associated with a Reynolds number $Re = 1$ \citep[e.g.][]{deAvillez02}. In structure formation, turbulence can also arise during the virialization of cosmological halos \citep{Wise07a, Greif08}. After the turbulent cascade creates metallicity gradients on small enough scales, molecular diffusion will homogenize the heavy elements. \citeauthor{deAvillez02} performed AMR simulations that include turbulent diffusivity to study the mixing timescales in the ISM. In their simulations and the ones presented here, numerical diffusion provides the majority of the mixing at the resolution limit. They find that the mixing time is fairly independent of the diffusion scale and could depend on some inertial scale. The origin and magnitude of diffusion does not significantly affect the gas on large scales. In their resolution study, mixing timescales only decrease by 20\% when the resolution is increased by a factor of 4. Thus we could be overestimating the mixing timescales up to a factor of a few because it is difficult to resolve the smallest turbulent scale in cosmological simulations. Furthermore, we expect metals that are newly incorporated into the galaxy to be stirred by the turbulence in assembling halos, which is sustained in virialization and major mergers on the largest scales. As discussed above, the metallicity of the most massive halo fluctuates around $10^{-3} Z_\odot$. This is intriguingly the same value as a sharp cutoff in stellar metallicities in four local dSph's: Sculptor, Sextans, Fornax, and Carina \citep{Tolstoy04, Helmi06}. This differs with the galactic halo stars, whose metal-poor tail extends to $Z/Z_\odot = 10^{-4}$ \citep{Beers05}. We must take care when comparing our results to observations since we made the simplification that every Pop III star produces a PISN \citep[for a detailed semi-analytic model of metal enrichment, see][]{Tumlinson06}. As discussed in \citeauthor{Helmi06}, the galactic halo may be composed of remnants of galaxies that formed from high-$\sigma$ density fluctuations, and dwarf galaxies originate from low-$\sigma$ peaks. In this scenario, the objects (or its remnants) simulated here would most likely reside in galactic halos at the present day. If we attempt to match this metallicity floor of $10^{-4} Z_\odot$ in the galactic halo, this requires $\sim$8\Ms~of metals produced for every Pop III star or roughly one in ten Pop III stars ending in a PISN. More likely, the current nearby dSphs are hosted by larger dark matter halos than we have been able to simulate to date. Hence too simple chemical evolution inspired extrapolations may be premature. Radiative feedback from Pop III stars play an important role in shaping the first galaxies. We studied the effects of this feedback on the global nature of high-redshift dwarf galaxies, using a set of five cosmology AMR simulations that accurately model radiative transfer with adaptive ray tracing. Additionally, we focused on the metal enrichment of the star forming halos and their associated star formation histories. Our key findings in this paper are 1. Dynamical feedback from Pop III stars expel nearly all of the baryons from low-mass host halos. The baryon fractions in star forming halos never fully recover even when it reaches a virial temperature of 10$^4$ K. The baryon fraction is reduced as low as $\sim$0.05 with SNe feedback, three times lower than the cases without stellar feedback. 2. Baryons on average gain angular momentum as they are expelled in feedback driven outflows. When it is reincorporated into the halo, it increased the spin parameter up to a factor of 4 with radiative and SNe feedback. 3. The accurate treatment of radiative transfer produces a complex, multi-phase ISM that has densities and temperatures that can span up to 6 orders of magnitude at a given radius. 4. Pair-instability SN preferentially enrich the IGM to a metallicity an order of magnitude higher than the overdense filaments adjacent to the sites of star formation. 5. Once a SN explosion cannot expel the gas in its host halo, the mean metallicity fluctuates around \zstar{-2.6} as there may be a balance between SN outflows, cold inflows, and contained SNe ejecta. We conclude that Pop III stellar feedback plays an integral part to early galaxy formation as it determines the characteristics of the first galaxies. | 7 | 10 | 0710.3160 |
0710 | 0710.1165_arXiv.txt | Many physical properties of galaxies correlate with one another, and these correlations are often used to constrain galaxy formation models. Such correlations include the color-magnitude relation, the luminosity-size relation, the Fundamental Plane, etc. However, the transformation from observable (e.g. angular size, apparent brightness) to physical quantity (physical size, luminosity), is often distance-dependent. Noise in the distance estimate will lead to biased estimates of these correlations, thus compromising the ability of photometric redshift surveys to constrain galaxy formation models. We describe two methods which can remove this bias. One is a generalization of the $V_{\rm max}$ method, and the other is a maximum likelihood approach. We illustrate their effectiveness by studying the size-luminosity relation in a mock catalog, although both methods can be applied to other scaling relations as well. We show that if one simply uses photometric redshifts one obtains a biased relation; our methods correct for this bias and recover the true relation. | The `configuration space' we use to describe galaxies is large, but galaxies do not fill it. The luminosity, color, size, surface brightness, stellar velocity dispersion, morphology, stellar mass, star formation history and spectral energy distribution of a galaxy are all correlated with one another. These correlations encode important information about galaxy formation, and so quantifying them provides important constraints on models. Current (e.g. SDSS, Combo-17, MUSYC, Cosmos) and planned surveys (e.g. DES, LSST, SNAP) go considerably deeper in multicolor photometry than in spectroscopy, or are entirely photometric. For such surveys, reasonably accurate photometric redshift estimates are or will be made. The question then arises as to which galaxy observables and correlations are affected by the noisy distance estimate associated with a photometric rather than spectroscopic redshift. The most widely studied property is luminosity - clearly, errors in the distance result in incorrect luminosity estimates. If not accounted for, this leads to a biased estimate of the luminosity function (e.g. Subbarao et al. 1996). Hence, there has been considerable effort devoted to the question of how to correct for this bias (e.g. Chen et al. 2003), and the problem has now been solved (Sheth 2007). The next step is to recover an unbiased estimate of not just the luminosity function, but the joint distribution of luminosity, color, size, etc., from photometric redshift datasets. The main goal of the present work is to provide an algorithm which does this for a magnitude limited photometric redshift survey. Because the same distance error which leads to a mis-estimate of the luminosity will produce a correlated mis-estimate of the size, we have chosen to phrase the discussion in terms of the size-luminosity relation - it exhibits all the features of interest. Section 2 illustrates the nature of the problem by showing the bias in the size-luminosity relation which results from treating photometric redshifts as though they were spectroscopic redshifts. This is done by constructing a mock galaxy sample and then perturbing the true redshifts to mimic photometric redshift errors. Section 3 places this problem in the more general context of inverse problems in statistical astronomy, and argues that a deconvolution algorithm, such as that due to Lucy (1974), is well-suited to removing the bias. It shows the result of applying this non-parametric deconvolution technique to a mock galaxy sample. Section~4 provides a maximum-likelihood formulation and solution of the problem. A final section summarizes our findings and discusses possible further studies and applications. Where necessary, we write the Hubble constant as $H_0 = 100h~{\rm km~s}^{-1}~{\rm Mpc}^{-1}$, and we assume a spatially flat cosmological model with $(\Omega_M,\Omega_{\Lambda}, h)=(0.3, 0.7, 0.7)$, where $\Omega_M$ and $\Omega_{\Lambda}$ are the present-day densities of matter and cosmological constant scaled to the critical density. We use $D_{\rm L}(z)$ to denote the luminosity distance; the angular diameter distance is $D_{\rm A}(z) = D_{\rm L}(z)/(1+z)^2$. | We presented two algorithms for reconstructing the intrinsic correlations between distance-dependent quantities in apparent magnitude limited photometric redshift datasets. One was a generalization of the non-parametric $V_{\rm max}$ method (Section~\ref{Vmax}), and the other used a maximum-likelihood approach (Section~\ref{ml}). Both our reconstruction methods assume that the distribution of photo-$z$ errors is known accurately. In practice, this means that spectroscopic redshifts are available for a subset of the data. The question then arises as to whether or not the number of spectra which must be taken to specify the error distribution reliably is sufficient to also provide a reliable (spectroscopic!) estimate of these scaling relations. If so, what is the basis for deciding that it is worth reconstructing these relations from the photo-$z$ data? This is the subject of work in progress, although the methods presented in this paper assume that such reconstructions will indeed be necessary. E.g., if the spectra are not simply a random subset of the magnitude limited photometric sample, then it may be difficult to quantify and so correct for the selection effects associated with the spectroscopic subset. We used the size-luminosity relation in a mock catalog which had realistic choices for the correlation to illustrate the biases which are present and must be corrected if photometric redshift datasets are to provide reliable estimates of galaxy scaling relations (Figures~\ref{M_Me_R_Re_early_types} and~\ref{M_R_correlation_rec}). We showed that our iterative deconvolution scheme provides a simple and reliable correction of this bias (Figures~\ref{M_R_correlation_rec}--\ref{p_R_given_M_rec}). Note that although we have illustrated our methods using a 2-dimensional distribution, the extension to $n$-correlated variables is trivial. Because our algorithm permits the accurate measurement of many scaling relations for which spectra were previously thought to be necessary (e.g. the color-magnitude relation, the size-surface brightness relation, the Photometric Fundamental Plane), we hope that our work will permit photometric redshift surveys to provide more stringent constraints on galaxy formation models at a fraction of the cost of spectroscopic surveys. Our results may have other applications. For example, Bernardi (2007) has highlighted a bias associated with the correlation between stellar velocity dispersion $\sigma$ and luminosity $L$ which arises if the distance indicator used to estimate $L$ is correlated with $\sigma$ (as may happen in the local Universe, where peculiar velocities make spectroscopic redshifts unreliable distance estimators). It may be that the methods presented here would allow an accurate reconstruction of the true relation from the biased one. This is the subject of on-going work. | 7 | 10 | 0710.1165 |
0710 | 0710.1023_arXiv.txt | We report on the discovery of the $z=1.016$ cluster RzCS 052 using a modified red sequence method, followup spectroscopy and X-ray imaging. This cluster has a velocity dispersion of $710 \pm 150$ km s$^{-1}$, a virial mass of $4.0 \times 10^{14}\ M_{\odot}$ (based on 21 spectroscopically confirmed members) and an X-ray luminosity of $(0.68 \pm 0.47) \times 10^{44}$ ergs s$^{-1}$ in the [1-4] keV band. This optically selected cluster appears to be of richness class 3 and to follow the known $L_X - \sigma_v$ relation for high redshift X-ray selected clusters. Using these data, we find that the halo occupation number for this cluster is only marginally consistent with what expected assuming a self-similar evolution of cluster scaling relations, suggesting perhaps a break of them at $z\sim1$. We also rule out a strong galaxy merging activity between $z=1$ and today. Finally, we present a Bayesian approach to measuring cluster velocity dispersions and X-ray luminosities in the presence of a background: we critically reanalyze recent claims for X-ray underluminous clusters using these techniques and find that the clusters can be accommodated within the existing $L_X - \sigma_v$ relation. | Clusters of galaxies are not only a powerful tool to study galaxy evolution but can also be used to constrain cosmological parameters, resolving several parameter degeneracies (e.g., Allen et al. 2004; Albrecht et al. 2006). In particular, clusters at high redshifts ($z > 1$), of which only a handful are currently known, provide the greatest leverage in determining the nature of the acceleration constant (e.g., Rapetti 2007). These determinations, however, rely on an accurate estimate of the cluster mass, whose uncertainty is arguably the dominant contributor to the error budget in deriving cosmological parameters from cluster statistics (Henry 2004; Albrecht et al. 2006). Ideally, one wish to apply the virial theorem to get a direct measurement of cluster masses. In fact, the dark matter velocity dispersion is an extremely good tracer of the halo masses in all simulations (Evrard et al. 2007), and galaxies are nearly unbiased velocity tracer (Evrard et al. 2007 and references therein; Rines Diaferio \& Natarajan 2007), in good agreement previous works (Biviano et al. 2006, Tormen et al. 1997). The measurement of the cluster velocity dispersion requires a large number of radial velocities, which are observationally expensive to obtain, especially for high redshift clusters. For this reason and because each mass estimator carries some key informations, more commonly the scaling between pairs of more easily observable mass-related quantities is studied, such as X-ray luminosity, temperature or the $Y_X$ (Kravtsov et al. 2006) parameter, or optical richness. These studies often look for outliers, however their search is blessed by data limitation: for example in the search of clusters X-ray dim for their optical richness, Donahue et al. (2001) and Gilbank et al. (2004), both mostly worked with putative clusters (i.e. not spectroscopically confirmed) and X-ray undetections. \begin{figure*} \psfig{figure=RzCS052_150.ps,width=12truecm,clip=} \caption[h]{True-color ($z'[3.6][4.5]$) degraded-resolution (to make galaxies not too small when printed) image of a region of a 24 Mpc$^2$ area around RzCS 052. Spectroscopically confirmed clusters and isodensity contours for red galaxies are also marked. Note the number density contrast of reddish galaxies between the cluster center and the right part of the image. The ruler is 1 arcmin long; North is up, East is to the left.} \end{figure*} Only few works directly address the relative quality of different mass estimators with velocity dispersion: Borgani \& Guzzo (2001) compare the scatter of two mass estimators, X-ray luminosity and richness, and found that the former is a better mass tracer than the latter when the former is uniformly measured and the latter is taken from a 50 years old paper reporting eye-estimate of the cluster optical richness (the Abell 1958 catalog). In both CNOC and nearby clusters, mass correlates better with richness than with X-ray luminosity (Yee \& Ellinson 2003; Popesso et al. 2005). Eke et al. (2004) found that optical luminosity is a better proxy of mass than velocity dispersion in common conditions, i.e. when velocities are available for a small sample of galaxies. A related issue, which we will examine below, is whether it exists clusters X-ray dim for their mass (velocity dispersion), (e.g. Lubin, Mulchaey \& Postman 2004, Fang et al. 2007, Johnson et al. 2006). The relation between richness and mass has received some recent attention in the form of the halo occupation function (Berlind \& Weinberg 2002; Lin et al. 2004 and references therein) whose first moment is the halo occupation number (HON), the average number $N$ of galaxies per cluster of mass $M$. In order to address the evolution of the HON, velocity dispersion information is often unavailable for a large cluster sample, mass and cluster size are inferred from other mass-related quantities (for example the X-ray temperature), and assumed to evolve self similarly. The evolution of the HON with redshift is still unclear: the initial study by Lin et al. (2004) claimed that the HON increases at high redshift, but Lin et al. (2006) find evidence that it does not evolve strongly out to $z \sim 1$, suggesting that the galaxy population in clusters was established and assembled at early epochs. Muzzin et al. (2007) confirms the above, with a sample of reduced redshift leverage and hence reduced evolution sensitivity, but available velocity dispersion information. \begin{figure*} \centerline{\psfig{figure=figpapI1.ps,width=8truecm,clip=} \psfig{figure=figpapI2.ps,width=8truecm,clip=}} \caption[h]{Spectra of RzCS 052 members coming the VLT run. We have vertically shifted the spectra and zoomed on a reduced wavelength range for display purpose.} \end{figure*} \begin{figure*} \psfig{figure=cmRz.ps,width=18truecm,clip=} \caption[h]{ {\it Left:} Colour--magnitude diagram for galaxies (close circles) within 1 arcmin from the cluster center or with a known redshift (red closed circles for members, blue crosses for interlopers) using CTIO discovery data. $R$ and $z'$ mag completeness limits are show with dashed lines. The green line is the expected CM at the cluster redshift, from Kodama \& Arimoto (1997) {\it Right:} Colour histograms of galaxies within 2 arcmin from the cluster center (solid histogram), and of the average control field (measured on a 0.36 deg$^2$ area, dashed histogram), normalized to the cluster area. A clear excess is seen, especially at $R-z>1.5$ mag. The hashed histogram is the colour distribution of spectroscopically confirmed member galaxies. In both panels, a few objects with spectroscopic redshift are missing because their fall on bad CCD regions or have extreme colours.} \end{figure*} Here, we present the photometric discovery, spectroscopic confirmation and X-ray properties of a new $z=1.016$ cluster of galaxies (RzCS 052), a cluster optically rich but undetected in the XMM-LSS survey (Pierre et al. 2007), and hence possibly X-ray dark (i.e. dim for its mass). We derive its global properties (richness, X-ray luminosity, velocity dispersion and mass) and study these in the context of cluster scaling relations ($L_X - \sigma$, HON) at high redshift. In particular, we test the claim that the HON (the way galaxies populate cluster-scale haloes) has not changed $z \sim 1$ (Lin et al. 2006) under far less assumptions than the original claim. We also present a Bayesian approach to the determination of cluster velocity dispersion and X-ray luminosity and use it to critically examine recent claims about the existence of underluminous X-ray clusters. A companion paper (Andreon et al. 2007) addresses the use of RzCS 052 as a laboratory for studying galaxy formation and evolution. We adopt $\Omega_\Lambda=0.7$, $\Omega_m=0.3$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Magnitudes are quoted in their native photometric system (Vega for $R$, SDSS for $z'$). | We have identified a distant cluster from a modified red-sequence method and followed it up spectroscopically. RzCS 052 is a richness class 3 cluster at $z=1.016$ with a velocity dispersion of $710 \pm 150$ km s$^{-1}$ and an X-ray luminosity of $0.68 \pm 0.47 \times 10^{44}$ ergs s$^{-1}$. In spite of its optical detection, RzCS 052 obeys to the high redshift $L_X-\sigma_v$ relationship as other X-ray selected clusters to the high redshift $L_X-\sigma_v$ relationship, whereas in principle variations in the dynamical state of the clusters or in the thermal history of the intracluster medium may have moved it away from the $L_X-\sigma_v$ relation. Analysis of the $N-M$ scaling shows that RzCS 052 has the right number of galaxies (actually, a bit less) that it should have for its mass, ruling out intense merging (among galaxies) activities in clusters from $z=1$ to today, in agreement with Lin et al. (2006). We present a Bayesian approach to measuring cluster velocity dispersions (most useful for sparsely sampled data and in presence of a background) and X-ray luminosities or upper limits (essential in the case of poorly determined parameters). Critical re-analysis of the data of clusters/groups claimed to be outliers of the $L_X-\sigma_v$ relationship leads to conclude that there are no known thus far examples of clusters X-ray underluminous for their velocity dispersion. The above result is quite reassuring for ongoing X-ray surveys: there is thus far no example of cluster missed because an anomalous $L_X$ for the cluster mass. \begin{table*} \caption{Luminosity and velocity dispersion of clusters at $z>0.8$} \begin{tabular}{llcrcrr} \hline Name & z & $\log L_X [1-4]$ & ref & $ \sigma_v$ & N \hfill & ref \\ \hline RXJ1716+6708 & 0.813 & $44.68\pm0.03$ & 12 & $1522 \pm 180 $& 37 & 1 \\ RXJ1821.6+6827 & 0.816 & $44.62\pm0.01$ & 2 & $775 \pm 122 $& 18 & 2 \\ MS1054-30321 & 0.830 & $44.91\pm0.04$ & 12 & $1153 \pm 80 $& .. & 3 \\ RXJ0152-1357S & 0.830 & $44.42\pm0.02$ & 12 & $737 \pm 126 $& 18 & 4 \\ RXJ0152-1357N & 0.835 & $44.59\pm0.03$ & 12 & $919 \pm 168 $& 16 & 4 \\ RzCS 530 & 0.839 & $44.20\pm0.09$ & 5 & $780 \pm 126 $& 17 & 5 \\ 1WGA1226+3333 & 0.890 & $45.17\pm0.01$ & 12 & $997 \pm 245 $& 12 & 6 \\ Cl1604+4304 & 0.900 & $44.00\pm0.06$ & 14,7 & $962 \pm 141 $& 67 & 7 \\ RzCS 052 & 1.016 & $43.83\pm0.37$& this work & $710 \pm 150 $& 21 & this work \\ RXJ0910+5422 & 1.106 & $44.00\pm0.05$ & 12 & $675 \pm 190 $& 25 & 8 \\ RXJ1252-2927 & 1.237 & $44.37\pm0.07$ & 12 & $747 \pm 79 $& 38 & 9 \\ LynxW & 1.270 & $43.70\pm0.23$ & 12 & $650 \pm 170 $& 9 & 10 \\ 1WGAJ2235.3 & 1.393 & $44.57\pm0.05$ & 11 & $762 \pm 265 $& 12 & 11 \\ \hline \end{tabular} \hfill \break \footnotesize{RzCS 530 is also known as XLSSC 003; References: 1: Gioia et al. (1999); 2: Gioia et al. (2004a); 3: Gioia et al. (2004b); 4: Demarco et al. (2005); 5: Valtchanov et al. (2004); 6: Maughan et al. (2004); 7: Gal \& Lubin (2004); 8: Mei et al. (2006); 9: Demarco et al. (2007); 10: Stanford et al. (2001); 11: Mullis et al. (2005); 12: Ettori et al. (2004); 14: Lubin et al. (2004). \hfill \break } \end{table*} | 7 | 10 | 0710.1023 |
0710 | 0710.1812_arXiv.txt | \label{Sec:intro} \setcounter{equation}{0} \subsection{Experiments and theories} \label{Sec:ExTh} An ongoing and forthcoming array of experiments (e.g. WMAP, \cite{Spergel:2006hy}, SDSS \cite{York:2000gk, Tegmark:2003uf, Tegmark:2006az, McDonald:2004xn}, SNLS \cite{Astier:2005qq}, ACBAR \cite{Kuo:2006ya}, Planck \cite{Planck}, ACT \cite{ACT}, Spider \cite{Spider}) is measuring the cosmic microwave background (CMB) and the large scale structure of the universe with unprecedented precision. This provides exciting opportunities to reveal the nature of the early universe and the underlying fundamental theories. The leading theoretical candidate for creating the initial conditions of our universe is inflation \cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi}. However inflation remains a paradigm, which can be implemented by a variety of models underpinned by differing microphysical constructions; as the constraints from data tighten, there is the hope that we might identify the specific scenario that describes our universe. With the natural ingredients of such model-building being supergravity and string theory, the process of better measuring the properties of the early universe is also a process of understanding better the theory of quantum gravity. In contrast to recent debates on the predictivity of the string theory landscape, here we use a more conventional approach to investigate the predictivity of string theory by studying the properties and exploring the dynamics of our own vacuum. We first scan the parameter space of inflationary models subject only to the requirement that they provide enough inflationary $e$-folds to solve the flatness and horizon problems. This is because the natural creation of a homogeneous and isotropic universe is the leading problem that we want to solve, and is perhaps the most attractive feature of the inflationary paradigm. After that, we study the observational consequences of all the viable parameter spaces, with the goal of looking for distinctive signatures. Some of these can be compared with observations and used to narrow down the parameter space. Despite of the vastness of all possible vacua in the string landscape, this process can be rather effective since certain observational features rely on distinctive dynamics. As we will see, such dynamics can either be field-theoretic with strong motivations from string theory, or completely stringy in nature. There are many candidate observable signatures in inflationary models. The most generic ones are the amplitude of the primordial power spectrum and its spectral index. Since most viable models built from a fundamental theory have adjustable parameters to fit these two observables, this leaves a large number of viable models that are consistent with the data, and even leaves the nature of the inflaton field ambiguous. In principle, Nature is not obligated to provide more information within our experimental abilities, and indeed there is no evidence for further parameters required to describe the current data. But anticipating her generosity, possible distinctive observables that might be measurable in the future include the scale-dependence (``running'') of the spectral index, departures from Gaussianity of the primordial fluctuations, a tensor contribution to the primordial power spectrum, and cosmic strings. These will be crucial to successfully carry out the program that we have outlined. With the rapidly improving quality of cosmological data, it will become increasingly interesting to implement the above program by comparing specific models to data, starting directly from microscopic parameters of theories. Modern cosmological data analyses make use of the powerful method known as Markov Chain Monte Carlo (MCMC) to implement the comparison to data, providing an efficient way of estimating posterior distributions of the microscopic parameters. However, in practice, when directly using microscopic parameters as MCMC parameters, highly non-linear relationships between the parameters and observables may introduce severe obstacles for MCMC to efficiently search the parameter space. Therefore, a reparameterization according to the specific nature of the model often becomes necessary. So instead of a straightforward exercise, implementing MCMC becomes a rather interesting model-dependent art. It is also a purpose of this paper to use an example to illustrate this process and extract certain model-independent procedures of such reparameterization which may be of more general interest. \subsection{Brane inflation} The inflationary models that we study in this paper belong to the brane inflation scenario proposed by Dvali and Tye \cite{Dvali:1998pa,HenryTye:2006uv}.\footnote{For recent reviews on other types of string inflation models, see Ref.~\cite{Burgess:2007pz,Linde:2007fr,Kallosh:2007ig,Cline:2006hu}.} We are interested in these models precisely because they can give rise to a large number of distinctive observational signatures. This happens even in the simplest scenarios that provide inflation. One of the most important reasons that makes it possible is that brane inflation can be achieved via two different mechanisms, namely slow-roll and Dirac-Born-Infeld (DBI) inflation. The original models of brane inflation \cite{Dvali:1998pa,Burgess:2001fx,Dvali:2001fw,Kachru:2003sx} are slow-roll inflationary models \cite{Linde:1981mu,Albrecht:1982wi}, where branes and anti-branes slowly approach each other in a flat potential. A model that uses this mechanism in the framework of the string theory flux compactification \cite{Douglas:2006es} is studied by Kachru, Kallosh, Linde, Maldacena, McAllister and Trivedi (KKLMMT) \cite{Kachru:2003sx}. As in the F-term inflation models in supergravity \cite{Lyth:1998xn}, it is found that the generic shape of the potential is too steep to achieve the slow-roll inflation, in this case due to the moduli stabilization. Again, similar to those supergravity models, it is possible that several contributions to the potential manage to cancel to a certain precision so that the potential becomes sufficiently flat. There are effective parameters in the model controlling the inflaton mass that can be adjusted to fit the observed spectral index \cite{Bean:2007hc}. The running of the spectral index, non-Gaussianities and tensor modes are all too small to be observed in the near future.\footnote{Some observables become measurable if there are sharp features in the potential \cite{Adams:2001vc,Peiris:2003ff,Covi:2006ci,Chen:2006xj,Hailu:2006uj}. In addition, there are other important observational possibilities of brane inflation -- cosmic strings and those related to reheating \cite{Polchinski:2004ia,HenryTye:2006uv} -- which apply to both slow-roll and DBI inflation.} Another inflationary mechanism that is so far uniquely found in brane inflation is the DBI inflation \cite{Silverstein:2003hf,Alishahiha:2004eh,Chen:2004gc,Chen:2005ad}. In DBI inflation, the rolling velocity of inflaton branes is not determined by the shape of the potential but by the speed-limit of the warped internal space. Such warped spaces are naturally present in the extra dimensions due to fluxes used to stabilize the string compactification \cite{Giddings:2001yu}. The first model that uses such a mechanism is that of Silverstein, Tong and Alishahiha (STA) \cite{Silverstein:2003hf,Alishahiha:2004eh}. In this model, as the branes roll into a throat from the UV side of the warped space under a quadratic potential, its velocity gets restricted by the large warping in the IR side of the warped space. However, instead of having a potential with a generic mass term, a rather steep potential, characterized by a large inflaton mass, is required to achieve this UV DBI inflation. The reason is that, when the branes enter from the UV side of the warped space in the GKP-type warped compactification \cite{Giddings:2001yu}, the energy provided by the antibranes sitting at the IR side is not large enough to drive DBI inflation even if there is the speed-limit, since the antibrane tension has been warped down correspondingly. Therefore an extra, steep, potential has to be added to raise the inflationary energy. In addition, embedded in the same warped compactification, the model generates large non-Gaussianities that exceed the experimental bound \cite{Bean:2007hc,Peiris:2007gz}, as well as excessive probe brane backreaction which we will address in Sec.~\ref{Sec:UV}. This is because in this model, the levels of non-Gaussianity and probe brane backreaction sensitively depend on the inflaton value, and it is viable only if the inflaton field is of (super-)Planckian size. However, the range of the inflaton field is restricted by some geometric conditions of the compactification and is sub-Planckian \cite{Chen:2005fe,Chen:2006hs,Baumann:2006cd,Bean:2007hc,Peiris:2007gz}. To fully make use of the speed-limit of the warped space, it is better to make the branes roll out from the IR end, and use antibranes in other throats to provide the inflationary energy. In this way the speed-limit of the branes and the inflationary energy become relatively independent of each other, leaving a rather flexible shape of the inflaton potential which has been the main problem of model-building. This is the model proposed in Ref.~\cite{Chen:2004gc,Chen:2005ad}. It can be generically realized in the multi-throat brane inflation scenario \cite{Chen:2004gc}. It happens that in this IR DBI inflation model, the large non-Gaussianities can also be small enough to satisfy the current observational bound \cite{Chen:2005fe}. This is partly because no matter how small the warp-factor (and consequently, how big the non-Gaussianity) the branes begin with, the level of non-Gaussianity decreases as the branes roll out and approaches its minimal value at the end of the inflation. Therefore, in the segment of the warped space traversed during the last 60 $e$-folds, the level of non-Gaussianity is among the smallest in the entire DBI inflation trajectory. Moreover, the geometric conditions that put a tight constraint on the STA model are automatically satisfied in the IR DBI model and has no effect on the non-Gaussianities. Besides providing a speed-limit to the inflaton, another important property of warped space is the reduction of the local fundamental string scale \cite{Randall:1999ee}. This turns out to have important consequences on density perturbations in DBI inflationary models. During the epoch when the string scale is red-shifted below the Hubble parameter, the quantum fluctuations on the inflaton branes become stringy.\footnote{Notice that such a stringy phase only happens in the inflaton sector, which is the deep IR side of a warped space with energy density of order $H^4$, so it does not backreact significantly on the Hubble expansion. We also note that such a stringy phase will backreact on the IR side of the warped geometry, but it is estimated that this still leaves a large enough portion of the geometry for DBI inflation to take place \cite{Chen:2005ad,Chen:2006ni}. We will discuss this more in Sec.~\ref{Sec:IR} \& \ref{Sec:Power}.} The density perturbations are no longer fully described by the usual field theory approximation, and acquire distinctive stringy signatures. In the IR DBI model, this stringy phase corresponds to earlier inflationary $e$-folds, and therefore larger scales in the sky. It is estimated that such a phase transition will give rise to a large transient (regional) running of the spectral index \cite{Chen:2005ad,Chen:2005fe}. In this paper, we make this prediction more quantitative and compare it to observations. \subsection{Outline} Following the strategy that we outlined in Sec.~\ref{Sec:ExTh}, in this paper, we first summarize the overall features of brane inflation using phase diagrams that describe the parameter spaces spanned by both inflationary mechanisms, \emph{i.e.} slow-roll vs. DBI (Sec.~\ref{Sec:phase}), reviewing the key observational predictions in the different parts of the parameter space. The main focus of this paper is to compare the IR DBI brane inflation model to observations (Sec.~\ref{Sec:IRDBI}). We derive analytical and numerical model predictions for the shape of the power spectrum, non-Gaussianity, and tensor modes, giving a quantitative estimate of the effect of the Hubble-expansion-induced stringy phase transition on density perturbations (Appendix \ref{Sec:Width}). We then proceed to compare these results to the observational data from cosmic microwave background and large scale structure (Sec.~\ref{Sec:MCMC}). We outline how such a comparison should be generally implemented using MCMC. The current data give a number of interesting constraints on the microscopic parameters of the model (Sec.~\ref{conc}), including the mass of the brane moduli potential, the fundamental string scale, the charge or warp factor of throats, and the number of the mobile branes. We also quantify some distinctive observable signatures of this model, such as the level of the non-Gaussianity and the running of the spectral index. We discuss how the latter is observationally different from two other cases that may also give large running spectral index: slow-roll inflation with mild features on potential (Appendix \ref{Sec:Mild}), and slow-roll or DBI inflation with a non-Bunch-Davies vacuum (Appendix \ref{Sec:nonBD}). These results illustrate how string theory can make testable predictions which might be subject to observational constraints. For convenience, all the variables used in this paper are summarized in Table~\ref{Tb:variables}. \begin{table*}[!ht] \begin{center} {\small \begin{tabular}{cll} \hline Variable & Description & Notes\\ \hline \hline $\mpl$ & 4d reduced Planck mass & $\mpl = (8\pi G)^{-1/2}=2.4\times 10^{18} {\rm GeV} $\\ $m_s$ & Mass scale of fundamental strings & $m_s \equiv \alpha'^{-1/2}$ \\ $g_s$ & String coupling & $g_s<1$ \\ $T_3$ & D3-brane tension & Eq.~(\ref{RTDef}) \\ $R$ & Length scale of warped throat & Eq.~(\ref{RTDef}) \\ $M$, $K$ & Flux numbers in warped throat & Integers \\ $n_A$ & Number of antibranes in A-throats & \\ $n_B$ & Number of branes (inflatons) in B-throat & \\ $N_B$ & Effective charge of B-throat & $N_B = a_B MK$ \\ $a_B$ & Multiplicative factor from orbifolding & $a_B\sim 1$ in data analysis \\ $\lambda_B$ & $\lambda_B = n_B N_B/2\pi^2$ & Eq.~(\ref{lambdaDef}) \\ $r$ & Radial coordinate of throats & \\ $\phi$ & Canonical inflaton field & $\phi=r\sqrt{nT_3}$ \\ $h_A$ & Minimum warp factor of A-throat & \\ $h_B(\phi)$ & Warp factor at location $\phi$ in B-throat & $h\equiv r/R=\phi R/\sqrt{\lambda}$ \\ $\beta$ & Characterization of shape of potential & Inflaton mass $m^2 = \beta H^2$ \\ $\gamma$ & Lorentz factor of inflaton & \\ $c_s$ & Sound speed in 4d & $c_s = 1/\gamma$ for DBI inflation \\ $V_0$ & Inflationary energy density & \\ $N_e$ & Number of $e$-folds to the end of inflation\\ $N_e^{\rm DBI}$ & Number of $e$-folds to the end of IR DBI inflation& \\ $N_{\rm tot}^{\rm NR}$ & Total $e$-folds of non-relativistic roll inflation & Typically fast-roll \\ $k_c$ & Critical scale of the stringy phase transition & Eq.~(\ref{kcdef})\\ $N_c$ & Critical DBI $e$-fold at $k_c$ & Eq.~(\ref{Ncdef})\\ $P(k)$ & Power spectrum &\\ $r_{TS}$ & Tensor to scalar ratio & \\ $f_{NL}^{\rm eq}$ & Estimator of the non-Gaussianity & Equilateral shape\\ \hline \end{tabular} } \caption{\label{Tb:variables} \small Description of variables. In the text, subscripts $A$ and $B$ are frequently added to some of the variables, referring to the quantities of the A- or B-throat.} \end{center} \end{table*} | \label{conc} We conclude by highlighting the main results and discussing their physical implications. The quoted ranges are at the $95\%$ confidence level, and we have combined constraints from both the power spectrum and non-Gaussianity. The detailed $68\%$ and $95\%$ CL marginalized constraints and the maximum likelihood values are listed in Table \ref{Tab_cosmo}. \subsection{Microscopic parameters} \begin{itemize} \item {\em Shape of the inflaton brane moduli potential: $1.3<\beta < 3.7$.} The lower bound is due to constraints from the power spectrum, while the upper bound is due to the non-Gaussianity constraint. It is encouraging that, while IR DBI inflation can happen for a range of $\beta$ that varies over nearly 10 orders of magnitude, $0.1 \lesssim\beta <10^9$ (see Eq.~(\ref{mrange})), comparison with data picks out a very small range around $\CO(1)$ which is generically expected theoretically. This makes an explicit construction of such potentials a more interesting question. \item {\em Fundamental string scale: $-9.6 < \log_{10}(m_s/\mpl)/g_s^{1/4} < -5.3 $.} The upper bound on the string scale is due to the large charge, and hence length scale, of the B-throat required to fit the amplitude of the density perturbations. The lower bound is due to the fact that a smaller string scale tends to increase the total number of $e$-folds of non-relativistic fast-roll inflation, and make the running of the spectral index too large (Fig.~\ref{Fig:2d}). The model prefers an intermediate fundamental string scale, $10^8\ {\rm GeV} < m_s/g_s^{1/4} < 10^{13}\ {\rm GeV}$, and therefore an intermediate large volume compactification, $8.9\times 10^7 < V^{1/6} \mpl <4.8\times 10^{13}$, where $V$ is the compactification volume. \item {\em B-throat charge: $8.8< \log_{10} N_B < 10.4 $; Number of inflaton branes: $3.9<\log_{10} n_B< 5.1$.} In terms of the GKP-type warped compactification, this implies flux numbers $K\sim M \sim \sqrt{N_B} \sim \CO(10^5)$. Explicit construction remains an open question as discussed in Sec.~\ref{Sec:OpenQu}. In the multi-throat brane inflation scenario, inflaton branes are generated from flux-antibrane annihilation. The number of branes generated in this process is roughly determined by the flux number $M$. Indeed, a small number of inflaton branes is ruled out by the data. \item {\em A-throat minimum warp factor: $-2.4<\log_{10} h_A \le 0$.} This is from combining the constraint on $n_B$ and $n_A h_A^4$, $h_A = (n_A h_A^4/n_B)^{1/4}$. A smaller $h_A$ leads to larger $N^{\rm NR}_{\rm tot}$ and larger running of the spectral index (Fig.~\ref{Fig:2d}). So the A-throat tends to be short. This makes tunneling reheating possible, where many interesting phenomena can occur, such as an intermediate matter-dominated epoch. \end{itemize} \subsection{Secondary derived parameters} \begin{itemize} \item {\em Inflationary phases.} In this model, not all $e$-folds comes from IR DBI inflation. The last $13< N^{\rm NR}_{\rm tot}< 24$ $e$-folds come from non-relativistic fast-rolling inflation, which is possible because inflatons are close to the top of the potential. \item {\em The stringy phase transition.} The Hubble-expansion induced stringy phase transition happens at the largest scales in the sky, $-6.0<\log_{10} k_c/{\rm Mpc} <-2.9$. However its impact on density perturbations extends over to shorter scales, such as generating a transient large running of the spectral index. \item {\em Inflation scale: $-10.0<\log_{10}V_0^{1/4}/\mpl<-5.1$.} This gives a very small tensor to scalar ratio $r_{TS}<10^{-13}$. \item {\em Cosmic string tension: $-23<\log_{10} G\mu_D +\log_{10}g_s^{1/2} < -14$.} Here the cosmic strings refer to the D-strings left over from the brane-antibrane annihilation in the A-throat, whose tension is $G\mu_D= (m_s h_A/g_s^{1/4} \mpl)^2/(16\pi^2 g_s^{1/2})$. There is an unconstrained freedom coming from the additive factor $\log_{10}g_s^{1/2}$, but it is not expected to give any significant contributions. The F-string tension differs by a factor of $g_s$, $\mu_F = g_s \mu_D$. \end{itemize} \subsection{Observational predictions} \begin{itemize} \item {\em Large, but regional, running of spectral index: $-0.046< dn_s/d\ln k(k=0.02/{\rm Mpc}) <-0.010$.} A reconstructed full-scale power spectrum and the running of the spectral index are shown in Fig.~\ref{Fig:Pk} \& \ref{Fig:nrun}. This prediction is stringy in nature. A better understanding of the theoretical details and better measurements of both the power spectrum and non-Gaussianities on the relevant scales may reveal finer structures. In future experiments, Planck is expected to achieve $\sigma(dn_s/d\ln k) = 0.005$ \cite{Planck}. \item {\em Large non-Gaussianities: $-272<f_{NL}^{\rm eq}(k=0.02/{\rm Mpc}) <-70$.} A reconstructed full-scale prediction is in Fig.~\ref{Fig:fnl}, which shows the running of the non-Gaussianities. This prediction is strictly speaking field-theoretic, but with strong string theory motivations, such as warped compactification and the DBI brane action. This field theoretic regime is $k>k_c$; the theoretical analysis for non-Gaussianities at $k\lesssim k_c$ is currently unavailable and remains an interesting open question. In future experiments, on CMB scales, Planck can achieve $\sigma(f^{\rm eq}_{NL}) =67$ \cite{Smith:2006ud,Hikage:2006fe}; on large scale structure scales, some high-$z$ galaxy surveys can reach similar or better precision \cite{Sefusatti:2007ih}. \end{itemize} As seen from these results, constraints from cosmological data, and even relatively loose constraints such as the non-Gaussianity constraint, are already putting strong restrictions on models which aim to provide self-consistent microphysical descriptions of the early universe. With the bounty of precision cosmological data expected in the future, the hope of probing not just field-theoretic, but string-theoretic early universe physics burns brightly. \medskip | 7 | 10 | 0710.1812 |
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0710 | 0710.5600_arXiv.txt | We study magnetic effects induced by rigidly rotating plates enclosing a cylindrical MHD Taylor-Couette flow at the finite aspect ratio $H/D=10$. The fluid confined between the cylinders is assumed to be liquid metal characterized by small magnetic Prandtl number, the cylinders are perfectly conducting, an axial magnetic field is imposed $\Ha \approx 10$, the rotation rates correspond to $\Rey$ of order $10^2-10^3$. We show that the end-plates introduce, besides the well known Ekman circulation, similar magnetic effects which arise for infinite, rotating plates, horizontally unbounded by any walls. In particular there exists the Hartmann current which penetrates the fluid, turns into the radial direction and together with the applied magnetic field gives rise to a force. Consequently the flow can be compared with a Taylor-Dean flow driven by an azimuthal pressure gradient. We analyze stability of such flows and show that the currents induced by the plates can give rise to instability for the considered parameters. When designing an MHD Taylor-Couette experiment, a special care must be taken concerning the vertical magnetic boundaries so they do not significantly alter the rotational profile. | Motion of a fluid confined between two concentric, rotating cylinders is a classical problem in hydrodynamics and, if the fluid is conducting and an external magnetic field is applied, magnetohydrodynamics (MHD). The flow of this type, usually referred to as the Taylor-Couette flow, has been first studied by Couette~\citep{Couette1888} and later was subject of a seminal work by Taylor~\citep{Taylor1923}, who experimentally confirmed theoretical results of a linear stability analysis. In the field of MHD, an important work was done by Velikhov~\citep{Velikhov59} who has shown that for the conducting fluid a weak magnetic field can play a destabilizing role and can lead to an instability which today is called magnetorotational instability (MRI~\citep{Jietal01}). When studying the Taylor-Couette system it is common to assume some simplifications, the small gap approximation or large aspect ratio. In the former it is assumed that the gap between the cylinders $D = \Rout-\Rin$ is small compared to the radii, i.e., $D/\Rout \ll 1$, this allows the neglect of terms of order $1/R$, $R$ being distance from the center of rotation. When considering the large aspect ratio, one assumes that the height of the cylinders $H$ is much larger than the gap width $\Gamma=H/D \gg 1$, which guarantees that a secondary flow due to the plates bounding the cylinders is insignificant and does not disturb the rotational profile of the fluid. On the other hand, there is also plenty of work done for small aspect ratio $\Gamma \approx 1$, where the rigidly rotating end-plates play crucial role and simply introduce a new class of problems. When $\Gamma$ becomes an important parameter it is possible to observe a wide family of different states (including non-axisymmetric ones or peculiar asymmetric patterns -- anomalous modes) for the same parameters, so that the observed results depend on their path through the parameters space from an initial state. Therefore this system is an excellent subject to the bifurcation theory~\citep{1988JFM...191....1P, LopezMarques03, Mullinetal02, Furukawaetal02, Kageyamaetal04, Youd06}. In the present work we focus on the case of wide gap $\Rin/\Rout=1/2$ and $\Gamma=10$ which is an intermediate aspect ratio, between very short and long containers, yet in purely hydrodynamical contest the influence of the vertical boundaries is small, at least for Reynolds numbers of order \ord{10^2-10^3}. However, if the rotation rates are large enough, so that the corresponding Reynolds number is \ord{10^5} and larger, the plates can easily dominate the flow in the entire container. This is due to the Taylor-Proudman theorem, from which follows that in rapidly rotating systems the flow tends to align itself along the axis of rotation. For such rotations it is necessary that $\Gamma$ would have to be several thousand in order to obtain the rotational profile which is not profoundly altered by the end-plates~\citep{HollerbachFournier04}. Results of a recent MRI experiment PROMISE~\citep{MRIexpA, MRIexpB, Stefani07}, as well as nonlinear simulations~\citep{SzklRud06,Szklarski07} indicate that for a flow with relatively small Reynolds number $ \approx 10^3$, and parameters resembling essentially MHD stable flow in the limit of infinitely long cylinders, there exist unexpected time-dependent fluctuation of the velocity field. These disturbances arise as an effect of the vertical boundary conditions, moreover the simulations show that they are much stronger if the end-plates bounding the cylinders are assumed to be perfectly conducting. The plates induce a well known hydrodynamical effect -- the Ekman circulation, which is a result of unbalanced pressure gradients in vicinity of the vertical no-slip boundary conditions. There the Ekman layer develops in which the fluid velocity from the bulk of the container must match the velocity imposed by the end-plates. It seems that for MHD Taylor-Couette flow, magnetic effects, unlike the classical hydrodynamical Ekman layer, induced by the plates have been overlooked. In this paper we argue that the rigidly rotating plates together with an imposed axial magnetic field give rise to a similar layer which develops for an infinite, rotating plate serving as a boundary for the conducting fluid. One of the most important features of such flow is the existence of the Hartmann current (absent in the conventional Hartmann problem,~\citep{2004JFM...504..183K}) which leaves the boundary layer and then interacts with the magnetic field. In particular, this becomes important for conducting plates which was the case for the PROMISE experiment, since one of the end-plates was made from copper. We discuss properties of Ekman-Hartmann layers for infinite, rotating plates and relate it to the end-plates enclosing the cylinders in a Taylor-Couette setup. It is shown that for considered radial boundary conditions, the induced current turn eventually in the radial direction and acting in concert with the imposed axial magnetic field gives rise to a body force. We demonstrate that magnetic effects induced by the end-plates enclosing the cylinders can profoundly alter flow properties. In particular the rotational profile can become significantly different from the expected parabolic Couette solution. Moreover, if the Hartmann current is strong enough, it is likely that the local Rayleigh criterion for stability will be violated and the flow becomes centrifugally unstable. In an MRI experiment it is crucial to rule out such instabilities and a special care concerning the vertical boundary conditions is needed in order to obtain the desired rotational profile. | \citet{GilmanBenton68} have shown with a linear theory that in vicinity of a rotating plane which serves as a border for rotating conducing fluid there develops the Ekman-Hartmann layer if $\Omplain \ne \Omfluid$ and an axial magnetic field is applied. The most important feature of Ekman-Hartmann layers is their ability to induce both mass fluxes and electric currents in the region outside the boundary layer. If $\Omplain < \Omfluid$ these fluxes are directed outwards the layer (``blowing''); when $\Omplain > \Omfluid$ towards the layer (``suction''). For the conducting plates the fluxes are much stronger since additional currents are drawn from/into the plates. Outside the Ekman-Hartmann layer exists the magnetic diffusion region, in which the electric current has only radial components. The current, together with the axial magnetic field, produces an electromagnetic body force acting on the fluid. We have shown in this paper that similar effects arise for the MHD Taylor-Couette flow when the rotating cylinders are bounded by two rigidly rotating end-plates. Near the plates the Ekman-Hartmann layer forms and, consequently, there exists the Hartmann current which penetrates bulk of the fluid. In the presence of an axial magnetic field such problem can be compared with the Taylor-Dean flow -- a flow between, possibly rotating, cylinders which is additionally driven by an azimuthal pressure gradient. We find that under certain conditions the resulting flow becomes unstable, Taylor vortices can be observed and the rotational profile is significantly different from the standard Couette solution $\Omegab$. The instability has essentially a centrifugal character as the Rayleigh criterion is locally vialoted. This is an undesirable effect from the point of view of an MRI experiment. In such experiment it is necessary to obtain a state resembling $\Omegab$ in the major part of the container, for parameters characterizing stable MHD flows. It is necessary to take into account the magnetic effects induced by the plates so that the MRI can be clearly identified rather then any other instability. The fluxes induced in the Ekman-Hartmann layer are a direct consequence of a shear close to the boundaries. Exemplary methods of reducing the shear have been proposed in~\citep{Szklarski07}. For rotation rates characterized by $\Rey$ of order \ord{10^3} all the effects can by significantly reduced by allowing the end-plates to rotate independently of the cylinders~\citep{Abshagenetal04}. Since for $\Omend = \Omin$ there is the Ekman suction, and for $\Omend=\Omout$ the Ekman blowing, there exists $\Omout < \Omend < \Omin$ for which the generated mass and charge fluxes are minimal. Alternatively, one can divide the plates into independently rotating rings~\citep{JiBurin06}. | 7 | 10 | 0710.5600 |
0710 | 0710.0025_arXiv.txt | {Data taken at the Pierre Auger Observatory are used to search for air showers initiated by ultra-high energy (UHE) photons. Results of searches are reported from hybrid observations where events are measured with both fluorescence and array detectors. Additionally, a more stringent test of the photon fluxes predicted with energies above $10^{19}$~eV is made using a larger data set measured using only the surface detectors of the observatory.} \begin{document} | It has been suggested that the excess of cosmic-rays with energies above the predicted Greisen-Zatsepin-Kusmin (GZK) steepening, observed by AGASA, could be due to cosmic-rays being produced as the by-products of 'top-down' scenarios and not by more conventional acceleration mechanisms, such as the diffusive shock process, in astrophysical objects. If the former were the case, it would be expected that a significant proportion of the spectrum of cosmic rays at the highest energies would be UHE photons: predictions vary from model to model but are in the range 10\% to 50\% above 10~$\eev$. The extensive air showers (EAS) generated by UHE photons reach shower maximum \xmax, the atmospheric depth at which the number of charged particles in the shower is greatest, at much greater depths than their nuclear counterparts. This is due to the much lower multiplicity in particle production in the electromagnetic-dominated photon showers than in the hadronic interactions present for nuclear primaries. The depth of maximum is further increased above 30~$\eev$ because of the suppression of Bethe-Heitler pair production due to the LPM effect \cite{LPMEffect_LP}\cite{LPMEffect_M}, which is not important for other cosmic-ray primaries. \xmax~ can thus be used to discriminate between photon and nucleonic UHE primaries. With the Pierre Auger Observatory this can be done using hybrid events which are observed with both fluorescence and array detectors, and in which \xmax~ is measured directly. This allows the determination of a limit to the integral fraction of photons above 10~$\eev$. Parameters measured using only the surface array of detectors, which reflect shower maximum, can also be compared to predictions for photons from simulations to provide a stronger limit on the flux of photons at several energies \cite{SDParametersPhotonLimit}. Two such observables are used which behave differently for nuclear primaries when compared to photons: these are the risetime of the signal in the water-Cherenkov detectors at 1000 m from the core and the radius of curvature of the shower front. | \subsection{Hybrid Results} In \cite{OriginalHybridPhotonLimit} we reported a limit to the fraction of photons in the integral cosmic-ray flux of 16\% (95\% c.l.) above 10~$\eev$ based on 29 high-quality hybrid events recorded in the period Jan. 2004 - Feb. 2006. We have updated the analysis with data collected until March 2007, keeping the analysis cuts as in the original paper. In total, 58 events are now available. The measured \xmax~distribution is shown in figure \ref{HybridXmaxDistribution} along with the calculated distribution from 10~$\eev$ photons made on an event-by-event basis. Even the largest observed value of \xmax, $\sim$~900~$\gcm$, is well below the average value expected for photons (about $\sim$~1000~$\gcm$, see e.g. Table 1 in \cite{OriginalHybridPhotonLimit}). With the updated sample, the upper limit becomes 13\% (95\% c.l.) above 10~$\eev$. \begin{figure}[htbp] \vspace{-10pt} \begin{center} \includegraphics[width=0.45\textwidth]{icrc0602_fig03.eps} \end{center} \caption[font=small,labelfont=bf, aboveship=2pt]{Black: The distribution of \xmax from 58 hybrid events with energies above 10~$\eev$ that meet selected cuts\cite{OriginalHybridPhotonLimit} . The dashed line shows a distribution for 10~$\eev$ photons arriving over a range of zenith angles. All events reach shower maximum at depths which are too shallow to be considered as a result of photon primaries.} \label{HybridXmaxDistribution} \end{figure} \begin{figure}[htbp] \vspace{-0pt} \begin{center} \includegraphics[width=0.45\textwidth, height=0.4\textwidth]{icrc0602_fig04.eps} \end{center} \caption[font=small,labelfont=bf, aboveship=2pt]{The principle component between risetime and radius of curvature as a function of energy for data (black) and simulated photons (red). Data lying above the dashed line, which is the mean of the distribution for photons, are identified as photon candidates. No events meet this requirement.} \label{UCLAPrincipleComponentPlot} \end{figure} \subsection{Surface Detector Results} Photon candidates were searched for by combining the risetime and radius of curvature for each event into one observable using a principle component analysis, where the angle of rotation is found using 5\% of the data. The remaining 95\% of showers were then identified as nuclei or photon candidates using an \emph{a priori} cut, where showers are excluded as photon-like if the principle component measurement is less than the mean of that predicted by a spectrum of photonic showers generated by Monte Carlo simulations. Events which were deemed as photonic were assigned energies using a reconstruction designed from photon simulations, reflecting the fact that these showers have low muon content and are more strongly attenuated than hadronic showers \cite{Billior}. The spectrum of simulated photons incorporates the efficiency of photon detection and reconstructions, as well as including the possibility that photons may interact with the geomagnetic field before arriving in the atmosphere. The results are summarised in figure \ref{UCLAPrincipleComponentPlot} where the principle component is plotted for data (black) and the simulated photons (red) as a function of energy. There are 0 events above the dashed line, where the cut is defined, and as such there are no candidate photons for events above 10 $\eev$. This constrains the maximum flux of photons to $3.8 \times 10^{-3}$,~$2.5 \times 10^{-3}$~and $2.2 \times 10^{-3}$~\fluxunit{} above 10, 20 and 40 $\eev$ respectively. The upper limit on the photon fraction, based on the measured spectrum of events in \cite{SommersSpectrum}, is calculated as 2.0\%, 5.1\% and 31\% at 10, 20 and 40 $\eev$ respectively. These limits are shown against the predictions for photons from top-down models based on the AGASA spectrum (see \cite{review}) in figure \ref{ModelsVsLimits}. Also shown are upper limits to the flux and fractions of photons from previous experiments. This work constrains the photon fractions and fluxes to more stringent limits than previously measured and disfavours the proposed top-down models as the sources of UHECRs and of the AGASA events. On completion of the surface array, the Pierre Auger Observatory will reach sensitivities of $4 \times 10^{-4}$~\fluxunit{} for the integrated flux and 0.7\% for the fraction of photons above 20 $\eev$ (95\% c.l.) after 5 years of operation. \begin{figure}[htbp] \vspace{-120pt} \subfigure[Upper limits to photon fractions] { \centering \includegraphics[width=0.45\textwidth]{icrc0602_fig05.eps} } \subfigure[Upper limits to photon flux] { \centering \includegraphics[width=0.45\textwidth]{icrc0602_fig06.eps} } \caption[font=small,labelfont=bf, aboveship=2pt]{(a) The upper limits to the fraction of photons including limits from previous experiments and predictions from top-down models based on the AGASA spectrum. The limits from this work are shown in black. The limit derived from the hybrid data is also shown and labelled FD. Labels A, HP and Y refer to the limits set by AGASA, Haverah Park and Yakutsk, for references see \cite{review}. (b) The upper limits to the flux of photons along with predictions of top-down models.} \label{ModelsVsLimits} \end{figure} | 7 | 10 | 0710.0025 |
0710 | 0710.2399_arXiv.txt | Using the IRAM 30 m telescope, a mapping survey in optically thick and thin lines was performed towards 46 high mass star-forming regions. The sample includes UC~H{\sc ii} precursors and UC~H{\sc ii} regions. Seventeen sources are found to show "blue profiles", the expected signature of collapsing cores. The excess of sources with blue over red profiles ([$N_{\rm blue}$ -- $N_{\rm red}$]/$N_{\rm total}$) is 29\% in the HCO$^+$ $J$=1--0 line, with a probability of 0.6\% that this is caused by random fluctuations. UC~H{\sc ii} regions show a higher excess (58\%) than UC~H{\sc ii} precursors (17\%), indicating that material is still accreted after the onset of the UC~H{\sc ii} phase. Similar differences in the excess of blue profiles as a function of evolutionary state are not observed in low mass star-forming regions. Thus, if confirmed for high mass star-forming sites, this would point at a fundamental difference between low- and high-mass star formation. Possible explanations are inadequate thermalization, stronger influence of outflows in massive early cores, larger gas reserves around massive stellar objects or different trigger mechanisms between low- and high- mass star formation. | Inflow motion is a fundamental phenomenon during stellar formation. Although the search for inflow is usually more difficult than that for outflow, studies of inflow have made great progress since the 1990s. In low-mass star forming regions, inflow motions were detected at different evolutionary stages, including Class --I, Class 0 and Class I cores \citep{zho93, mmt97, lee99, gre00, evans03}. Recently, a number of inflow candidates were found in high mass star formation regions. Among a sample of 28 massive cores, 12 were found to show line profiles that peak at blue-shifted velocities (hereafter "blue profiles"; see Sect.\,3.1), the expected signature of inflow \citep{we03}. \citet{fws05} (hereafter FWS05) detected such asymmetric profiles in 22 cores within a sample of 77 high-mass proto-stellar objects (HMPOs). Most recently, \citet{wyr06} detected 9 sources with a blue profile in a sample of 12 ultracompact (UC)~H{\sc ii} regions. Variation of inflow motion with time is critical for high mass star formation. It has been indicated that when a protostar reaches $>$10 M$_{\odot}$ it can generate enough radiation pressure to halt spherical infall and inhibit its mass increasing\citep{wc87}. Observationally, however, it is not yet clear how inflow is related to the evolution of massive (proto)stars. To study this problem, we have carried out a survey for a sample including both cores of UC~H{\sc ii} regions and precursors of UC~H{\sc ii} regions. While previous surveys using single point observations provided some statistical evidence for the occurrence of infall within massive cores, blue profiles can also be caused by rotation. Therefore maps of the molecular environment are indispensable. Mapping also allows us to locate the center of the inflow and to identify cores that are simultaneously showing evidence for in- and outflow. Therefore, we conducted a mapping survey including 46 high mass star-forming regions which were selected applying three criteria: (1) The sources must have been mapped in the submillimeter or millimeter wavelengths with continuum or spectroscopy; (2) signal-to-noise ratios should be $>$5 at 350\,$\mu$m (Mueller et al. 2002) and higher at other wavelengths; (3) there should be no other core within one arcmin \citep{zhs97, hnb98, tie98, hat00, mbc00, bsp02, mse02}. With respect to their stellar content, we can divide the sample into two different groups of targets: (1) Thirty three sources lack 6 cm continuum emission and are precursors of UC~H{\sc ii} regions or HMPOs \citep{mbc00, bsp02}. Among these, thirty are hosting a luminous IRAS source. The remaining three are associated with IRAC (the InfraRed Array Camera on the Spitzer Space Telescope) point sources (W3-W and W3-SE) or are not hosting an IRAC source (18454--3). All 33 cores comprise `group~I'. (2) Thirteen UC~H{\sc ii} regions are assigned to `group~II'. This letter presents a list of the identified collapse candidates and provides the statistics of blue excesses. Detailed properties of individual cores will be analyzed in a future paper. | \subsection{Blue profile identification} For self-absorbed optically thick lines, the classical signature of inflow is a double peaked profile with the blue-shifted peak being stronger, or a line asymmetry with the peak skewed to the blue side. While optically thin lines should show a single velocity component peaking at the line center. Among the 46 cores observed, five (05490+2658, G31.41+0.31, 18454-3, 18454-4, 19266+1745) will be ignored because they show either too complex spectral profiles, inhibiting a detailed analysis, or a lack of optically thin lines. Estimates of optical depths were obtained from line ratios between different isotopomers of CO and CS and from the relative intensities of individual hyperfine components in the case of C$^{17}$O and N$_2$H$^+$. C$^{18}$O, C$^{17}$O, C$^{34}$S and N$_2$H$^+$ tend to be optically thin, while CS is optically thick. HCO$^+$ opacities could not be estimated. However, the similarity of HCO$^+$ and CS line shapes (see Sect.\,3.2) as well as the results of \cite{gre00} and FWS05 clearly indicate that HCO$^+$ is also optically thick. The 41 remaining sources were detected in at least one optically thick and one optically thin line. A blue profile caused by inflow motion with velocity $v \propto r^{-1/2}$ in a region with higher excitation temperature ($T_{ex}$) inside requires $T_{\rm A}$*(B)/$T_{\rm A}$*(R) $>$ 1. Here $r$ is the radius of the collapsing core \citep{zho93}. $T_{\rm A}$*(B) and $T_{\rm A}$*(R) are the blue and red peak intensities of the optically thick line. We also define a dimensionless asymmetry parameter following \cite{mmt97}, $\delta V$ = ($V_{\rm thick}$-$V_{\rm thin}$)/$\Delta$$V_{\rm thin}$. $V_{\rm thick}$ is the peak velocity of the opaque line, $V_{\rm thin}$ and $\Delta$$V_{\rm thin}$ denote the peak velocity and width of the optically thin line. Only for $\delta V < -0.25$ or $> 0.25$ the line profile is rated blue or red, respectively. Our sources (Table~2) discriminate among five main types of line shapes: (1) cores with lines showing a ``blue profile'' (in the following denoted with B); (2) cores with lines showing a "red profile" (R); (3) cores exhibiting blue and red profiles at different spatial positions (BRS); (4) cores where some lines show a blue profile, while others display a red profile (BRL); (5) cores without obvious asymmetric lines (S). Only cores showing at least one line of type B, but no lines of type R are identified as targets potentially undergoing inflow motion. \subsection{Collapse candidates and their profile ``excess''} \label{bozomath} With the criteria outlined in Sect.\,3.1, seventeen inflow candidates are identified (see Table~2). Ten belong to group~I and seven are part of group~II. To provide a typical example, Fig.\,1 shows the infall signature of the group~I core W3-SE. Fig. 1a displays the HCO$^+$\,(1--0) spectra, showing the angular size of the core. Fig. 1b shows a number of profiles towards the central position. The HCO$^+$\,(1--0) and (3--2) lines as well as the CS\,(3--2) transition show the blue asymmetry. For the HCO$^+$\,(1--0) line this is also demonstrated in the position-velocity (P-V) diagram of Fig.\,1c. For comparison, Fig.\,1d shows a P-V diagram of the optically thin C$^{18}$O\,(1--0) emission. The quantity ``excess'' as defined by \cite{mmt97} is $E$ = ($N_{\rm B} - N_{\rm R}$)/$N_{\rm T}$, where $N_{\rm B}$ and $N_{\rm R}$ mark the numbers of sources with blue and red profiles. $N_{\rm T}$ is the total number of sources. For our survey the excess was calculated for the two HCO$^+$ transitions and the CS\,(3--2) line. Fig. 2 shows the log[$T_{\rm A}$*(B)/$T_{\rm A}$*(R)] and $\delta$V (see Sect.\,3.1) distributions of the three individual lines. Statistical results are given in Table~3. The observed excess derived from the HCO$^+$\,(1--0) and (3--2) lines is 0.29 and 0.11, respectively. Both are larger than those obtained by FWS05 for the same lines (0.15 and 0.04). For the CS transition we obtain 0.29. To evaluate the statistical significance of the determined values, we conducted the binomial test (see FWS05 and references therein). Probabilities that the excesses are a product of a random distribution are given in the last column of Table~3. These are 0.006 and 0.01 for HCO$^+$\,(1--0) and CS (3--2) respectively. Apparently, both lines are sensitive tracers of potential inflow motion in massive cores. To evaluate differences between the two classes of cores (I and II; see Sect.\,1) with respect to the excess, we used the HCO$^+$\,(1--0) line, which was mapped in the largest number of sources. The results listed in the lower part of Table~3 include 16 sources with profiles of type B. The excesses observed for group~I and II are 0.17 and 0.58, respectively. Twenty of our 46 sources overlap with those of FWS05. Among them are 19 group~I sources (out of 33), but only one source is from group~II (out of 13). Our study includes various CO and CS lines. We also made maps. Thus we can view the common objects from a different perspective and can check, how far the choice of different molecular transitions and the presence of maps is leading to contradictions with previously published results. Differences are indeed significant. For eight of the 19 overlapping type I cores we obtain different line asymmetry classification, emphasizing the need for detailed maps. Nevertheless, the overall difference in the HCO$^+$\,(1--0) excess is negligible (0.17 versus 0.15). To summarize, both data sets indicate that the HCO$^+$\,(1--0) excess is low for UC~H{\sc ii} precursors. For UC~H{\sc ii} regions, our results and those of \cite{wyr06} suggest that the excess is larger and more significant. From the binomial test for group~I and II, the probability that the blue excesses (0.17 and 0.58 respectively) are arising by chance is 0.13 and 0.008, respectively. \subsection{A comparison with low mass star-forming surveys} While low mass star-forming regions show infall from the Class --I to the Class I stages of evolution, high mass star-forming regions also exhibit infall signatures from their earliest stages till a UC H{\sc ii} region has formed \citep{wzx05,bkl06,qzm06}. In low mass cores the profile excess was found to be 0.30, 0.31 and 0.31 for Class --I, 0 and I core samples in the HCO$^+$\,(3--2) line (\cite{evans03} and references therein). There seem to be no significant differences among the cores in different evolutionary phases. However, our samples show the excess of UC~H{\sc ii} regions far surpassing that of the UC~H{\sc ii} precursors. This may point to fundamental differences between low and high mass star-forming conditions. Possible causes to the higher blue excess in Group II sources may be: (1) The molecular gas surrounding UC H{\sc ii} regions may be more adequately thermalized to show the blue excess, i.e. the excitation temperature of specific lines may increase more monotonically towards the center. Thus all lines may produce blue profiles indicating infall motion, while in younger cores still some lines may show red profiles. (2) The amount of dense cool gas is larger towards younger objects. Outflows of dense molecular gas may be more active around Group I objects, shaping more red profiles. (3) Low mass cores are relatively isolated and their gas supply is limited. Simulations showed that this may halt the increase of inflow \citep{vor05}. However, high mass stars form in giant molecular clouds and their inflow motions are not easily halted by the exhaustion of molecular gas before most of it is dispelled. (4) In low mass cores, star formation may be spontaneous. In high mass cores, collapse may be trigged by extrinsic disturbances and the collapse may take more time to develop. With respect to potential selection effects, we used the same criteria to identify the targets of the two separate groups of sources. Since this study is based on a limited number of sources, more data quantifying the blue excess as a function of evolutionary stage would be highly desirable. | 7 | 10 | 0710.2399 |
0710 | 0710.2350_arXiv.txt | {Gamma-ray bursts (GRBs) have been detected up to GeV energies and are predicted by many models to emit in the very high energy (VHE, $>$ 100 GeV) regime too. Detection of such emission would allow us to constrain GRB models. Since its launch, in late 2004, the Swift satellite has been locating GRBs at a rate of approximately 100 per year. The rapid localization and follow-up in many wavelengths has revealed new and unexpected phenomena, such as delayed emission in the form of bright X-ray flares. The Milagro gamma-ray observatory is a wide field of view (2 sr) instrument employing a water Cherenkov detector to continuously ($>$ 90\% duty cycle) observe the overhead sky in the 100 GeV to 100 TeV energy range. Over 100 GRBs are known to have been in the field of view of Milagro since January 2000, including 57 since the launch of Swift (through May 2007). We discuss the results of the searches for prompt emission from these bursts, as well as for delayed emission from the X-ray flares observed in some of the Swift bursts.} \begin{document} | Almost 40 years after the detection of the first gamma-ray burst (GRB), many questions remain about these powerful explosions. Progress in the field has accelerated in the decade since the detection of the first X-ray afterglow~\cite{costa}. The launch of Swift~\cite{gehrels}, in late 2004, has resulted in the rapid and accurate localization of $\sim$100 GRBs per year. This has enabled an unprecedented number of multiwavelength follow-up observations in every band of the electromagnetic spectrum, from radio to the highest energy gamma rays, paving the way for a more complete understanding of these phenomena. The highest energy emission to be detected conclusively from GRBs was seen by the EGRET instrument. EGRET detected several bursts above 100 MeV, with no evidence for a cut-off in the spectrum~\cite{dingus01}. One burst in particular, GRB 940217, emitted an 18 GeV photon over 90 minutes after the start of the burst~\cite{hurley94}. Another burst, GRB 941017, was found to display a second, higher energy, spectral component which extended up to at least 200 MeV and decayed more slowly than the lower energy component~\cite{gonzalez03}. On this evidence alone, the expectations are that the upcoming launch of GLAST, with its much higher sensitivity than EGRET, will result in a large number of detections at high energies, leading to a better understanding of the GeV properties of GRBs. At very high energy (VHE, $>$100 GeV), there have been no conclusive detections of GRBs, although there have been several tantalizing hints of emission. Milagrito, a prototype of Milagro, searched for emission coincident with 54 BATSE bursts and reported evidence for emission above 650 GeV from GRB 970417a, at the 3$\sigma$ level~\cite{atkins00a,atkins03}. The HEGRA group reported evidence at the 3 sigma level for emission above 20 TeV from GRB 920925c~\cite{padilla98}. Follow-up observations above 250 GeV by the Whipple atmospheric Cherenkov telescope~\cite{connaughton97} failed to find any high energy afterglow for 9 bursts studied, though the delay in slewing to observe these bursts ranged from 2 minutes to almost an hour. More recent efforts by the MAGIC~\cite{magicGRBlimits}, Whipple~\cite{whippleGRBlimits}, and Milagro~\cite{milagro1} telescopes have resulted only in upper limits. There is no shortage of models predicting (or explaining) very high energy emission from GRBs (e.g.~\cite{dermer00,pilla98,zhang01,granot03,gupta07}). Many of the proposed emission mechanisms predict a fluence at TeV energies comparable to that at keV-MeV energies. One such mechanism involves the inverse Compton upscattering of lower energy (synchrotron) photons by the energetic electrons which emitted them. The strong magnetic fields and large bulk Lorentz factors present in GRB jets can result in a high energy component peaked at TeV energies. The strength of such a component is highly dependent on the environments of the particle acceleration and the gamma ray production. Milagro\cite{atkins00b,atkins01} is a TeV gamma-ray detector, located at an altitude of 2630 m in northern New Mexico. Milagro uses the water Cherenkov technique to detect extensive air-showers produced by VHE gamma rays as they interact with the Earth's atmosphere. The Milagro field of view is $\sim$2 sr and duty cycle is $>$90\%. The effective area is a function of zenith angle and ranges from $\sim50$ m$^2$ at 100 GeV to $\sim10^5$ m$^2$ at 10 TeV. A sparse array of 175 4000-l water tanks, each with a PMT, was added in 2002. These ``outriggers,'' extend the physical area of Milagro to 40000 m$^2$. The combination of large field of view and high duty cycle make Milagro the best instrument currently available for conducting a search for prompt VHE emission from GRBs. Milagro is also able to operate in ``scaler'' mode, where individual tube rates can be monitored to search for GRB emission in the 1--100 GeV energy range coincident in time with known satellite bursts (see ~\cite{taylor,cesar}). In addition to conducting a search for emission from satellite-detected GRBs, Milagro is capable of performing a blind search for emission at all locations in its field of view and many different durations (see ~\cite{vlasios}). \subsection{The GRB Sample} Twenty-five satellite-triggered GRBs occurred within the field of view of Milagro between January 2000 and December 2001. No significant emission was detected from any of these bursts~\cite{atkins05}. In the period from January 2002 to December 2004 (post-BATSE and pre-Swift), there were only 11 well-localized GRBs within the Milagro field of view. Since the launch of Swift (through the end of May 2007), however, there have been a total of 57 bursts in the field of view of Milagro~\footnote{A good source of information on well-localized GRBs is J. Greiner's web page {\tt http://www.mpe.mpg.de/$\sim$jcg/grbgen.html}}, many of them with measured redshift. Table~\ref{grb_table} lists the 57 GRBs in this sample, along with some of their properties. Due to the absorption of high-energy gamma rays by the extragalactic background light (EBL), detections at VHE energies are only expected for redshifts less than $\sim$0.5. The degree of gamma-ray extinction from this effect is uncertain, because the amount of EBL is not well known. In this work we use the model of \cite{primack05}, which predicts an optical depth of roughly unity for 500 GeV (10 TeV) gamma rays from a redshift of 0.2 (0.05). \subsection{Bright X-ray Flares} One of the highlights of the Swift mission has been the detection of bright X-ray flares in the afterglow phase of many GRBs~\cite{burrows05}. These flares can sometimes be as bright as the prompt component of the GRB itself. While the exact nature of the flares is not clear, with various different proposals for what may be causing them (e.g. \cite{guetta07,kobayashi07,krimm07}), there are reasonable expectations that these flares will result in the emission of delayed very high energy photons~\cite{wang06}. The first survey of X-ray flares from GRBs observed by Swift in the first year of operations~\cite{chincarini} shows that approximately one third of GRBs show significant flaring activity. Table~\ref{flare_table} lists the 10 flares from the flare catalog of ~\cite{chincarini,falcone} which were in the field of view of Milagro. | 7 | 10 | 0710.2350 |
|
0710 | 0710.0355_arXiv.txt | An analytical approximation to periodic orbits in the circular restricted three-body problem is provided. The formulation given in this work is based in calculations known from classical mechanics, but with the addition of the necessary terms to give a fairly good approximation that we compare with simulations, resulting in a simple set of analytical expressions that solve periodic orbits on discs of binary systems without the need of solving the motion equations by numerical integrations. | \label{sec:introd} The majority of low-mass main-sequence stars seem to be grouped in multiple systems preferentially binary (Duquennoy \& Mayor 1991, Fischer \& Marcy 1992). These systems have attracted more attention since the discovery that many T-Tauri and other pre-main-sequence binary stars possess both circumstellar and circumbinary discs from observations of excess radiation at infrared to millimeter waves and direct images in radio (Rodr{\'i}guez \etal 1998; for a review see Mathieu 1994, 2000). The improvement in the observational techniques allows, on the other hand, the testing of theories on these objects. Moreover, extrasolar planets have been found to orbit stars with a stellar companion, e.g., 16 Cygni B, $\tau$ Bootis, and 55 $\rho$ Cancri (Butler \etal 1997; Cochran et al. 1997). All these facts make the study of stellar discs in binary systems, as well as the possibility of stable orbits on them that could be populated by gas or particles, a key element for better understanding stellar and planetary formation. In particular the study of simple periodic orbits of a test particle in the restricted three-body problem results a very good approximation to the streamlines in an accretion disc in a very small pressure and viscosity regime, or for proto-planetary systems and its debris (Kuiper belt objects). Extensive theoretical work has been done in this direction (Lubow \& Shu 1975; Paczy\'nski 1977; Papaloizou \& Pringle 1977; Rudak \& Paczy\'nski 1981; Bonnell \& Bastien 1992, Bate 1997; Bate \& Bonnell 1997). Several studies are directed to find the most simple and important geometric characteristic of the discs which is the size of both circumstellar and circumbinary discs, going from analytical approximations: Eggleton (1983), who provides a simple analytical approximation to the Roche lobes; Holman \& Wiegert (1999) radii in planetary discs; Pichardo, Sparke \& Aguilar (2005, PSA05 hereafter) radii in eccentric binary stars. Based in the approximation studied in classical mechanics to periodic orbits given by Moulton (1963), we provide in this work a fairly good approximation by calculating the necessary terms, for the case of circular binaries, to periodic orbits. This formulation gives not only the radii of the circumstellar and circumbinary discs, but gives a good description of any periodic orbit at a given radius for any of the discs, in the form of an analytical approximation. In this manner, the set of equations we provide can be used to find any periodic orbit or its initial conditions to run it directly in simulations of the three-body restricted problem, or as a good approximation to initial conditions in the eccentric case as well, from the point of view of particle or hydrodynamical simulations, without the need of solving the restricted three body problem numerically. The low viscosity regime described by the periodic orbits representation, is important in astrophysics since it gives an idea of how bodies would respond only influenced by the effects of the potential exerted by the binaries, and they could permit to link the periodic orbits to physical unknown characteristics of the discs like viscosity that needs to be artificially introduced to hydrodynamical codes. It also might permit the calculation of physical characteristics like dissipation rates, etc. Studies like these result easier given in the form of approximated analytical expressions, instead of the usual numerical approximations to the periodic orbits since an analytical formulation is much faster to add to any hydrodynamical or particle code that requires the characteristics of a given family of periodic orbits, or simply their initial conditions in a more precise form than assuming Keplerian discs until the approximation fails. The periodic orbits show regions where the orbits are compressed (or decompressed), these regions could trigger density fluctuations maybe able to drive material to form important agglomerations in the discs that, depending on their positions on the disc and the density, could give origin to planets. An analytical approximation with a given density law for the discs would allow these kind of studies to be much faster and easier. In Section \ref{sec:method}, we show the strategy and equations used to find the approximation to the periodic orbits in circumstellar discs. In Section \ref{sec:orbits2}, we present the equations approximating the periodic orbits in circumbinary discs, and an analysis of stability in Section \ref{sec:stability}. A comparison of the application of the formulation and numerically calculated periodic orbits is given in Section \ref{sec:comparison}. Conclusions are presented in Section \ref{sec:conclusions}. | \label{sec:conclusions} We have constructed an analytical set of equations based on perturbative analysis that result simple and precise to approximate the solution for the circular restricted three body problem. We choose some terms in the expanded equations of motion, which are relevant to the solution of the problem. The original equations are approximated with a set of equations that, in many cases, have an analytical solution. We show the goodness of our analytical approximation by direct comparison of geometry and rotation curves with the full numerical solution. The set of equations we prsent can provide any periodic orbit for a binary system, either circumstellar or circumbinary, and not only the outer edge radius of circumstellar discs or the inner edge radius (gap) of the circumbinary disc, without the need to solve the problem numerically. The relations provided are simple and straightforward in such a way that our approximation can be used for any application where initial conditions of periodic orbits or complete periodic orbits are needed, or for direct study of the three body restricted circular problem. For instance, in order to introduce on a hydrodynamical or particle code, our initial conditions would result in much more stable discs than with keplerian initial conditions, or than by constructing the discs directly from hydrodynamical simulations of accretion to binaries that will require long times to obtain stable discs. Since our approximation is completely analytical it will have the obvious additional advantage of being computationally, extremely cheap, and easier to implement than the numerical solution. The periodic orbits respond uniquely to the binary potential and do not consider other physical factors such as gas pressure, or viscosity that work to build the fine details of discs structure. They are however, the backbone of any potential and their shape and behavior give the general discs phase space structure. In this manner, apart of all the possible hydrodynamical or particle discs applications, we can directly use them to study from the general discs geometry to rotation curves, or rarification and compression zones by orbital crowding, etc. | 7 | 10 | 0710.0355 |
0710 | 0710.5093_arXiv.txt | The Unified Model of AGN predicts the sole difference between Seyfert 1 and Seyfert 2 nuclei is the viewing angle with respect to an obscuring structure around the nucleus. High energy photons above 20 keV are not affected by this absorption if the column is Compton thin, so their 30--100~keV spectra should be the same. However, the observed spectra at high energies appear to show a systematic difference, with Seyfert 1's having $\Gamma \sim $2.1 whereas Seyfert 2's are harder with $\Gamma \sim$ 1.9. We estimate the mass and accretion rate of Seyferts detected in these high energy samples and show that they span a wide range in $L/L_{Edd}$. Both black hole binary systems and AGN show a correlation between spectral softness and Eddington fraction, so these samples are probably heterogeneous, spanning a range of intrinsic spectral indices which are hidden in individual objects by poor signal-to-noise. However, the mean Eddington fraction for the Seyfert 1's is higher than for the Seyfert 2's, so the samples are consistent with this being the origin of the softer spectra seen in Seyfert 1's. We stress that high energy spectra alone are not necessarily a clean test of Unification schemes, but that the intrinsic nuclear properties should also change with $L/L_{Edd}$. | The simplest version of the Unification model of Antonucci \& Miller (1985) is that the central engine (black hole, its accretion disc and broad line region: BLR) are the same in all AGN, but that this is embedded in an obscuring torus. The nucleus is seen directly only for inclination angles which do not intersect this material, giving the classic Seyfert 1 AGN signature of a strong and variable UV/X-ray continuum and broad emission lines. Conversely, where the obscuration is in the line of sight, these features are hidden, and the presence of an AGN can only be inferred from the high excitation lines produced in the narrow line region on much larger scales (Seyfert 2's). A key piece of evidence for this scenario is the detection of polarized broad emission lines in classic Seyfert 2 galaxies (most notably in NGC 1068), showing that the BLR is present, but can only be seen via scattered light (e.g. Antonucci \& Miller 1985; Tran 1995, Heisler et al. 1997). This scattering medium filling the 'hole' in the torus is also detected in transmission in Seyfert 1's through the partially ionised absorption signatures seen in soft X-rays (e.g. Blustin et al 2005). \begin{figure*} \begin{center} \begin{tabular}{l} \leavevmode \epsfxsize=12cm \epsfbox{panel.eps} \end{tabular} \end{center} \caption{Plots showing the distribution of log mass and log accretion rate for individual instruments and the full sample (bottom panels) with Seyfert 1 sources shown by a black dashed line and Seyfert 2 sources by a solid cyan line. The log mean accretion rate for each type is given in the top right hand corner.} \label{} \end{figure*} The intrinsic X-ray spectrum of a type 1 AGN can be well described by a power law (of photon spectrum $N(E)\propto E^{-\Gamma}$) together with its Compton reflection from the optically thick accretion disc and/or torus (e.g. Nandra \& Pounds 1994). Again, the Unified models are supported by X-ray observations, which show that this intrinsic spectrum is substantially suppressed at low energies due to absorption in Seyfert 2's (Awaki et al 1991; Smith \& Done 1996; Turner et al 1997; Bassani et al 1999; Cappi et al 2006). However, the strong energy dependence of photoelectric absorption means that this is unlikely to affect samples above 10 keV for columns which are not Compton thick, yet high energy experiments show that the intrinsic spectra of Seyfert 2 galaxies are systematically harder ($\Gamma$ $\sim$ 1.9 - 2) than than the Seyfert 1's, which have $\Gamma > $ 2 (Zdiarski et al. 1995, Gondek et al. 1996, Perola et al. 2002, Malizia et al. 2003, Beckmann et al. 2006). This is not consistent with the idea that these nuclei are identical, unless the intrinsic emission is anisotropic, being harder in the equatorial plane. However, all accreting black hole systems (both stellar and supermassive) generally show softer intrinsic spectra at higher accretion rates (hereafter parameterised as Eddington fraction, $L/L_{Edd}$), e.g. Laor (2000); Remillard \& McClintock (2006). Thus inclination alone is not the sole determinant of the observed spectrum, and Seyferts at different $L/L_{Edd}$ should not be expected to have the same intrinsic emission. Here we collate estimates of black hole mass and mass accretion rate for the Seyfert 1 and 2's detected out to $>50$~keV from CGRO (OSSE), BeppoSAX (PDS) and INTEGRAL (IBIS) instruments. We show that these high energy samples of Seyferts span a large range $L/L_{Edd}$, so should include a range of intrinisic spectral slopes. The mean $L/L_{Edd}$ for the Seyfert 1's in the sample is higher than that for Seyfert 2's, consistent with the steeper spectral slope inferred for the Seyfert 1's being due to the intrinsic spectrum softening at higher $L/L_{Edd}$. We stress that samples of Seyfert 1 and 2 AGN need to be matched on intrinsic properties such as $L/L_{Edd}$ in order to explore differences in orientation. \begin{table*} \begin{center} \begin{minipage}{125mm} \bigskip \begin{tabular}{ll|l|c|c} \hline Source & Instrument & logM/M$_{\odot}$ & L/L$_{Edd}$ &Reference\\ \hline \vspace{1mm}MCG--6-30-15 & {\em OSSE }, {\em BeppoSAX}, {\em INTEGRAL} & 6.81 $^{+0.54}_{-0.54}$ $^{\sigma}$ & 0.25 $^{x}$ & 6 (18) + 23\\ \vspace{1mm}IC 4329A & {\em OSSE }& 6.70 $^{+0.56}_{-peg}$ (mean) $^{r}$ & 0.79 $^{f}$ & 20+ 20\\ \vspace{1mm}&&6.85 $^{+0.55}_{-peg}$ (rms) $^{r}$ &&\\ \vspace{1mm}MR 2251-178& {\em INTEGRAL}& 6.92 $^{+0.22}_{-0.22}$ $^{m}$ & 4.68 $^{z}$ & 10 + 21 \\ \vspace{1mm}NGC3783 & {\em OSSE }, {\em BeppoSAX}& 6.97 $^{+0.30}_{-0.97}$ (mean) $^{r}$ & 0.23 $^{f}$ & 2 + 20 \\ \vspace{1mm}&& 7.04 $^{+0.30}_{-0.96}$ (rms) $^{r}$ &&\\ \vspace{1mm}NGC7469 & {\em OSSE }, {\em BeppoSAX} & 6.81 $^{+0.30}_{-peg}$ (mean) $^{r}$ & 2.14 $^{f}$ & 20 + 20\\ \vspace{1mm}&& 6.88 $^{+0.29}_{-peg}$ (rms) $^{r}$ &&\\ \vspace{1mm}ESO 141-55 & {\em OSSE }, {\em BeppoSAX} & 7.10 $^{+0.57}_{-0.57}$ $^{m}$ & 2.51 $^{ox}$ & 1\\ \vspace{1mm}MCG--2-58-22 & {\em OSSE } & 7.14 $^{+0.57}_{-0.57}$ $^{m}$ & 3.47 $^{ox}$ & 1\\ \vspace{1mm}NGC6814 & {\em OSSE } & 7.08 $^{+0.57}_{-0.57}$ $^{r}$ & 0.05 $^{f}$ & 22 + 2\\ \vspace{1mm}3C120 & {\em BeppoSAX} & 7.36 $^{+0.22}_{-0.28} $ (mean)$^{r}$ & 0.65 $^{f}$ & 20 + 20 \\ \vspace{1mm}&&7.48 $^{+0.21}_{-0.27}$ (rms) $^{r}$ &&\\ \vspace{1mm}NGC3516 & {\em OSSE } & 7.36 $^{+0.59}_{-0.59}$ $^{r}$ & 0.07 $^{f}$ & 2\\ \vspace{1mm}Mrk279 & {\em OSSE } & 7.54 $^{+0.10}_{-0.13}$ $^{r}$ & 0.13 $^{z}$ & 5\\ \vspace{1mm}NGC3227 & {\em OSSE } & 7.59 $^{+0.19}_{-peg}$ (mean) $^{r}$ & 0.01 $^{f}$ & 20 + 2\\ \vspace{1mm}&&7.69 $^{+0.18}_{-peg}$ (rms) $^{r}$ &&\\ \vspace{1mm}NGC5548 & {\em OSSE }, {\em BeppoSAX} & 8.09 $^{+0.09}_{-0.07}$ (mean) $^{r}$ & 0.05 $^{f}$ & 20 + 20\\ \vspace{1mm}&& 7.97 $^{+0.08}_{-0.07}$ (rms) $^{r}$ &&\\ \vspace{1mm}1H 1934-063 & {\em INTEGRAL} & 7.86 $^{+0.63}_{-0.63}$ $^{m}$ & 0.07 $^{ox}$ & 1\\ \vspace{1mm}Fairall 9 & {\em BeppoSAX} & 7.90 $^{+0.11}_{-0.31}$ (mean) $^{r}$ & 0.16 $^{s}$ & 20 + 20\\ \vspace{1mm}&&7.91 $^{+0.11}_{-0.32}$ (rms) $^{r}$ &&\\ \vspace{1mm}NGC7213 & {\em OSSE } & 7.99 $^{+0.64}_{-0.64}$ $^{\sigma}$ & 0.02 $^{f}$ & 2\\ \vspace{1mm}MCG +8-11-11 & {\em OSSE } & 8.06 $^{+0.64}_{-0.64}$ $^{m}$ & 0.35 $^{ox}$ & 1\\ \vspace{1mm}NGC526A & {\em OSSE }, {\em BeppoSAX} & 8.11 $^{+0.65}_{-0.65}$ $^{\sigma}$ & 0.17 $^{ox}$ & 6 (18) + 1\\ \vspace{1mm}Mrk841 & {\em OSSE } & 8.49 $^{+0.68}_{-0.68}$ $^{r}$ & 0.17 $^{f}$ & 22 + 2\\ \vspace{1mm}Mrk509 & {\em OSSE }, {\em BeppoSAX}, {\em INTEGRAL} & 8.16 $^{+0.04}_{-0.04}$ $^{r}$ & 0.45 $^{ox}$ & 5 + 1\\ \vspace{1mm}III Zw 2 & {\em OSSE } & 8.42 $^{+0.67}_{-0.67}$ $^{r}$ & 0.13 $^{z}$ & 3\\ \vspace{1mm}PG 1416-129 & {\em INTEGRAL} & 8.75 $^{+0.70}_{-0.70}$ $^{r}$ & 0.12 $^{z}$ & 3\\ \vspace{1mm}3C111 & {\em INTEGRAL} & 9.56 $^{+0.76}_{-0.76}$ $^{m}$& 0.01 $^{z}$ & 9\\ \hline \end{tabular} \end{minipage} \end{center} \caption{Seyfert 1 sub-sample. Name superscripts indicate the hard X-ray satellite that observed the source, o: {\it OSSE}, i: \inte ~\& b: \sax. Mass superscripts refer to the measurement method, $\sigma$: stellar velocity dispersion, r: reverberation mapping, m: other. Accretion rate superscripts denote the method by which the bolometric luminosity was determined, x: correction from X-ray, f: integration of the SED flux, ox: correction from O{\small III}, s: SED modelling, z: other. References, 1: Wang \etal (2007), 2: Woo \& Urry (2002), 3: Hao \etal (2005), 4: McHardy (1988), 5: Vestergaard \& Peterson. (2006), 6: Garcia-Rissmann \etal (2005), 7: Bian \& Gu (2007), 8: Greenhill \etal (1997), 9: Grandi \etal (2006), 10: Morales \& Fabian (2002), 11: Marconi \etal (2006), 12: Whysong \& Antonucci (2004),13: Awaki \etal (2005), 14: Czerny \etal (2001), 15: Tadhunter \etal (2003), 16: van Bemmel \etal (2003), 17: Gu \etal (2006), 18: Tremaine \etal (2002), 19: Risaliti \etal (2005), 20: Kaspi \etal (2000), 21: Monier \etal (2001), 22: Laor \etal (2001) 23: Dadina (2007). Where more than one reference is given, the first refers to the black hole mass and the second to the accretion rate. Where reference 18 is given in brackets, the relation from Tremaine \etal (2002) has been used to determine the black hole mass via stellar velocity dispersion.} \end{table*} \begin{table*} \begin{center} \begin{minipage}{110mm} \bigskip \begin{tabular}{ll|l|c|c} \hline Source & Instrument & logM/M$_{\odot}$ & L/L$_{Edd}$ & Reference\\ \hline \vspace{1mm}NGC6300 & {\em INTEGRAL} & 5.45 $^{+0.44}_{-0.44}$ $^{m}$ & 0.91 $^{x}$ & 13\\ \vspace{1mm}NGC7314 & {\em BeppoSAX} & 6.03 $^{+0.48}_{-0.48}$ $^{\sigma}$ & 0.41 $^{x}$ & 17 (+18)\\ \vspace{1mm}NGC4945 & {\em OSSE }, {\em INTEGRAL} & 6.15 $^{+0.49}_{-0.49}$ $^{\sigma}$ & 0.17 $^{x}$ & 8\\ \vspace{1mm}MCG--5-23-16 & {\em OSSE }, {\em BeppoSAX} & 6.29 $^{+0.50}_{-0.50}$ $^{m}$ & 0.09 $^{ox}$ & 1\\ \vspace{1mm}Circinus & {\em INTEGRAL} & 6.42 $^{+0.51}_{-0.51}$ $^{\sigma}$ & 0.32 $^{ox}$ & 7\\ \vspace{1mm}NGC5506 & {\em OSSE } & 6.65 $^{+0.53}_{-0.53}$ $^{\sigma}$ & 2.51 $^{ox}$ & 7\\ \vspace{1mm}NGC4593 & {\em INTEGRAL}, {\em BeppoSAX} & 6.91 $^{+0.55}_{-0.55}$ $^{r}$ & 0.12 $^{f}$ &2\\ \vspace{1mm}ESO 103-G35 & {\em INTEGRAL}, {\em BeppoSAX} & 7.14 $^{+0.57}_{-0.57}$ $^{m}$ & 0.01 $^{x}$ & 14\\ \vspace{1mm}NGC4388 & {\em OSSE }, {\em INTEGRAL} & 7.22 $^{+0.58}_{-0.58}$ $^{\sigma}$ & 0.79 $^{ox}$ & 7\\ \vspace{1mm}NGC1068 & {\em INTEGRAL} & 7.23 $^{+0.58}_{-0.58}$ $^{\sigma}$ & 0.44 $^{f}$ & 8 + 2\\ \vspace{1mm}NGC7582 & {\em OSSE } & 7.25 $^{+0.58}_{-0.58}$ $^{\sigma}$ & 0.35 $^{ox}$ & 7\\ \vspace{1mm}NGC5674 & {\em BeppoSAX} & 7.36 $^{+0.59}_{-0.59}$ $^{\sigma}$ & 0.17 $^{x}$ & 17 (+18)\\ \vspace{1mm}ESO 323-G077 & {\em INTEGRAL} & 7.39 $^{+0.59}_{-0.59}$ $^{m}$ & 0.28 $^{ox}$ & 1\\ \vspace{1mm}NGC4507 & {\em OSSE }, {\em INTEGRAL} & 7.58 $^{+0.61}_{-0.61}$ $^{\sigma}$ & 0.27 $^{ox}$ & 7\\ \vspace{1mm}NGC4258 & {\em BeppoSAX} & 7.62 $^{+0.61}_{-0.61}$ $^{\sigma}$ & 0.01 $^{f}$ & 2\\ \vspace{1mm}NGC7172 & {\em OSSE }, {\em BeppoSAX} & 7.67 $^{+0.61}_{-0.61}$ $^{\sigma}$ & 3x10$^{-3}$ $^{ox}$ & 7\\ \vspace{1mm}NGC2992 & {\em BeppoSAX} & 7.72 $^{+0.62}_{-0.62}$ $^{\sigma}$ & 0.01 $^{f}$ & 2\\ \vspace{1mm}NGC5252 & {\em BeppoSAX} & 8.04 $^{+0.64}_{-0.64}$ $^{\sigma}$ & 0.17 $^{s}$ & 2\\ \vspace{1mm}Centaurus A & {\em INTEGRAL} & 8.04 $^{+0.64}_{-0.64}$ $^{\sigma}$ & 7x10$^{-4}$ $^{s}$ & 11 + 12\\ \vspace{1mm}NGC1365 & {\em BeppoSAX} & 8.18 $^{+0.65}_{-0.65}$ $^{m}$ & 2x10$^{-3}$ $^{x}$ & 19\\ \vspace{1mm}NGC2110 & {\em OSSE }, {\em BeppoSAX} & 8.30 $^{+0.66}_{-0.66}$ $^{\sigma}$ & 5x10$^{-3}$ $^{f}$ & 2\\ \vspace{1mm}NGC1275 & {\em OSSE }, {\em INTEGRAL} & 8.51 $^{+0.68}_{-0.68}$ $^{\sigma}$ & 0.03 $^{f}$ & 2\\ \vspace{1mm}Mrk3 & {\em OSSE }, {\em INTEGRAL} & 8.65 $^{+0.69}_{-0.69}$ $^{\sigma}$ & 6x10$^{-3}$ $^{f}$ & 2\\ \vspace{1mm}Cyg A & {\em INTEGRAL} & 9.40 $^{+0.75}_{-0.75}$ $^{\sigma}$ & 8x10$^{-3}$ $^{x}$ & 15 + 16\\ \hline \end{tabular} \end{minipage} \end{center} \caption{Seyfert 2 sub-sample. Superscripts and references as for Table 1.} \end{table*} | The X-ray spectra of stellar mass black hole binaries can be well described by a combination of emission from the accretion disc and a Comptonised tail (see e.g. the reviews by Remillard \& McClintock 2006; Done, Gierlinski \& Kubota 2007). The relative importance of the disc and tail can change dramatically, as does the shape of the tail, giving rise to the well known 'spectral states' as a function of (average) $L/L_{Edd}$. In the low/hard state, generally seen at luminosities below a few percent of Eddington, the tail dominates the X-ray emission and its spectral index is fairly well correlated with (average) mass accretion rate, with $\Gamma\sim 1.5$ at the lowest luminosities softening to $\Gamma\sim 2.1$ just before the major transition to the soft, or disc dominated states. The tail in these states can be very weak, carrying less than 20 per cent of the total bolometric power, when it has $\Gamma\sim 2.1$ (disc dominant state). Conversely, the tail can also be much stronger and softer, with $\Gamma\sim 2.5$ (very high state or steep power law state). AGN spectra should show analogous spectral states if the properties of the accretion flow scale simply with black hole mass. There is evidence for this simple scaling, from the radio-X-ray `fundamental plane' relations (Merloni, Heinz, \& Di Matteo 2003; Falcke, K\"ording \& Markoff 2004) and X-ray variability properties (McHardy et al. 2007; Gierlinski et al. 2007). In this picture (see e.g. Boroson 2002), LINERS, with their hard X-ray spectra and weak UV disc, are the analogues of the low/hard state (but see Maoz 2007), while Seyfert 1 and QSO's at higher $L/L_{Edd}$ correspond to the disc dominated states, with the characteristic strong UV disc emission and consequent broad emission lines, and the Narrow line Seyfert 1's at the highest $L/L_{Edd}$ may be the analogues of the very high state (Pounds, Done \& Osborne 1995; Middleton, Done \& Gierli\'nski 2007). Observations of type 1 AGN spectra in the 2--10~keV band are consistent with this predicted softening of the intrinsic spectrum with increasing $L/L_{Edd}$ (e.g. Laor 2000), so it seems highly likely that the high energy X--ray spectra considered here should also change with $L/L_{Edd}$. We show that current samples of Seyferts detected at high energies are clearly highly heterogeneous in this parameter (see Fig 1), so the mean spectral index will be determined by a signal--to--noise weighted average of $L/L_{Edd}$. The $L/L_{Edd}$ distributions for the total Seyfert 1 and Seyfert 2 samples presented here are not significantly different (statistically) however the individual OSSE, BeppoSAX and INTEGRAL samples considered here are all consistent with having a higher mean $L/L_{Edd}$ for the Seyfert 1's than for the Seyfert 2s. Thus, despite the limited sample size the apparently steeper spectra seen in the Seyfert 1's at high energy may be be due to an intrinsic softening with increasing $L/L_{Edd}$. Future studies with larger samples of high energy spectra of AGN should be able to test unambiguously whether $L/L_{Edd}$ is the major parameter in determining the shape of the 20--100~keV spectrum in Compton thin Seyferts. \label{lastpage} | 7 | 10 | 0710.5093 |
0710 | 0710.1059_arXiv.txt | In this paper we show that all supergravity billiards corresponding to $\sigma$-models on any $\mathrm{U/H}$ non compact-symmetric space and obtained by compactifying supergravity to $D=3$ are fully integrable. The key point in establishing the integration algorithm is provided by an upper triangular embedding of the solvable Lie algebra associated with $\mathrm{U/H}$ into $\slal(\mathrm{N},\mathbb{R})$ which always exists. In this context we establish a remarkable relation between the arrow of time and the properties of the Weyl group. The asymptotic states of the developing Universe are in one-to-one correspondence with the elements of the Weyl group which is a property of the Tits Satake universality classes and not of their single representatives. Furthermore the Weyl group admits a natural ordering in terms of $\ell_T$, the number of reflections with respect to the simple roots and the direction of time flows is always towards increasing $\ell_T$, which plays the unexpected role of an entropy. | Notwithstanding its length and its somewhat pedagogical organization, the present one is a research article and not a review. All the presented material is, up to our knowledge, new. Due to the combination of several different mathematical results and techniques necessary to make our point, which is instead physical in spirit and relevant to basic questions in supergravity and superstring cosmology, we considered it appropriate to choose the present somewhat unconventional format for our paper. After the theoretical statement of our result, we have illustrated it with the detailed study of a few examples. These case-studies were essential for us in order to understand the main point which we have formalized in mathematical terms in part I and we think that they will be similarly essential for the physicist reader. The table of contents helps the reader to get a comprehensive view of the article and of its structure. \tableofcontents \part{Theory: Stating the principles} | In this paper we have made a few steps forward in developing the general programme of supergravity billiards as a paradigm for superstring cosmology. Our results are both of physical and mathematical nature. \par On the physical side, which for us means supergravity/superstring theory, the essential points are the following ones: \begin{description} \item[1)] We have shown that all supergravity billiards are completely integrable, irrespectively whether they are defined on a maximally split coset manifold $\mathrm{U/H}$ as it happens in the case of maximal susy or a non maximally split $\mathrm{U/H}$, as it happens in all lower supersymmetry cases. We have provided the explicit integration algorithm which just depends on the triangular embedding of the solvable Lie algebra $Solv(\mathrm{U/H})$ into that of $Solv(\mathrm{SL(N)/SO(N)})$. \item[2)] We have discovered a new principle of time orientation of the cosmic flow which relies on the natural ordering of the Weyl group elements (or of the permutations) according to their length $\ell_T$ in terms of transpositions. Cosmic evolution is always in the direction of increasing $\ell_T$ which plays the role of an entropy. There is a fascinating similarity, in this context between the laws of cosmic evolution and those of black hole thermodynamics. \item[3)] We have clarified the meaning of Tits Satake universality classes, introduced in \cite{contoine}, at least from the vantage point of cosmic billiards. The asymptotic states, the type of available flows and the critical surfaces in parameter space are properties of the class and do not depend on the representative manifold in the class. \end{description} On the mathematical side the highlights of our paper are the following ones: \begin{description} \item[1)] We have introduced the notion of generalized Weyl group for a non compact symmetric space $\mathrm{U/H}$ and shown that the factor group with respect to its normal subgroup is just the Weyl group of the Tits Satake subalgebra $\mathbb{U}_{TS} \, \subset \, \mathbb{U}.$ Moreover, we have demonstrated that not only the factor group is isomorphic to the Weyl group of the Tits Satake projection but even the generalized Weyl group is also isomorphic $\mathcal{W}\left(\mathbb{U}\right) \, \sim \, \mathcal{W}\left(\mathbb{U}_{\mathrm{TS}}\right)$. At least this is true in the considered examples and we make the conjecture that it is true in general. \item[2)] We have established a remarkable conjecture encoded in property \ref{kkminorprinc} of the main text: the constraints on minors of the diagonalizing orthogonal matrix for the Lax operator commute with the Toda flow. \item[3)] We have proposed a very simple efficient method of calculating the Toda flow asymptotics at $t=\pm \infty$ for the Lax operator of a $\sigma$-model with target space any non compact-symmetric coset space ${\mathrm{U}}/{\mathrm{H}}$. Our algorithm requires only the knowledge of the corresponding Weyl group $\mathrm{Weyl}(\mathbb{U})$ as well as that of the small group $\mathrm{H}$. \item[4)] We have posed the question how the equations cutting out algebraic loci in compact group or coset manifolds and defined in terms of vanishing minors in the defining representation can be lifted to the abstract group level and extended to all irreducible representations. \end{description} | 7 | 10 | 0710.1059 |
0710 | 0710.1745_arXiv.txt | Using the Green Bank Telescope (GBT) and Pulsar Spigot at 350\,MHz, we have surveyed the Northern Galactic Plane for pulsars and radio transients. This survey covers roughly 1000 square degrees of sky within $75^{\circ} < l < 165^{\circ}$ and $|b| < 5.5^{\circ}$, a region of the Galactic Plane inaccessible to both the Parkes and Arecibo multibeam surveys. The large gain of the GBT along with the high time and frequency resolution provided by the Spigot make this survey more sensitive by factors of about 4 to slow pulsars and more than 10 to millisecond pulsars (MSPs), compared with previous surveys of this area. In a preliminary, reduced-resolution search of all the survey data, we have discovered 33 new pulsars, almost doubling the number of known pulsars in this part of the Galaxy. While most of these sources were discovered by normal periodicity searches, 5 of these sources were first identified through single, dispersed bursts. We discuss the interesting properties of some of these new sources. Data processing using the data's full-resolution is ongoing, with the goal of uncovering MSPs missed by our first, coarse round of processing. | Here we define the ``Northern Sky'' as the sky above the northern Arecibo declination limit of $38^{\circ}$. Previous pulsar surveys of the Northern Sky between $300-500$\,MHz include an all-sky survey above $\delta > 38^{\circ}$ using the Jodrell Bank telescope at 428\,MHz, which, largely due to RFI, discovered only one pulsar \citep{nll+95}; a survey at Green Bank using the 140-ft (43-m) equatorially mounted telescope at 370\,MHz, which found eight new pulsars \citep{snt97}; and a survey using the Green Bank 300-ft (91-m) transit telescope at 390\,MHz, which provided partial coverage of the Northern Sky and discovered 20 pulsars \citep{sstd86}. The GBT's large gain (2\,K/Jy) and maneuverability make it the most sensitive radio telescope on Earth over a large portion of the sky. Furthermore, the primary beam width of the GBT at 350\,MHz is $\sim 0.6^{\circ}$, large enough to permit very efficient large-scale, single-pixel surveys of the sky. This motivated us to commence a 350-MHz survey, which we will refer to as the GBT350 survey, of the Northern Galactic Plane, corresponding to Galactic longitudes $75^{\circ} < l < 165^{\circ}$ and out to latitudes $|b| < 5.5^{\circ}$. Here we present a short summary of the survey design and analysis techniques, and then discuss some of the new sources. | 7 | 10 | 0710.1745 |
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0710 | 0710.3740_arXiv.txt | We use the very large Millennium Simulation of the concordance $\Lambda$CDM cosmogony to calibrate the bias and error distribution of Timing Argument estimators of the masses of the Local Group and of the Milky Way. From a large number of isolated spiral-spiral pairs similar to the Milky Way/Andromeda system, we find the interquartile range of the ratio of timing mass to true mass to be a factor of $1.8$, while the 5\% and 95\% points of the distribution of this ratio are separated by a factor of $5.7$. Here we define true mass as the sum of the ``virial'' masses $M_{200}$ of the two dominant galaxies. For current best values of the distance and approach velocity of Andromeda this leads to a median likelihood estimate of the true mass of the Local Group of $5.27\times 10^{12}\msun$, or $\log M_{LG}/M_\odot = 12.72$, with an interquartile range of $[12.58, 12.83]$ and a 5\% to 95\% range of $[12.26, 13.01]$. Thus a 95\% lower confidence limit on the true mass of the Local Group is $1.81\times 10^{12}\msun$. A timing estimate of the Milky Way's mass based on the large recession velocity observed for the distant satellite Leo I works equally well, although with larger systematic uncertainties. It gives an estimated virial mass for the Milky Way of $2.43 \times 10^{12}\msun$ with a 95\% lower confidence limit of $0.80 \times 10^{12}\msun$. | \label{intro_section} During the 1970's it became generally accepted that most, perhaps all, galaxies are surrounded by extended distributions of dark matter, so-called dark halos \citep{eks74,opy74}. These were soon understood to play an essential role in driving the formation and clustering of galaxies \citep{wr78}. With the introduction of the Cold Dark Matter (CDM) paradigm, these ideas took more concrete form, allowing quantitative predictions to be made both for the population properties \citep{blumenthal84} and for the large-scale clustering \citep{davis85} of galaxies. Measurements of the fluctuation spectrum of the Cosmic Microwave Background \citep{smoot92,spergel03} and of the apparent acceleration of the cosmic expansion \citep{riess98,perlmutter99} elevated the CDM model, in its variant with a cosmological constant ($\Lambda$CDM), to the status of a standard paradigm. At the same time improving numerical techniques and faster computers have enabled detailed simulation of the formation and evolution of the galaxy population within this paradigm throughout a significant fraction of the observable Universe \citep{mr_nature}. Nevertheless, direct observational evidence for halos as extended as the paradigm predicts around galaxies like our own has so far come only from statistical analyses of the dynamics of satellite galaxies \citep[e.g.][]{zaritsky97, prada03} and of the gravitational lensing of background galaxies \citep[e.g.][]{seljak02, mandelbaum06} based on large samples of field spirals. The earliest observational indication that the effective mass of the Milky Way must be much larger than its stellar mass came from the Timing Argument (hereafter TA) of \citet{kw59}. These authors noted that the Local Group is dominated by the two big spirals, and that these are currently approaching each other at about $100\kms$. (The next most luminous galaxy is M33 which is probably about a factor of $10$ less massive than M31 and the Galaxy.) This reversal of the overall cosmic expansion must have been generated by gravitational forces, and since the distance to the nearest external bright galaxy is much greater than that between M31 and the Milky Way, these forces are presumably dominated by material associated with the two spirals themselves. \citeauthor{kw59} set up a simple model to analyse this situation -- two point masses on a radial orbit. These were at pericentre (i.e. at zero separation) at the Big Bang and must have passed through apocentre at least once in order to be approaching today. Clearly this requires an apocentric separation larger than the current separation and an orbital period less than twice the current age of the Universe. Together these requirements put a lower limit on on the total mass of the pair. A more precise estimate of the minimum possible mass is obtained from the parametric form of Kepler's laws for a zero angular momentum orbit: \begin{equation} r=a(1-\cos\chi) \end{equation} \begin{equation} t=\bigg(\frac{a^{3}}{GM}\bigg)^{1/2}(\chi-\sin\chi) \end{equation} \begin{equation} \frac{dr}{dt}=\sqrt{\frac{GM}{a}}\frac{\sin\chi}{1-\cos\chi} \end{equation} where $r$ is the current separation, $dr/dt$ is the current relative velocity, $a$ is the semi-major axis, $\chi$ is the eccentric anomaly, $t$ is the time since the Big Bang (the age of the universe) and $M$ is total mass \citep{lb81}. Given observationally determined values for $r$, $dr/dt$ and $t$, these equations have an infinite set of discrete solutions for $\chi$, $a$ and $M$ labelled by the number of apocentric passages since the Big Bang. The solution corresponding to a single apocentric passage gives the smallest (and only plausible) estimate for the mass, which is about $5\times 10^{12}\msun$ for current estimates of $r$, $dr/dt$ and $t$. Note that this is still a lower limit on the total mass, even within the simple point-mass binary model, since any non-radial motions in the system would increase its present kinetic energy and so increase the mass required to reverse the initial expansion and bring the pair to their observed separation by the present day \citep[see][]{el82}. As \citeauthor{kw59} realised, this timing estimate of the total mass of the Local group exceeds by more than an order of magnitude the mass within the visible regions of the galaxies, as estimated from their internal dynamics, in particular, from their rotation curves. Thus 90\% of the mass must lie outside the visible galaxies and be associated with little or no detectable light. Modern structure formation theories like $\Lambda$CDM predict this mass to be in extended dark matter halos with $M(r)$ increasing very roughly as $r$ out to the point where the halos of the two galaxies meet. Such structures have no well-defined edge, so any definition of their total mass is necessarily somewhat arbitrary. In addition, their dynamical evolution from the Big Bang until the present is substantially more complex than that of a point-mass binary. Thus the mass value returned by the Timing Argument cannot be interpreted without some calibration against consistent dynamical models with extended dark halos. A first calibration of this type was carried out by \citet{kc91} using simulations of an Einstein-de Sitter CDM cosmogony. Here we use the very much larger Millennium Simulation \citep{mr_nature} to obtain a more refined calibration based on a large ensemble of galaxy pairs with observable properties similar to those of the Local Group. We find that the standard timing estimate is, in fact, an almost unbiased estimate of the sum of the conventionally defined virial masses of the two large galaxies. \citet{zaritsky89} attempted to measure the halo mass of the Milky Way alone by measuring radial velocities for its dwarf satellites and assuming the population to be in dynamical equilibrium in the halo potential. They noted, however, that one of the most distant dwarfs, Leo I, has a very large recession velocity and as a result provides a interesting lower limit on the Milky Way's mass by a variant of the original Timing Argument. To reach its present position and radial velocity, Leo I must have passed pericentre at least once since the Big Bang and now be receding from the Galaxy for (at least) the second time. Applying the point-mass radial orbit Equations (1) -- (3) to this case gives a lower bound of about $1.6\times 10^{12}\msun$. This seems likely to be a significant underestimate, since Leo I could not have passed through the centre of the Milky Way without being tidally destroyed so its orbit cannot be purely radial. Below we calibrate the Timing Argument for this case also, finding it to work well although with significantly more scatter than for the Local Group as a whole. This is because the $\Lambda$CDM paradigm predicts that the dynamics on the scale of Leo I's orbit ($\sim 200\kpc$) is typically more complex than on the scale of the Local Group as a whole ($\sim 700\kpc$). Our paper is organised as follows. In Section~\ref{data_section}, we briefly describe the Millennium Simulation and the selection criteria we use to define various samples of `Local Group-like' pairs and of `Milky Way-like' halos. In Section~\ref{result_section}, we plot the ratio of true total mass to Timing Argument mass estimate for these samples, and we use its distribution to define an unbiased TA estimator of true mass with its associated confidence ranges. In Section~\ref{app_LG_section} this is then applied to the Local Group in order to obtain an estimate its true mass with realistic uncertainties. Section~\ref{app_MW_section} attempts to carry out a similar calibration for the TA-based estimate of the Galaxy's halo mass from the orbit of Leo I. We conclude in Section~\ref{conclusions} with a summary and brief discussion of our results. | \label{conclusions} The statistical argument underlying the analysis of this paper is more subtle than it may at first appear, so it is worth restating it somewhat more formally in order to understand what is being assumed in deriving the mass estimates for the Local Group and for the Milky Way given above. We believe that the mass distributions around galaxies are much more extended than the visible stellar distributions, and that these have been assembled from near-uniform ``initial'' conditions in a manner at least qualitatively resembling that in a $\Lambda$CDM universe. Thus the assembly histories of the Local Group and of the Milky Way's halo differ in major ways from those assumed by the original Timing Arguments of \citet{kw59} and \citet{zaritsky89}. In addition, the meaning of the derived mass values needs clarification. We wish to use the Millennium Simulation to calibrate the TA estimates against conventional measures of halo mass, and to test the general applicability of the Timing Argument. However, we want to do this in a way which avoids any significant dependence on the details of the $\Lambda$CDM model, for example, on the exact density profiles, abundances and substructure properties which it predicts for halos. Our method uses the simulation to estimate the distribution of the ratio of ``true'' mass to TA mass estimate for samples of objects whose properties ``resemble'' those of the observed Local Group and Milky Way -- Leo I systems. Our restrictions on separation and radial velocity implement this similarity requirement in a straightforward way, but our constraints on $V_{max}$ have a more complex effect. Although the true $V_{max}$ values for M31 and the Milky Way are very likely within our looser range ($150\kms \le V_{max} < 300\kms$) the simulation exhibits a tight correlation between $V_{max}$ and $M_{200}$ . Imposing fixed limits on $V_{max}$ is thus effectively equivalent to choosing a fixed range of $M_{200}$. As a result, we are in practice estimating the distribution of $A_{200}$ or $B_{200}$ for systems of given {\it true} mass, subject to the assumed constraints on separation and radial velocity. However, when we apply our results to estimate true masses for the Local Group and the Milky Way, we implicitly assume that our distributions of $A_{200}$ and $B_{200}$ are appropriate for samples of given TA mass estimate, again subject to our constraints on separation and radial velocity. It is thus important to understand when these two distributions can be considered the same. The relation can be clarified as follows. From the simulation we compile the distribution of $M_{tr}/M_{TA}$, or equivalently of $\Delta \equiv \ln M_{tr} - \ln M_{TA}$, for systems with $\ln M_{tr}$ in a given range. We then implicitly assume that this distribution does not depend on $M_{tr}$, at least over this range, so that the result can be taken as an estimate of the probability density of $\Delta$ at given $M_{tr}$. Bayes Theorem then gives us the probability density function (pdf) for $\Delta$ at fixed $M_{TA}$: \begin{eqnarray} f[\Delta | \ln M_{TA}] = \frac{f[\Delta, \ln M_{TA}]}{f[\ln M_{TA}]}\nonumber\\ = \frac{f[\Delta, \ln M_{tr}]}{f[\ln M_{TA}]} \nonumber\\ = \frac{f[\Delta | \ln M_{tr}]\, f[\ln M_{tr}]}{ f[\ln M_{TA}]}\nonumber\\ = f[\Delta | \ln M_{tr}] \label{Bayes} \end{eqnarray} The first line here simply writes the conditional pdf of $\Delta$ at given $M_{TA}$ in terms of the joint pdf of the two quantities and the pdf of $M_{TA}$. The second line then rewrites the joint pdf in terms of the equivalent variables $\Delta$ and $M_{tr}$, using the fact that the Jacobian of the transformation is unity. The third line re-expresses the joint pdf as the product of the pdf of $\Delta$ at given $M_{tr}$ times the pdf of $M_{tr}$. The final line then follows from the normalisation condition, {\it provided} that $f[\ln M_{tr}]$ is constant and $f[\Delta | \ln M_{tr}]$ is independent of $M_{tr}$. Thus, when estimating $M_{tr}$ from $M_{TA}$, we assume a uniform prior on $\ln M_{tr}$ and that the distribution of $\Delta$ does not depend on true mass. Both these assumptions appear natural and appropriate. The analysis underlying the Timing Argument (Equations (1) -- (3)) assumes that the relative orbit of the two objects is bound and has conserved energy since the Big Bang. Recently, \citet{sales07} have shown that in $\Lambda$CDM models this assumption is significantly violated for a substantial number of satellites within halos comparable to that of the Milky Way. In particular, they demonstrate the presence of a tail of unbound objects which are being ejected from halos as a result of $3$-body ``slingshot'' effects during their first pericentric passage. These objects are typically receding rapidly from their ``Milky Way'', as assumed in the \citet{zaritsky89} argument, but they violate its assumption that the present orbital energy can be used to infer the period of the initial orbit (i.e. the time from the Big Bang to first pericentric passage). Clearly such objects should also be present in the Millennium Simulation, although lack of resolution might make them under-represented in comparison to the simulations analysed by \citet{sales07}. Thus our analysis takes the possibility of such ejected satellites into account, at least in principle. Objects of this type will show up as systems with anomalously large TA estimates for their halo mass, and Fig.~\ref{mw_mass_scatter} shows a number of outliers which could well be explained in this way. Issues of this kind do not effect TA-based estimates of the mass of the Local Group since the two big galaxies are currently approaching for the first time. The only kinematic information about the relative orbit of M31 and the Milky Way used in our analysis is their current approach velocity. \citet{vdmg07} show that geometric arguments can already constrain the transverse component also, and future astrometry missions such as SIM might allow $V_{tr}$ to be measured directly. Thus it is interesting to ask if our TA mass estimate could be significantly refined by measuring the full 3-D relative motion of the two galaxies, rather than just its radial component. We address this in Fig.~\ref{a200_vtr_scatter} which plots $A_{200}$, the ratio of true mass to TA estimate, against $V_{tr}$ for a sample of Local Group analogues with our preferred morphology, isolation and radial velocity cuts, and with $150\kms \le V_{max} < 300\kms$. The median $V_{tr}$ for this sample is $86 \kms$, comparable to the \citet{vdmg07} estimate for the real system. There is no apparent correlation of $A_{200}$ with $V_{tr}$, and indeed, the medians of $A_{200}$ for the high and low $V_{tr}$ halves of the sample are both close to 1 and do not differ significantly. Pairs with high $V_{tr}$ do show larger {\it scatter} in $A_{200}$ than pairs on near-radial orbits. For given separation, radial velocity and age, the Kepler model implies a mass which increases monotonically with $V_{tr}$. The absence of a detectable trend in Fig.~\ref{a200_vtr_scatter} shows that uncertainties in $V_{tr}$ do not dominate the scatter in our TA mass estimate, and that a measurement of $V_{tr}$ will not substantially increase the precision with which the true mass can be measured. \begin{figure} \includegraphics[width=0.5\textwidth]{a200_vtr.ps} \caption{Scatter plot of $A_{200}$ versus transverse velocity for Local Group analogues in our sample with preferred morphology, isolation and radial velocity cuts and with $150\kms \le V_{max} < 300\kms$. The vertical dashed line indicates the median $V_{tr}$. The distributions on either side of this line are each further split in half at the median values of $V_{tr}$ (the solid horizontal lines.) There is essentially no correlation in this plot, indicating that a measurement of the transverse velocity will not significantly improve the TA mass estimate.} \label{a200_vtr_scatter} \end{figure} In conclusion, our analysis shows the Timing Argument to produce robust estimates of true mass both for the Local Group and for the Milky Way, as long as ``true mass'' is understood to mean the sum of the conventional masses of the major halos. For the Local Group as a whole, the estimate and confidence limits given in Section~\ref{app_LG_section} and in the Abstract appear reliable given the excellent statistics provided by the Millennium Simulation, the lack of any substantial dependence on our isolation and morphology cuts, and the relatively simple dynamical situation. Although the results based on Leo I's orbit also appear statistically sound, the more complex dynamical situation offers greater scope for uncertainty, particularly when trying to place a lower limit on the mass of the Milky Way's halo. On the other hand, our best estimate of this mass is just under half of our estimate of the sum of the halo masses of M31 and the Galaxy. Thus the picture presented by the data appears quite consistent, and gives no reason to be suspicious of the Milky Way results. | 7 | 10 | 0710.3740 |
0710 | 0710.2436_arXiv.txt | The luminosities of the optical afterglows of Gamma Ray Bursts, 12 hours (rest frame time) after the trigger, show a surprising clustering, with a minority of events being at a significant smaller luminosity. If real, this dichotomy would be a crucial clue to understand the nature of optically dark afterglows, i.e. bursts that are detected in the X--ray band, but not in the optical. We investigate this issue by studying bursts of the pre--{\it Swift} era, both detected and undetected in the optical. The limiting magnitudes of the undetected ones are used to construct the probability that a generic bursts is observed down to a given magnitude limit. Then, by simulating a large number of bursts with pre--assigned characteristics, we can compare the properties of the observed optical luminosity distribution with the simulated one. Our results suggest that the hints of bimodality present in the observed distribution reflects a real bimodality: either the optical luminosity distributions of bursts is intrinsically bimodal, or there exists a population of bursts with a quite significant grey absorption, i.e. wavelength independent extinction. This population of intrinsically weak or grey--absorbed events can be associated to dark bursts. | Since the detection of the first long Gamma Ray Burst (GRB) optical afterglow (Van Paradijs et al. 1997), the non--detection of any optical source in the direction of the gamma--ray trigger of some events stimulated the interest about the possible differences between the nature of the afterglow emission of the optically bright and faint GRBs. During the last 9 years the increasing number of optical detections and spectroscopic redshift determinations allowed us to study the intrinsic features of the optical afterglow emission of long GRBs. Despite the improved (i.e. made more promptly) observations of optical afterglows, in almost half of the observed long GRBs no optical counterpart is still found. These events have been called in the literature as Dark Burst, or Failed Optical Afterglows GRBs (Lazzati et al. 2002). In Nardini et al. (2006a) \footnote{Liang \& Zhang (2006) independently found similar results.} we showed the optical $R$ band luminosity light curves of a sample of 24 pre--{\it Swift} GRBs with known spectroscopic redshift and published estimate of host galaxy dust absorption. We found a strong clustering of that luminosities for the GRBs in our sample. Most (i.e. 21/24) of the $R$ band luminosities at 12 hours in the source frame are clustered within a log--normal distribution centred around a mean value $\log L_{\nu_R}=30.65$ [erg s$^{-1}$ Hz$^{-1}$] with a dispersion $\sigma=0.28$. We also found 3 GRBs showing dimmer luminosities, a factor 15 (from 3.6 to 4.6 $\sigma$) smaller than the mean of the higher luminosity distribution. No GRB was found in the luminosity range between these two ``families''. In Fig \ref{istobs} we show the histogram of the $R$ band luminosities of our sample of GRBs. In a recent update (Nardini et al. 2006b) we added 8 new GRBs detected by the {\it Swift} satellite (Gehrels et al. 2004) for whose an estimate of the host galaxy dust extinction has been published. This small sample of {\it Swift} GRBs confirms both the clustering and the bimodality of the optical luminosities found by us with pre--{\it Swift} bursts. We also evaluated the optical luminosities for 17 other {\it Swift} GRBs without any published $A_V^{host}$ estimate, and found that they are consistent with our previous findings. \begin{figure} \vskip -0.5 true cm \centerline{\psfig{figure=istolumvecchi.eps,angle=0,width=10cm}} \vskip -0.8 true cm \caption{ Histogram of the monochromatic optical luminosities 12 hours (rest frame) after the trigger for the 24 GRBs analysed in Nardini et al. (2006a). Data have been de-reddened both for Galactic and host extinction. } \label{istobs} \end{figure} The discovery of a family of optically dim GRBs is an important clue for the understanding of the nature of dark bursts. The few underluminous observed events could be the tip of the iceberg of a population of GRBs which are intrinsically less luminous. Therefore, a fraction of dark GRBs could belong to this family whose distance, optical absorption, or observing conditions do not allow any optical detection. Because of its potential importance for the understanding of dark bursts, we investigate, in this paper, if the observed bimodal luminosity distribution of optically--bright bursts is due to any selection effect (related to the search/detection of GRB optical counterparts) or if it reflects the existence of two GRB populations. To this aim we simulate through a Montecarlo method a sample of GRBs with a redshift distribution traced by the cosmic star formation rate (Porciani \& Madau 2001), assuming different shapes of their intrinsic optical luminosity function. We also simulate different values of dust absorption within the host galaxy to all the simulated events. This value is calculated assuming the standard extinction curves (Pei 1992). We infer a limiting magnitude distribution obtained by the analysis of the deepest $R$ band upper limits of all the pre--{\it Swift} GRBs with no detection of their optical afterglow. This is the key point of our study: the use of the upper limits on the optical flux to construct the probability that a simulated bursts would be detected or not. It is this probability distribution that allows us to perform, meaningfully, our simulations. We then compare the resulting luminosity distribution of the detectable simulated events with the one obtained in Nardini et al. (2006a) and shown in Fig \ref{istobs}. The scope of our simulation is to check if for any conceivable combination of the input assumptions (i.e. luminosity function, extinction and redshift distribution) we can reproduce a simulated sample whose $R$--band luminosity distribution (12h rest frame) is consistent with that observed with the sample of 24 pre--{\it Swift} GRBs. For our analysis we use only the pre--{\it Swift} GRBs because they represent an homogeneous sample: their optical light curves were sampled from hours to days after the trigger and the estimates of the host galaxy extinction have been published. The same study can be repeated using the more recent {\it Swift} GRBs once we will have a sufficient number of events with published estimates of the host galaxy extinction (now this number is too small, see Nardini et al. 2006b). | In this study we analysed the possible importance of observational selection effects on the clustering and the bimodality found in the long GRBs optical afterglow luminosity distribution. \begin{itemize} \item We studied the $R$ band upper limits of all the pre--{\it Swift} dark GRBs and we showed that they are consistent with the ``typical'' afterglow detections. The distribution of these upper limits extrapolated at 12 h after trigger enables the definition of a Telescope Selection Function, which can define the probability, for any burst, to be pointed with a telescope and exposure time corresponding to some limiting magnitude. \item Through Montecarlo simulations, we have studied which combinations of optical luminosity functions and absorption in the host galaxy can be consistent with the observations. In doing so, we have found that the gap between the two observed luminosity distribution is real and it is not the result of a small number statistics. \item If the absorption is chromatic (i.e. a ``standard" one) no unimodal intrinsic luminosity distribution agrees with the observed one. If a good fraction of events (but not all) is absorbed by ``grey" (i.e. achromatic) dust, then a unimodal luminosity distribution is possible. An achromatic absorption would not be recognised through the standard $A_V^{host}$ estimate methods. This can lead to underestimate the intrinsic luminosity for a fraction of bursts that could thus appear as members of an underluminous family. \item Dark bursts could then be associated either to an intrinsically optically underluminous family, or to those bursts being characterised by a relatively large achromatic absorption. In the first case one should explain why, optically, the luminosity distribution is bimodal, while other properties are not (e.g. the distribution of the energetics of the prompt emission), while in the second case one should explain why a fraction of bursts live in a different environment, characterised by grey dust. \end{itemize} | 7 | 10 | 0710.2436 |
0710 | 0710.0096_arXiv.txt | We have computed the size distribution of silicate grains in the outer radiative region of the envelope of a protoplanet evolving according to the scenario of \citet{pollack96}. Our computation includes grain growth due to Brownian motion and overtake of smaller grains by larger ones. We also include the input of new grains due to the breakup of planetesimals in the atmosphere. We follow the procedure of \citet{podolak03}, but have speeded it up significantly. This allows us to test the sensitivity of the code to various parameters. We have also made a more careful estimate of the resulting grain opacity. We find that the grain opacity is of the order of $10^{-2}\unit{cm^2~g^{-1}}$ throughout most of the outer radiative zone as \citet{hubickyj} assumed for their low opacity case, but near the outer edge of the envelope, the opacity can increase to $\about{1}\unit{cm^2~g^{-1}}$. We discuss the effect of this on the evolution of the models. | The planets in our solar system were formed from material left over after the formation of the Sun, but the details of the formation process are still under investigation. It is believed that the gas giant planets of our solar system, Jupiter and Saturn, have a solid core embedded within a gaseous envelope \citep{saumon04}. The most popular theory for the formation of these planets supposes that they were formed in two stages: A solid core was first formed by accretion of planetesimals in the protoplanetary disk, and, when the core became massive enough to gravitationally attract and capture gas from the disk, the gaseous envelope was added. Detailed numerical simulations of these formation stages have been performed for this core accretion model \citep{pollack96,hubickyj}. These simulations show three distinct phases. During the first phase the planet still consists primarily of solid material, the gas accretion rate is small, and the planetesimal accretion rate increases rapidly until the planet's feeding zone is depleted. The planet then enters a second phase during which the rates of solid and gas accretion are both small and nearly constant. When the gas mass of the planet is about equal to its solid mass the third phase begins, characterized by runaway accretion of gas. Simulations of this kind, using reasonable parameters, have been successful in producing planets with properties similar to those of the actual giant planets, however the formation time for a Jupiter-like planet at $5\unit{AU}$ from a Sun-like star is uncomfortably close to the estimated lifetime of the protoplanetary disk in which it is to be formed. Observations of disks around young stars give estimates of 0.1--10~{Myr} for their lifetimes \citep{haisch01}, while numerical simulations are able to produce a planet in times varying from a few to many millions of years, depending on several key parameters \citep{hubickyj}. One of these key parameters is the opacity of the gaseous envelope during the second of the phases described above. In particular, the contribution to the opacity from aerosols is an important factor whose value is uncertain. Grain opacity is difficult to estimate, because it depends on the size distribution as well as the composition of the aerosols. It has been demonstrated \citep{hubickyj} that arbitrarily lowering the opacity in the outer radiative zone by a factor of 50 can reduce the formation time by roughly a factor of 3. \citet{podolak03} has shown that the size distribution of grains in the envelope can be quite different from the interstellar size distribution usually assumed in opacity calculations, and that this can result in a significantly lowered opacity. The purpose of this work is to make this conclusion more robust by repeating the calculation with greater accuracy, and by exploring further the parameter space. We attempt to identify under which circumstances it is correct to assume that the grain opacity, in the context of the core accretion model, is much lower than the interstellar value, and to suggest the value that should be used instead. We use the same microphysics model that was described in \citep{podolak03}, but we have made the code much more efficient so that we can more readily test the sensitivity to various assumptions. We have also done a more careful computation of the actual opacity of the computed distribution. In section~\ref{sec:model} we review the details of the microphysical model. In section~\ref{sec:results} we present the results of our calculations and some tests of the sensitivity of the model to the various input parameters, and in section~\ref{sec:conclusion} we give our conclusions. | Conclusions} Before any conclusions can be drawn from the above results about the effect of grain growth on planet formation it is important to note again one important assumption of the calculations made here, the assumption of a static atmosphere. As was previously mentioned, the radius-density-temperature relation of the atmosphere remains constant throughout the calculation. This assumption was justified because the time to reach a steady state of the grain size distribution is much shorter than the time scale of the atmosphere's evolution. We must remember however that the structure of the atmosphere which is used in the present calculation was determined \emph{without} knowledge of the grain size distribution, and with previously chosen grain opacity. The grain opacity itself will, of course, influence the evolution of the planet's atmosphere. In this sense, the calculations carried out here are not completely self consistent. It is possible that using the opacity profiles shown above in a planet evolution model would result in different atmosphere structures which, if used as background for the grain growth process, would result in different opacity profiles. With this reservation in mind, there is at least one conclusion that can still be stated with some confidence about the effects of grain growth on the atmosphere's opacity, and that is that grain growth \emph{does affect} the atmosphere's opacity. More precisely: The conditions during phase 2 of core accretion are such that growth and settling are efficient processes by which grains are removed from the radiative zone of the planet's atmosphere, reducing their contribution to the opacity. For most of the radiative zone the resulting grain opacity is significantly lower than it would be if grain growth were not accounted for. The implication for giant planet formation theory would be that the formation time of a Jupiter-like planet through the core accretion scenario may be shortened, alleviating one of the major points of concern about this model. As for the actual value of grain opacity that should be used in planet evolution models, our simulations indicate that the opacity is not reduced by a constant factor throughout the atmosphere. Instead, it is roughly equal to the interstellar grain value in the uppermost layers, where the grains reside for a relatively short time, and grain growth is not an important process. Deeper in the atmosphere, where the grain microphysics has time to adjust the grain size distribution, the opacity can be reduced over the interstellar value by as much as three orders of magnitude. In the uppermost parts of the atmosphere, where the optical depth is $\tau\leq\sim 10$ the optical depth profile is close to that computed for the high opacity case of \citet{hubickyj}, but as you descend into the deeper layers of the radiative zone, the optical depth remains well below the profile derived for the high opacity case. Except for the deepest region of the radiative zone, the optical depth in our models is always higher than that derived for the low opacity case of \citet{hubickyj}. It seems that we must combine a microphysics calculation of the grain distribution together with an evolutionary model of the sort presented in \citet{hubickyj} in order to determine the true effect of the grain size distribution on the contraction time of the protoplanetary envelope. In the present work, evaporation of grains was completely neglected. If there are ice grains present, they will evaporate once the temperature exceeds $\about 200$ K. Organic material in the grains will either evaporate or pyrolyze in the hydrogen-rich atmosphere. Ice and organics are expected to make up about $2/3$ of the high-Z material by mass, so the grain abundance throughout most of the region we have considered will be only about 30\% of what we have considered for the baseline model. As we showed above, such a reduction in the rate of grain accretion has only a small effect on the resulting distribution. However, it must be remembered that in the upper layers, where the temperatures are low enough, grains can accumulate ice and organics. These components will be lost in the deeper layers. If the ice acts as a ``glue'' to hold the silicates together, then we would expect the larger grains to break up upon reaching the deeper, hotter layers. At the very least the grains will develop voids and possibly a fractal structure. This will lower their mean density and have a substantial effect on the size distribution. Thus our conclusions about the opacity must be seen as tentative, subject to future work. | 7 | 10 | 0710.0096 |
0710 | 0710.4136_arXiv.txt | The extragalactic cosmic gamma-ray background (CGB) is an interesting channel to look for signatures of dark matter annihilation. In particular, besides the imprint in the energy spectrum, peculiar anisotropy patterns are expected compared to the case of a pure astrophysical origin of the CGB. We take into account the uncertainties in the dark matter clustering properties on sub-galactic scales, deriving two possible anisotropy scenarios. A clear dark matter angular signature is achieved when the annihilation signal receives only a \emph{moderate} contribution from sub-galactic clumps and/or cuspy haloes. Experimentally, if galactic foregrounds systematics are efficiently kept under control, the angular differences are detectable with the forthcoming GLAST observatory, provided that the annihilation signal contributes to the CGB for a fraction $\agt$10-20\%. If, instead, sub-galactic structures have a more prominent role, the astrophysical and dark matter anisotropies become degenerate, correspondingly diluting the DM signature. As complementary observables we also introduce the cross-correlation between surveys of galaxies and the CGB and the cross-correlation between different energy bands of the CGB and we find that they provide a further sensitive tool to detect the dark matter angular signatures. | Astronomical and cosmological observations provide overwhelming evidence for the presence of dark matter (DM) (see e.g.\ \cite{Bertone:2004pz} for a review). In particular, the combination of various cosmological data sets provides a precise measurement of the amount of DM in the universe: $\Omega_c h^2 \simeq 0.11$ with a $2 \sigma$ precision of $\sim$ 5\% in the minimal $\Lambda$CDM model \cite{Tegmark:2006az,Spergel:2006hy,Seljak:2006bg} and $\sim$ 20\% in more extended models \cite{Hamann:2006pf}. However, despite the noticeable sensitivity to the cosmological abundance of matter (either dark or baryonic), such measurements only weakly constrain the properties and nature of the particle associated to DM, and very weak limits are available on the DM particle mass $m_\chi$ and on its couplings. The simplest DM candidate is the Weakly Interacting Massive Particle (WIMP) which is characterized by having been in thermal equilibrium in the early universe (as opposed to for example the sterile neutrino or super-heavy DM), and having decoupled from equilibrium while non-relativistic. In order to get the correct DM abundance the mass of such a particle cannot be larger than $\sim 30$ TeV \cite{Bertone:2004pz,Griest:1989wd}. On the other hand, collider experiments provide a lower bound on the mass of $\sim 50-100$ GeV \cite{Bertone:2004pz}, depending on the specific particle candidate. Mass of $\mathcal{O}$(GeV) are however possible if more exotic candidates are considered \cite{Gunion:2005rw}. The typical thermally averaged DM annihilation cross section in the WIMP scenario is $ <\!\!\! \sigma_\chi v \!\!\! >$$\sim 10 ^{-26}$ cm$^{3}$s$^{-1}$ \cite{Bertone:2004pz}. However, we stress that if the DM is produced out of equilibrium in the early universe, no bounds can be given and super-massive, GUT scale, DM particles ($m_\chi \sim 10^{15}$ GeV) and cross sections $ <\!\!\! \sigma_\chi v \!\!\! >$ $ \ll 10^{-26}$ cm$^{3}$s$^{-1}$ are in principle possible. From the point of view of particle physics WIMP candidates are very appealing and emerge naturally in Supersymmetric (SUSY) extensions of the standard model or in the Universal Extra Dimensions (UED) model \cite{Bertone:2004pz}. The sensitivity of accelerator experiments, notably the Large Hadron Collider, and of direct search experiments are approaching the levels required to test the WIMP hypothesis, and a direct discovery of DM WIMPs could happen in the not so distant future. DM WIMP candidates have thus typically a large annihilation cross section and pair-annihilate into standard model particles that subsequently decay and shower producing large numbers of photons and neutrinos. Such $\gamma$-rays from DM annihilation constitute an ideal target for astronomical searches. Thus, astrophysical and cosmological observations can provide a crucial test, complementary to a direct laboratory detection, in the search for the nature of DM particles. Various astrophysical environments have been discussed in detail as promising sites for observation of DM annihilation, among others the galactic center, satellite dwarf galaxies of the Milky Way and clumps of DM in the Milky Way halo. In the following we will focus instead on the all-sky diffuse signal expected in the extragalactic cosmic gamma-ray background (CGB) \cite{Bergstrom:2001jj,Ullio:2002pj,Elsaesser:2004ck,Elsaesser:2004ap}. Peculiar spectral and angular features can help in disentangling a signal produced by DM from emission by ``ordinary'' astrophysical sources. The spectrum of photons from DM annihilation is in general harder than the spectra arising from normal astrophysical processes and exhibit a pronounced cutoff at an energy near $m_\chi$ \cite{Bergstrom:2001jj,Ullio:2002pj}. The resulting emission thus appear like a ``bump'' in the background astrophysical energy spectrum in the energy range in which the DM signal gives a relevant contribution. However, although this kind of signature would constitute a strong hint of DM annihilation, astrophysical processes that could mimic such behavior are possible. Another signature, which has been widely studied, is direct annihilation into a state containing photons, resulting in a line in the background spectrum that would constitute a ``smoking gun'' signature of DM. However, by construction this process is necessarily loop suppressed and in most models the flux is quite small (see, however, \cite{Ullio:2002pj} for a more thorough discussion). Peculiar angular signatures thus offer a complementary signature to exclude the remaining degenerate astrophysical interpretations of a signal. An example is the clumpiness of DM at sub-galactic scales \cite{Berezinsky:2005py,Taylor:2002zd,Pieri:2007ir,Bergstrom:1998jj} investigated by recent zoomed high-resolution N-body simulations \cite{Diemand:2006ik,Diemand:2005vz}: Clumpiness would result in a population of high galactic latitude extended gamma emitters with a typical annihilating DM gamma spectrum. These kinds of objects could hardly be associated to astrophysical emitters (but see \cite{Baltz:2006sv}). In these models the size of the clumps is expected to have a characteristic distribution and thus the anisotropy of the integrated signal from all the clumps also exhibits a characteristic behavior \cite{Pieri:2007ir}. Likewise, the expected angular anisotropies both in the case of an astrophysical and of a DM origin of the CGB can be calculated, and have received increasing attention in the last few years \cite{Cuoco:2006tr,Ando:2005xg,Ando:2006mt,Ando:2006cr,Miniati:2007ke}. In the following we will further pursue this issue addressing the differences expected in the two cases and their detectability in the light of the improved statistics that will be available, when the GLAST observatory is launched and start to take data in the near future. We will compare throughout the paper our findings in particular with \cite{Ando:2005xg,Ando:2006cr} that deal specifically with anisotropies induced by DM annihilation. Already, there have been claims \cite{Elsaesser:2004ap,deBoer:2005tm} of a DM signal in the CGB as observed by EGRET (see also \cite{Bergstrom:2006tk,deBoer:2006ck,deBoer:2007zc}), although with the limited EGRET statistics and with the uncertainties in the galactic foregrounds, alternative astrophysical explanations cannot be ruled out. On the other hand, with the improved statistics from GLAST, a proper analysis of the anisotropy properties of the CGB should be able to prove, or disprove, the DM interpretation of features in the CGB spectrum. Complementary to previous studies we shall employ in the following a parametric approach characterizing the expected CGB signal in terms of a few key parameters, that catch the relevant physical aspects of the problem, and varying them in order to asses the robustness and/or model dependence of the possible signatures. A further advantage of this approach is to make explicit the various assumptions employed throughout on which the final signatures depend. The relevant parameters in the following will be the degree of correlation of the CGB sources with matter and the absolute normalization of the signal, or, equivalently, the expected collected statistics. Further, we will also consider complementary anisotropy observables like the cross-correlation between surveys of galaxies and the CGB and the cross-correlation between different energy bands of the CGB. Together with the auto-correlations of the CGB these represent a set of independent observables that can be jointly employed improving considerably the sensitivity to the DM signal. The paper is organized as follows: In section \ref{CGBmain} we present a discussion of the horizons within which the CGB signal is expected to come, relevant for the determination of the intensity of the CGB anisotropies itself. In section \ref{CGBanisotropies} we introduce the formalism to derive the CGB anisotropies in terms of the angular power spectrum. In section \ref{GLASTforecast} we present a forecast for the expected statistics from GLAST and we discuss the possibility of disentangling the DM annhilation signal from that of astrophysical processes. In sections \ref{crosscorr} and \ref{crosscorr_eb} we introduce the cross-correlation between the CGB and galaxy surveys and the cross-correlation between different energy bands of the CGB and similarly we discuss the different behavior and sensitivity in the two cases of interest. In section \ref{discussion} we discuss how the previous conclusions apply to different possible scenarios for the CGB and DM properties. In section \ref{final} we summarize and conclude. | \label{final} In the present work we have studied the kind of signatures that DM annihilation is expected to imprint in the anisotropies of the CGB, complementary to the signatures in the energy spectrum. We have addressed the main physical ingredients contributing to the DM signature and discussed the robustness of the signature with respect to various possible scenarios. We can summarize our findings as follows: \begin{itemize} \item{The DM annihilation signal traces in general the matter distribution quadratically due to its $\rho_\chi^2$ dependence. However, an effective linear correlation can arise if the signal is significantly enhanced by the presence of cuspy haloes or sub-galactic clumps. We have defined the two extreme ``linear'' and ``quadratic'' scenarios. The first corresponds to the case in which the cosmological DM annihilation signal is dominated by galactic or sub-galactic structures while in the second the signal is dominated by emission on scales larger than that of a galactic halo. We have chosen a phenomenological approach introducing a parameter $\xi$ that weights the two relative contributions exploring the DM signatures for different possible choices of $\xi$.} \item{The anisotropies are determined both by the intrinsic fluctuations in the source field and by the size of the emission horizon $z_{\mathcal{H}}$. For $E_{\rm cut}\agt 100$ GeV the horizon $z_{\mathcal{H}}$ is essentially fixed by photon absorption in the EBL. The bulk of the gamma-rays is expected to originate inside $z_{\mathcal{H}}\approx 0.5$, independent of whether they have a DM or an astrophysical origin. For $E_{\rm cut}\agt 10$ GeV, DM annihilation in the quadratic scenario has a redshift horizon $z_{\mathcal{H}}\approx 3-4$. The horizon is still significantly limited by PP losses at this energy, otherwise exceeding $z_{\mathcal{H}}\approx 10$. Blazars and DM in the linear scenario have degenerate horizons $z_{\mathcal{H}}\approx 1$. } \item{ In the quadratic scenario the DM anisotropy signal is sensibly enhanced with respect to blazars for $E_{\rm cut}=$10 GeV. Further, also the shapes of the angular spectra differ significantly \cite{Ando:2005xg,Ando:2006cr}. The signature remain standing also for $E_{\rm cut}=$ 100 GeV despite the decreased statistics and become particularly strong, being independent of uncertainties related to the blazar-matter bias or to the evolution of blazars. This scenario can easily be detected by GLAST and would constitute a strong signature of DM annihilation. The DM linear scenario, instead, exhibit the same level of fluctuations of blazars and the two thus have almost degenerate anisotropy features.} \item{The above signature in the angular spectrum remains quite robust as long as the the quadratic DM signal is at least at the 10-20\% level with respect to the linear DM or the blazar component. A further uncertainty to take into account is the normalization of the CGB (and thus the available statistics) that is likely to be reduced if part of the sources contributing to the CGB will be resolved by GLAST. If the normalization is reduced by an extreme factor of 5, (20\% of the present EGRET value), the pure quadratic DM scenario exhibits still a relevant anisotropy transition signature. If the quadratic DM contributes for a 20\% (thus, 4\% of the present EGRET value) then the detection of the signature becomes quite challenging.} \item{The cross-correlation between the CGB and a survey of galaxies and the cross-correlation between different energy bands of the CGB provide further independent and sensitive observables that can be employed in combination with the CGB auto-correlation. A joint analysis of all the anisotropy observables considerably improves the sensitivity to the DM signal and, more in general, the power of the statistical analysis. In principle, the exact contribution from DM annihilation in sub-galactic clumps and cuspy haloes can be treated as free parameters (instead of relying on a model) and inferred from the analysis.} \end{itemize} The above conclusions hold exactly if a perfect cleaning of the galactic foregrounds and a lossless extraction of the CGB signal is possible. The analysis of foregrounds will be likely the main challenge in the study of the CGB. Clearly, given the above shown potential of CGB anisotropies in looking for DM signatures, it would be worth to perform further detailed studies on the issue. The launch of the GLAST satellite is expected by the middle of 2008, while the satellite AGILE \cite{AGILE} launched in April 2007 is currently already taking data. The improvement in statistics compared to EGRET will allow for new, powerful tools to search for exotic contributions to the gamma-ray signal. The anisotropy analysis of the CGB in particular, if foregrounds contaminations can be efficiently kept under control, promises to provide a clear signature of DM annihilation or, in the case of a negative answer, to obtain new constraints on the DM properties, complementary to a pure energy spectrum analysis. \vspace{1pc} | 7 | 10 | 0710.4136 |
0710 | 0710.4070_arXiv.txt | {We investigate how different models that have been proposed for solving the dark matter problem can fit the velocity dispersion observed around elliptical galaxies, on either a small scale ($\sim$ 20kpc) with stellar tracers, such as planetary nebulae, or large scale ($\sim$ 200kpc) with satellite galaxies as tracers. Predictions of Newtonian gravity, either containing pure baryonic matter, or embedded in massive cold dark matter (CDM) haloes, are compared with predictions of the modified gravity of MOND. The standard CDM model has problems on a small scale, and the Newtonian pure baryonic model has difficulties on a large scale, while a fit with MOND is possible on both scales. | Measuring the velocity field in and around galaxies is the main way to test the dark matter distribution at small and intermediate scales. The observation of what are apparently non-Newtonian rotation curves around spiral galaxies (e.g. Rubin et al. 1980) has been first solved by assuming that galaxies are embedded in dark matter haloes. Numerical simulations, however, predict a radial distribution for the CDM model much more concentrated than what is observed (Gentile et al 2004; de Blok 2005). An alternative explanation was proposed by Milgrom (1983), as MOdified Newtonian Dynamics (MOND). When the Newtonian acceleration falls below the critical value $a_0 \sim 2 \times 10^{-10}$ m s$^{-2}$, the gravity law is empirically modified and then declines in 1/$r$ instead of 1/r$^2$. Around spherical systems, the modified acceleration $g$ satisfies the relation $$ g \mu(g/a_0) = g_n$$ where $g_n$ is the Newtonian acceleration. For non-spherical geometry, this is only an approximation. However, in this Letter we only consider spherical systems as representing elliptical galaxies and we adopt $\mu(x) = x/\sqrt{1+x^2}$. This model is very successful on a galactic scale; in particular, it explains a large number of rotation curves of galaxies, with some exceptions (Gentile et al. 2004), and naturally the Tully-Fisher relation, (e.g. Sanders \& McGaugh 2002). While the dark matter problem is observationally very clear around spiral galaxies, thanks to their rotation curve measured with the cold hydrogen gas at 21cm, which is nearly in circular orbits (e.g. Bosma 1981, Verheijen \& Sancisi 2001), the situation is much more complex around elliptical galaxies, with little or no rotation. Recently, planetary nebulae have been used as an efficient tool for measuring the velocity field at large radii in early-type galaxies, and they complement stellar absorption kinematical studies (Romanowky et al. 2003). In typical elliptical galaxies, the velocity-dispersion profiles were found to decline with radius, up to 5 effective radii, thereby requiring no dark matter at all. Dekel et al. (2005) show that the data are still compatible with the usual dark matter models, if the planetary nebulae tracers have particularly radial orbits in the outer parts, because of a recent merger with small impact parameter. However, the recent results from Douglas et al (2007) challenge this interpretation, prolonging the decline to more than 7 effective radii. On larger scales around early-type galaxies, from 50 to 300kpc, Klypin \& Prada (2007), KP07, have proposed to test gravity models with satellite galaxies as a tracer. From the Sloan Digital Sky Survey, they stack several thousand galaxies in 3 luminosity classes and determine the number density of satellites and their velocity dispersion around them. In each mass range, the radial distributions are obtained with around 1500 satellites, although about 1.5 satellite exist around each galaxy. This large-scale galaxy neighborhood has not been widely tested yet in modified gravity. The well-known difficulty of MOND in clusters has found a possible solution with neutrinos of 2 eV mass (Sanders 2003; Angus et al 2007a), and the escape velocity around giant galaxies like the Milky Way was shown to correspond to observations, when including the external field effect (Famaey et al 2007; Wu et al 2007). In this work, we solve the Jeans equation for the distribution of the velocity dispersion around elliptical galaxies, and in particular we fit the NGC3379 galaxy, where the most extended data is available for the velocities. We also further explore fits at larger radii for the special case of NGC 3379, with satellite galaxies as tracers, as done by KP07 and Angus et al (2007b). This is only statistically valid around a generic early-type galaxy, with a mass comparable to NGC 3379. This is the brightest galaxy of a group, but the observed companion velocity is not statistically significant. \begin{figure}[ht!] \center{ \includegraphics[angle=-90, width=8cm]{8569-f1.ps} \caption{ Fit of the blue-light distribution in NGC 3379. The data points are taken from Douglas et al (2007). The curve represents the light profile derived with the mass-to-light ratio given in Table \ref{massp} and a mass model of cored-Hernquist type (Table \ref{tracers}). } \label{logsigma} } \end{figure} \vspace{-1 cm} | In contrast to the conclusion of KP07, MOND does not predict constant velocity dispersion with radius in the neighborhood of elliptical galaxies. There is a wide latitude for varying the possible anisotropy parameter according to the scale and the tracer considered. Our best fit starts from a tangential or isotropic configuration near the center and evolves progressively to a radial one for each of the three scales considered, which appears quite realistic. While the Newtonian model without dark matter has problems in the outer parts, the CDM1 model encounters severe difficulties in the inner parts. The CDM2 model, with a reduced dark matter relative to the visible mass, can also fit the data quite well. Compared to the MOND model, it requires a larger radial anisotropy on each scale, and its $\beta$ profile is unusual for the satellite tracer with a negative slope. For $\beta$= constant, the $\sigma_{los}^{sat}$ would be too steep compared to the observations. And by introducing anisotropy increasing with radius, the $\sigma_{los}^{sat}$ slope is even more increased. In the MOND regime, the external field effect (EFE, e.g. Wu et al 07) would also change the predictions. This will modulate the actual force on the particle tracers, and on the velocity. Since there was still latitude for fitting the observations with MOND, we feel that it is still possible to consider it with other anisotropy parameters. However, we note that the external field effect is not even necessary, for reproducing the observed velocity dispersion slope. \vspace{-0.2 cm} | 7 | 10 | 0710.4070 |
0710 | 0710.1180_arXiv.txt | \noindent We use a gravitational bar torque method to compare bar strengths (the maximum tangential force normalized by radial force) in $B$ and $H$-band images of 152 galaxies from the Ohio State University Bright Spiral Galaxy Survey. Our main motivation is to check how much the difference in the rest-frame wavelength could affect comparisons of bar strengths in low and high redshift observations. Between these two bands we find an average bar strength ratio \QBH= $1.25$ which factor is nearly independent of the morphological type. We show that \QBH $> 1$ is mostly due to reduced bulge dilution of radial forces in the $B$-band. The bar torque method needs an estimate for the vertical scale height of the galaxy, based on the radial scale length of the disk and the galaxy's morphological type. Since these two might not always be possible to determine at high redshifts in a reliable manner, we also checked that similar results are obtained with vertical scale heights estimated from the radii corresponding to the $K$-band surface brightness of 20 mag/arcsec$^2$. Also, we made a simple test of the usability of the bar torque method at high redshifts by checking the effects of image degradation (nearest neighbour sampling without any adjustment of noise levels): we found that the estimated bar strengths varied by $\pm 10\%$ at most as long as the total extent of the bar was at least 10 pixels. Overall, we show that the gravitational bar torque method should provide a proficient tool for quantifying bar strengths also at high redshifts. | Bars are important drivers of galaxy evolution, which highlights the importance of studying bar fractions and the properties of bars at a large range of redshifts. Bars generally have fairly old stellar populations, thus being well visible in the near-IR, where also the obscuration by dust is minimal. However, the near-IR images we obtain from the distant galaxies are redshifted so that at redshift $z \sim 1$ their rest-frame wavelength lies in the optical. Therefore, it is important to compare the optical and near-IR properties of bars in local galaxies. For that purpose the $B$ and $H$-band images are very convenient, offering a wide separation in wavelengths. Namely, bars are still prominent in the $B$-band \citep{osu,whyte,sheth06}, where they can be identified in a fairly reliable manner, whereas, due to the Balmer break, bars typically disappear in the ultraviolet \citep{sheth}. On the other hand, the $H$-band is already fairly free of the effects of dust and thus offers a reliable tracer for the stellar mass distribution. If bars were formed due to instabilities in the disk, one would expect a large fraction of barred galaxies at high redshifts in the framework of the hierarchical clustering models. This is because these models indicate dynamically colder disks in the distant Universe \citep*{NFW}. Also, interactions of disk galaxies were probably more frequent in the past \citep*[and references therein]{fer}, promoting the formation of bars \citep*[discussed in][]{elm}. Therefore, it was surprising to find a lack of bars in galaxies at redshift of $z>0.7$ (\citealp{abe}; \citealp{vdb96}; \citealp{vdb00}; \citealp*{vdb01}). Since this might be an artifact due to noise, poorer image resolution, and redshifted wavelength, \citet{vdb} degraded about 100 $B$-band images of the Ohio State University Bright Spiral Galaxy Survey ({\sc osubsgs}) to study these effects. By mimicking the characteristics of the images in the Hubble Deep Field, using the Wide Field Planetary Camera 2 on HST, they created artificial images of those galaxies as they would appear in the $I$-band at a redshift $z=0.7$. They found that two-thirds of bars initially classified as strong bars (SB), were still detectable in the degraded images, whereas weak bars (SAB) would largely disappear. They concluded that while selection effects reduce the number of observed bars, they cannot completely explain the lack of bars at high redshift. The early studies of bar frequencies were challenged by \citet{sheth}, who found similar bar frequencies at high and at low redshifts for very large bars. However, the number statistics were a severe concern, since only four bars were identified. These authors used galaxies in the Hubble Deep Field-North, observed at $V$, $I$ and $H$-bands, paying particular attention to the resolution limits due to which small bars cannot be detected at high redshifts. In two subsequent studies \citep{elm,jogee}, higher resolution images were used, based on the optical Advanced Camera for Surveys (ACS) at the Hubble Space Telescope. Both studies found a constant bar fraction in the redshift range $z = 0-1$. \citet{jogee} also found that it holds separately for the distinct intervals $z = 0-0.7$ and $z = 0.7-1.0$, using broad $V$ and $z$-bands to identify the bars. Taking into account the wavelengths used in their studies this constant bar fraction is not completely expected. Namely, if the true bar fraction is the same at all redshifts and bars are more prominent in the near-IR, one would expect a larger number of detected bars in the intermediate redshifts. There the rest-frame band shifting is not yet important, so that both strong and weak bars should be detected \citep{sheth06}. However, in the above studies the number statistics are still too poor (less than 300 galaxies spread over the entire redshift range) for any reliable discussion of how the bar fraction evolves as a function of redshift. Also, using ACS very deep images are required for detecting weak bars at high redshift. The above studies have been superseded recently by \citet{sheth07}, using wide field {\sc cosmos} observations in $I$-band at $z < 0.84$, based on deeper images than used in the previous studies. Contrary to most previous studies they found that the fraction of bars declines rapidly with redshift. This drop of bars was also found to be more dramatic for strong bars. These studies of bar frequencies clearly show, besides the difficulty of measuring bar frequencies at high redshifts, also the importance of measuring the strengths of the bars at high redshifts in a reliable manner. In this study we use the {\sc osubsgs} to compare bar strengths derived from the $B$ and $H$-band images of nearby galaxies, using a gravitational bar torque method. Besides bar strength, our method gives simultaneously also the length and the relative luminosity of the bar. The outline of this paper is as follows. In Sections \ref{sample} and \ref{method} we describe the details of our sample and the method used, and also address the factors which might affect the bar strength measurements at higher redshifts. We then show comparisons between $B$ and $H$-band bar strengths in Section \ref{highz}; the results are discussed in Section \ref{discussion}, and summarized in Section \ref{conclusions}. | We have used a gravitational bar torque method to compare bar strengths derived from the $B$ and $H$-band images for the galaxies in the {\sc osubsgs} sample. We presented four different ways to do the comparison. Further we subjected the method to various tests simulating a high redshift environment. We estimate $z_0$ through $r_{K_{20}}$, and degrade the $B$-band images using a nearest neighbour algorithm before calculating the bar strengths. Our main conclusions can be summarized as follows: (i) Quite unexpectedly, we found that bars appear to have stronger torques in the optical than in the near-IR: the conversion factor between $B$ and $H$-bands is 1.25. We also showed that this is due to reduced dilution of radial forces by relatively smaller bulges in the optical. (ii) The scaled bar lengths, $barlen$/$r_{K_{20}}$, and $r_{Q_g}/barlen$ are similar in the $B$ and $H$-bands, with the bars in the $B$-band being on average slightly shorter than bars in the $H$-band. (iii) An approximation of the vertical scale height while calculating $Q_g$ can be estimated from $r_{K_{20}}$. Resulting bar strengths will be overestimated for early types and underestimated for late types, by less than 15\%. (iv) The $Q_g$ method is found to work well even on low resolution images. We degraded our $B$-band images using the nearest neighbour algorithm and find that as long as the total bar length is at least $\sim 10$ pixels, the resulting $Q_g$ value is typically within 10\% of the original one. We have shown that the $Q_g$ method provides a reliable tool for measuring bar strengths at high redshifts: with the pixel resolution of $0\farcs 05$ of the ACS, bars with $\sim 2-3$ kpc can still be reliably measured at $z \sim 1.0$. | 7 | 10 | 0710.1180 |
0710 | 0710.2847_arXiv.txt | {M\,87 is the first extragalactic source detected in the TeV range that is not a blazar. With the increasing performances of ground-based \v{C}erenkov telescopes, we can now probe the variability in the $\gamma$-ray flux at small timescales, thus putting strong constraints on the size of the emitting zone. The large scale jet of M\,87 is misaligned with respect to the line of sight. A modification of standard emission models of TeV blazars appears necessary to account for the $\gamma$-ray observations despite this misalignment. } {We explain TeV $\gamma$-ray spectra and fast variability of M\,87 by invoking an emission zone close to the central supermassive black hole, which is filled with several plasma blobs moving in the large opening angle of the jet formation zone. } {We develop a new multi-blob synchrotron self-Compton (SSC) model with emitting blobs set on a cap beyond the Alfv{\'e}n surface in the jet, at a distance of $\sim$100\,$r_g$ from the central engine to interpret the high energies inferred by new TeV observations. We present a SSC model that is explicitly adapted to deal with large viewing angles and moderate values of the Lorentz factor inferred from (general relativistic) magnetohydrodynamic models of jet formation. } {This scenario can account for the recent $\gamma$-ray observations of M\,87 made by the High Energy Stereoscopic System (H.E.S.S.) telescope array. We find individual blob radii of about $10^{14}$\,cm, which is compatible with the variability on timescales of days recently reported by the H.E.S.S. collaboration and is of the order of the black hole gravitational radius. Predictions of the very high energy emission for three other sources with extended optical or X-ray jet which could be misaligned blazars still with moderate beaming are presented for one Seyfert 2 radiogalaxy, namely Cen\,A, one peculiar BL\,Lac, PKS\,0521$-$36, and one quasar, 3C\,273. } {} | M\,87 is a well-known nearby giant elliptical galaxy \citep[z=0.00436,][]{2000MNRAS.313..469S} close to the center of the Virgo cluster, which shows a multi-spectral jet, signature of an active galactic nucleus (AGN). Its jet is one of the best known, at all scales, thanks to its nearby location and its strong synchrotron radiation in the optical band. M\,87 is classified as FR\,I based on its radio morphology. \citet{2002ApJ...568..133W} observed the jet with {\it Chandra} on July 29 and 30, 2000 and detected it up to a distance of $\sim$21$\arcsec$ from the core in the X-ray band, which implies that the jet is not as strongly aligned along the line of sight \citep[see also][]{1996MNRAS.283..873R} as in the case of blazars. At radio wavelengths, an impressive jet, which extends up to a few tens of kiloparsecs, can be seen. The central engine is thought to be a supermassive black hole (SMBH) with a mass of $M_\mathrm{BH} \sim 3 \times 10^9 M_{\sun}$ \citep{1997ApJ...489..579M}. The scale length is thus $r_g = G M_\mathrm{BH} / c^2 \sim 4.5 \times 10^{14}$\,cm $\sim 1.4 \times 10^{-4}$\,pc. Using the {\it Hubble Space Telescope} ({\it HST}), \citet{1999ApJ...520..621B} observed superluminal apparent motions of about $4c$--$6c$ beyond 400\,pc for the internal knots, between 1994 and 1998, thus confirming that the jet is relativistic. They conclude that the jet is oriented within $19 \degr$ of the line of sight. Due to the presence of a SMBH in the core and the presence of the jet, M\,87 was deemed an interesting candidate for TeV emission. \citet{2004ApJ...610..156L} reported an upper limit observed with Whipple in 2000 and 2001, simultaneously with X-ray flares observed by {\it RXTE}. HEGRA observed M\,87 in 1998 and 1999 for a total exposure of 77\,h after data quality selection \citep{2003A&A...403L...1A,2004NewAR..48..407B}. A 4.1$\sigma$ significance was recorded and an integrated flux ($E>730$\,GeV) of 3.3\% Crab was measured. Recently, \citet{2006Sci...314.1424A} have observed M\,87 with the High Energy Stereoscopic System (H.E.S.S.)\footnote{\url{http://www.mpi-hd.mpg.de/hfm/HESS/}} between 2003 and 2006 in 89\,h live-time with a $13\sigma$ detection and discovered variations on timescales of about 2 days, 10 times faster than that observed in any other waveband. This shows that the emission region is very compact, with a dimension of the order of a few Schwarzschild radii. These observations, thus confirming the detection by HEGRA \citep[][]{2005astro.ph..4395B}, are particularly interesting since M\,87 is the first non-BL\,Lac extragalactic object ever observed at TeV energy. Radio-loud galaxies contain AGNs with jets like blazars, but the jet emission is less strongly boosted due to larger viewing angles between the jet and the observer's line of sight. It is therefore a challenge for standard models of TeV blazars to explain the very high energy (VHE) emission of M\,87. In this paper, we present a modified synchrotron self-Compton (SSC) scenario to explain the VHE emission of M\,87. Classic SSC models \citep[e.g.][]{1979A&A....76..306G,1996ApJ...463..555I,1996ApJ...461..657B,1999MNRAS.306..551C,2001A&A...367..809K} are applied to blazars, which are beamed sources, and cannot account for the observations of radiogalaxies like M\,87. Our goal is to further develop one of these models to reconcile beamed and unbeamed sources in the same framework of models. Such propositions for unification of AGNs have already been studied considering orientation effects \citep[e.g.][]{1993ARA&A..31..473A,1995PASP..107..803U}, or radio/X-ray power among BL\,Lac objects, flat-spectrum radio-loud quasars (FSRQs) and FR\,Is \citep[e.g.][]{1998MNRAS.299..433F,1998MNRAS.301..451G,2000MNRAS.318..493C}. A short description of the leptonic blob-in-jet model for TeV blazars is found in Sect.~\ref{sec:bij}, and its development and application to M\,87 are described in Sect.~\ref{sec:multi_blobs} and Sect.~\ref{sec:M87}. In the framework of misaligned BL\,Lac-like objects, we then try to predict VHE fluxes for objects with optical/X-ray extended jets in Sect.~\ref{sec:predictions}. Implications on unification schemes of AGNs are discussed in Sect.~\ref{sec:unification}. | } We have presented a SSC model to interpret VHE emission of M\,87 as well as other misaligned sources with extended optical/X-ray jet. This model accounts in a simple way for a differential Doppler boosting effect by modeling the emission of several blobs of plasma located in the broadened formation zone of the jet close to the SMBH, just beyond the Alfv{\'e}n surface predicted by MHD models. Our scenario provides a reasonable interpretation of the H.E.S.S. VHE observations of M\,87 and provides the possibility to extend standard leptonic models of TeV blazars to other types of AGNs. However, we do not exclude other leptonic or hadronic models. For instance, \citet{2007arXiv0704.3282N} recently interpreted the H.E.S.S. observations from 2005 of M\,87 by invoking acceleration and radiation of electrons in the black hole magnetosphere, which is another kind of leptonic model. Hadronic models also cannot be excluded as efficient particle acceleration processes can occur in the close surroundings of the black hole. More observations are needed to constrain the emission models and especially to distinguish between hadronic and leptonic scenarii. The upcoming {\it GLAST} mission and the H.E.S.S.\,II project will certainly help to understand the mechanisms at work in the AGNs by exploring spectral ranges below TeV, which is decisive to constrain the shape of the inverse Compton bump. Moreover, the lack of genuine simultaneous multiwavelength campaigns on M\,87 needs to be filled, especially since this source is known to be variable at small timescales in VHE. Several types of active nuclei are potential emitters of VHE photons with predicted TeV fluxes detectable by present \v{C}erenkov arrays like H.E.S.S. and MAGIC, or by the next generation of instruments such as the CTA project. Such data will be crucial to test AGN unifying schemes. | 7 | 10 | 0710.2847 |
0710 | 0710.0590_arXiv.txt | We use the magnetic butterfly diagram to determine the speed of the magnetic flux transport on the solar surface towards the poles. The manifestation of the flux transport is clearly visible as elongated structures extended from the sunspot belt to the polar regions. The slopes of these structures are measured and interpreted as meridional magnetic flux transport speed. Comparison with the time-distance helioseismology measurements of the mean speed of the meridional flows at the depth of 3.5--12~Mm shows a generally good agreement, but the speeds of the flux transport and the meridional flow are significantly different in areas occupied by the magnetic field. The local circulation flows around active regions, especially the strong equatorward flows on the equatorial side of active regions affect the mean velocity profile derived by helioseismology, but do not influence the magnetic flux transport. The results show that the mean longitudinally averaged meridional flow measurements by helioseismology may not be used directly in solar dynamo models for describing the magnetic flux transport, and that it is necessary to take into account the longitudinal structure of these flows. | The largest-scale velocity fields on the Sun consist of differential rotation and meridional circulation. The differential rotation is defined as an integral of the zonal (East--West) component of the velocity field depending on the solar latitude, $b$, and radius. The meridional flow is calculated as an integral of the North--South component of the velocity field, generally depending again on the latitude and radius. Both the differential rotation and meridional circulation are the key ingredients of the solar dynamo. The differential rotation plays an important role in generating and strengthening of toroidal magnetic field inside the Sun, while the meridional flow transports the magnetic flux towards the solar poles resulting in cyclic polar field reversals \citep[for a recent review, see][]{Brandenburg2005}. The meridional flux transport seems to be an essential agent influencing the length, strength and other properties of solar magnetic cycles. Generally, the slower the meridional flows are, the longer the next magnetic cycle is expected. Dynamo models showed that the turn-around time of the meridional cell is between 17 and 21 years, and that the global dynamo may have some kind of memory lasting longer than one cycle \citep{2006GeoRL..3305102D}. The speed of the meridional flow and its variation with the solar cycle measured by local helioseismology in the subsurface layers of the Sun were used as an input in the recent flux-transport models \citep{2006ApJ...649..498D}. In local helioseismology measurements (e.g. \citeauthor{2004ApJ...603..776Z} \citeyear{2004ApJ...603..776Z}, \citeauthor{2006ApJ...638..576G} \citeyear{2006ApJ...638..576G}), the meridional flow was derived from a general subsurface flow field by averaging the North-South component of the plasma velocity over longitude for a Carrington rotation period. The studies revealed that the mean meridional flow varied with the solar activity cycle. These variations may significantly affect solar-cycle predictions based on the solar dynamo models, which assume that the magnetic flux is transported with the mean meridional flow speed \citep{2006ApJ...649..498D}. Our goal is to verify this assumption and to investigate the relationship between the subsurface meridional flows and the flux transport. In this study, we show that the mean meridional flows derived from the time-distance helioseismology subsurface flow maps are affected by strong local flows around active regions in the activity belts. These local flows have much less significant effect on the magnetic flux transport. | We have compared the measurements of the meridional speed derived from two different techniques: by time-distance local helioseismology and by measuring the flux transport speed using the magnetic butterfly diagram. We have found that the results agree quite well in general, but they differ in regions occupied by local magnetic fields. The detailed flow maps from helioseismology show that this is partly due to the presence of meridional counter-cells at the equatorial side of magnetic regions, which influence the time-distance derived meridional flow profile, but does not influence the magnetic flux transport. We have studied eleven non-consecutive Carrington rotations covering one solar cycle. The effect of the local flows around active regions and especially on their equatorial side is noticed in all the studied cases. Therefore, this behaviour seems to be a common property of the subsurface dynamics around active regions located in the activity belt. However, we have to keep in mind that both datasets are not directly comparable, since the time-distance flow maps represent the behaviour of flows during one Carrington rotation, while the butterfly diagram tracking procedure provide results averaged over few Carrington rotations. The more homogeneous data for local helioseismology are needed to study this effect in more detail. The results show that the speed of the magnetic flux transport towards the solar poles may significantly deviate from the longitudinally averaged meridional flow speed derived from local helioseismology measurements, which are affected by local circulation flows around active regions in the activity belt. Therefore, using the longitudinally averaged meridional flow profile from helioseismology in the solar cycle models for description the flux transport is not justified. The longitudinal structure of these flows should be taken into account. | 7 | 10 | 0710.0590 |
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