{ "1101/1101.3240_arXiv.txt": { "abstract": "We investigate the validity of the generalized second law of gravitational thermodynamics on the apparent and event horizons in a non-flat FRW universe containing the interacting dark energy with dark matter. We show that for the dynamical apparent horizon, the generalized second law is always satisfied throughout the history of the universe for any spatial curvature and it is independent of the equation of state parameter of the interacting dark energy model. Whereas for the cosmological event horizon, the validity of the generalized second law depends on the equation of state parameter of the model. ", "introduction": "The present acceleration of the universe expansion has been well established through numerous and complementary cosmological observations \\cite{Riess}. A component which is responsible for this accelerated expansion usually dubbed ``dark energy'' (DE). However, the nature of DE is still unknown, and people have proposed some candidates to describe it (for review see \\cite{Padmanabhan,Copeland} and references therein). One of the important questions in cosmology concerns the thermodynamical behavior of the accelerated expanding universe driven by DE. It was shown that the Einstein equation can be derived from the first law of thermodynamics by assuming the proportionality of entropy and the horizon area \\cite{Hawking,Bardeen,Jacobson}. In the cosmological context, attempts to disclose the connection between Einstein gravity and thermodynamics were carried out. It was shown that the differential form of the Friedmann equation in the Friedmann-Robertson-Walker (FRW) universe can be written in the form of the first law of thermodynamics on the apparent horizon \\cite{Cai05}. Further studies on the equivalence between the first law of thermodynamics and Friedmann equation has been investigated in various gravity theories like Gauss-Bonnet, Lovelock and braneworld scenarios \\cite{Cai05,Akbar,Sheykhi2}. Besides examining the validity of the thermodynamical interpretation of gravity by expressing the gravitational field equations into the first law of thermodynamics in different spacetimes, it is also of great interest to investigate the validity of the generalized second law (GSL) of thermodynamics in the accelerating universe driven by DE. The GSL of thermodynamics is as important as the first law, governing the development of the nature \\cite{Davies,Gong,Izquierdo2,Izquierdo,Zhau,Sheykhi3,Jamil,Wang2,Karami2,Karami3}. Here our aim is to investigate the validity of the GSL of gravitational thermodynamics for the interacting DE model with DM in a non-flat FRW universe enclosed by the dynamical apparent horizon and the cosmological event horizon. Note that in the literature, people usually have studied the validity of the GSL for a specific model of DE with special kind of interaction term. But we would like to extend it to any DE model with general interaction term. This paper is organized as follows. In section \\ref{II}, we study the DE model in a non-flat FRW universe which is in interaction with the DM. In section \\ref{III}, we investigate the validity of the GSL of gravitational thermodynamics for the universe enclosed by the apparent horizon and the event horizon which is in thermal equilibrium with the Hawking radiation of the horizon. Also we give an example for the case of the cosmological event horizon. In section \\ref{IV} we investigate the effect of interaction term between DE and DM on the GSL. Section \\ref{V} is devoted to conclusions. ", "conclusions": "\\label{V} Here the GSL of gravitational thermodynamics for the interacting DE with DM in a non-flat FRW universe is investigated. Some experimental data have implied that our universe is not a perfectly flat universe and it possess a small positive curvature \\cite{Bennett}. Although it is believed that our universe is flat, a contribution to the Friedmann equation from spatial curvature is still possible if the number of e-foldings is not very large \\cite{Huang}. The boundary of the universe is assumed to be enveloped by the dynamical apparent horizon and the cosmological event horizon. We assumed that the universe to be in thermal equilibrium with the Hawking temperature on the horizon. We found that for the dynamical apparent horizon, the GSL is respected throughout the history of the universe for any spatial curvature and it is independent of the EoS parameter of the interacting DE model. But for the cosmological event horizon, the GSL is satisfied for the special range of the EoS parameter of the model. The above results show that the dynamical apparent horizon in comparison with the cosmological event horizon, is a good boundary for studying cosmology, since on the apparent horizon there is the well known correspondence between the first law of thermodynamics and the Einstein equation \\cite{Gong5}. In the other words, the Friedmann equations describe local properties of spacetimes and the apparent horizon is determined locally, while the cosmological event horizon, Eq. (\\ref{eh}), is determined by global properties of spacetimes \\cite{Cai05}. Besides in the dynamic spacetime, the horizon thermodynamics is not as simple as that of the static spacetime. The event horizon and apparent horizon are in general different surfaces. The definition of thermodynamical quantities on the cosmological event horizon in the nonstatic universe are probably ill-defined \\cite{Wang2}. \\\\ \\\\ \\noindent{{\\bf Acknowledgements}}\\\\ The authors thank the reviewers for very valuable comments. The work of K. Karami has been supported financially by Research Institute for Astronomy $\\&$ Astrophysics of Maragha (RIAAM), Maragha, Iran." }, "1101/1101.1073_arXiv.txt": { "abstract": "We examine various implications from a dynamical and chemical model of globular clusters (GCs), which successfully reproduces the observed abundance patterns and the multiple populations of stars in these systems assuming chemical enrichment from fast rotating massive stars. Using the model of Decressin et al.\\ (2007) we determine the ratio between the observed, present-day mass of globular clusters and their initial stellar mass as a function of the stellar initial mass function (IMF). We also compute the mass of low-mass stars ejected, and the amount of hydrogen ionising photons emitted by the proto globular clusters. Typically, we find that the initial masses of GCs must be $\\sim$ 8--10 times (or up to 25 times, if second generation stars also escape from GCs) larger than the present-day stellar mass. The present-day Galactic GC population must then have contributed to approximately 5--8\\% (10--20\\%) of the low-mass stars in the Galactic halo. We also show that the detection of second generation stars in the Galactic halo, recently announced by different groups, provides a new constraint on the GC initial mass function (GCIMF). These observations appear to rule out a power-law GCIMF, whereas they are compatible with a log-normal one. Finally, the high initial masses also imply that GCs must have emitted a large amount of ionising photons in the early Universe. Our results reopen the question on the initial mass function of GCs, and reinforce earlier conclusions that old GCs could have represented a significant contribution to reionise the inter-galactic medium at high redshift. ", "introduction": "\\label{s_intro} Although for a long time thought to be among the most simple stellar systems, globular clusters (hereafter GCs) have been subject to intense studies both observationally and through theory and simulations. These include for example detailed work on the stellar content of GCs and on chemical abundances of GC stars,searches for viable proto-GCs, studies of dynamical effects on massive star cluster evolution, cosmological simulations of their formation, estimates of their contribution to cosmic reionisation, and other related Galactic and extragalactic astrophysical topics \\citep[see e.g.\\ reviews by][]{Gratton04,Brodie06,Piotto2009,Boily10,Elmegreen10iau266}. Despite these studies many open questions remain, concerning both GCs as individual objects and as a collective population. For example, the origin of Galactic halo stars and the contribution of GCs to this population, are still unclear \\citep[see e.g.][]{Hut92,Parmentier07,Bell08,Boley09}. Similarly, the shape of the globular cluster initial mass function (GCIMF), the nature of the present-day globular cluster mass function, and the processes and the timescales responsible for transforming the former into the latter, are debated \\citep[see][]{Fall01,Vesperini03,Parmentier05,Parmentier07,Elmegreen10}. Also, first steps are being made in order to understand GC formation in cosmological simulations \\citep{Bromm02,Kravtsov05,Boley09,Griffen10}. Finally \\citet{Ricotti02} has shown that GCs emit enough ionising photons to reionise the Universe, provided their escape fraction $f_{\\rm esc}$ is of order unity. Examining this question is also of interest in the present context, where it appears that galaxies found so far in deep surveys are insufficient to reionise the inter-galactic medium \\citep[see e.g.][]{Bunker09,Ouchi09,McLure10,Labbe10}, and where the main sources responsible of cosmic reionisation, presumably faint, low-mass galaxies, below the current detection limits \\citep[cf.][]{Choudhury07,Choudhury08}, remain thus to be identified. A major paradigm-shift has occurred recently in the GC community, that sheds new light on these key questions. Indeed, detailed abundance studies of their long-lived low-mass stars made possible with 8-10m class telescopes, together with high-precision photometry of Galactic GCs performed with HST, have revolutionised our picture of these stellar systems. It is now clear that individual GCs host multiple stellar populations, as shown by their different chemical properties and by multimodal sequences in the colour-magnitude diagrams \\citep{Bedin04,Piottoetal07,Milone08,Milone10,VillanovaPiotto10}. Indeed, although nearly all GCs\\footnote{With the notable exception of $\\omega$ Cen, M22, and M54 \\citep[see e.g.][ and references therein]{DaCosta_M22_09,JohnsonPilachowski10,Siegel07,Carretta10M54}.} appear to be fairly homogeneous in heavy elements \\citep[i.e., Fe-peak, neutron-capture, and alpha-elements, see e.g.][]{James04,Sneden05,Carrettaetal09Fe}, they all exhibit large star-to-star abundance variations for light elements from C to Al that are the signatures of hydrogen-burning at high temperature implanted at birth in their long-lived low-mass stars \\citep[see e.g.][ and reference therein]{Gratton01,Gratton04,Sneden05,PrantzosCCCI07,Carretta09,Charbonnel10}. In fact, the so-called O--Na anticorrelation is ubiquitous in Galactic GCs, and is now accepted as the decisive observational criterion distinguishing {\\it bona fide} GCs from other clusters \\citep{Carretta10}. The current explanation for these chemical patterns is the so-called ``self-enrichment'' scenario that calls for the formation of at least two stellar generations in all GCs during their infancy: The first generation stars were born with the proto-cluster original composition, which is that of contemporary field halo stars, while the second generation stars formed from original gas polluted to various degrees by hydrogen-burning processed material ejected by more massive, short-lived, first generation GC stars. Details and references can be found e.g. in \\citet{PrantzosCC06}, who discuss the pros and cons of two versions of this ``self-enrichment scenario\", invoking either massive AGB stars \\citep[e.g.][]{CottrellDaCosta1981,VenturaDAntona2001,VenturaDAntona09}, or fast rotating massive stars \\citep[e.g.][]{Decressin07b} as polluters. Whatever the actual polluting stars, an immediate consequence of this scenario is that, in order to reproduce the present proportion of first to second generation stars with acceptable values of the polluters IMF, the initial stellar masses of GCs must have been considerably larger than their present-day value \\citep{PrantzosCC06,Decressin07b,Dercole08,Decressin10,Carretta10}. However, most of the extra-galactic studies have not yet incorporated this revised picture, or have not yet explored the resulting implications. Furthermore, the recent discovery of stars with signatures characteristic of 2$^{\\rm nd}$ generation GC stars among the metal-poor halo population \\citep{Carretta10,Martell10} sheds new light on the amount of low-mass stars ejected from GCs and on the initial mass function of these clusters, as we shall show below. In the present paper, we explore several consequences of this new paradigm, based on the model that was developed by \\citet[][ hereafter DCM07]{Decressin07b} to describe the early chemical and dynamical evolution of GCs. In this model fast-rotating massive ($M \\ga$ 25 \\msun) stars are responsible for the GC pollution. The model successfully explains the observed abundance patterns of present-day GC stars, and has also been tested with N-body and hydrodynamical simulations \\citep[see][]{Decressin08,Decressin10}. Its main assumptions are briefly described and summarised in \\S \\ref{s_model}. Within this framework, we constrain the relation between the initial and the present stellar mass of GCs (\\S \\ref{initialmass}), as well as the contribution to the stellar halo (\\S \\ref{contribhalo}), taking the recent observational identification of second generation stars in the Galactic halo \\citep{Martell10, Carretta10} into account. Implications on the GCIMF are derived in \\S \\ref{GCIMF}. Finally, we derive in \\S \\ref{GCreionisation} a well-defined ionising photon production rate for proto-GCs, taking all the detailed observational constraints from nearby GCs into account, and estimate their contribution to cosmic reionisation. Our main conclusions are summarised in \\S \\ref{s_conclude}. ", "conclusions": "\\label{s_conclude} In light of the recently recognised, general existence of multiple stellar generations in globular clusters (GCs) implying significant losses of 1$^{\\rm st}$ generation stars from these clusters, we have re-examined the initial masses of GCs, the contribution of low-mass stars ejected from GCs to the stellar halo of our Galaxy, and the contribution of GCs to the ionising photon production necessary to reionise the inter-galactic medium at high redshift. These quantities have been estimated from the chemical and dynamical model of \\citet{Decressin07b}, which successfully reproduces the main observational constraints from 1$^{\\rm st}$ and 2$^{\\rm nd}$ generation stars, by invoking pollution from fast rotating massive stars. The main free parameters of this model are the slope of IMF for high masses ($>$ 0.8 \\msun, the IMF being fixed to the observed log-normal distribution for lower masses), the relative number of 1$^{\\rm st}$/2$^{\\rm nd}$ generations stars, given by the fraction $f_p=0.33 ^{+0.07}_{-0.08}$ of 1$^{\\rm st}$ generation stars determined from the detailed spectroscopic observations of \\citet{Carretta10}, and a dilution parameter $d\\approx 1.15$ inferred from the Li-Na anticorrelation observed in GCs \\citep{Decressin07b,Charbonnel10b}. The dynamical scenario we have explored allows for the evaporation of stars from the 1$^{\\rm st}$ generation (corresponding to an escape fraction of 2$^{\\rm nd}$ generation stars of zero, $\\esll=0$), or from both generations, as suggested by recent observations finding stars characteristic of the 2$^{\\rm nd}$ generation in GCs in the Milky Way halo \\citep{Carretta10,Martell10}. The latter case translates to $\\esll \\sim$ 0.43--0.65. We have obtained the following main results for an IMF with a Salpeter slope above 0.8 \\msun: \\begin{itemize} \\item The initial stellar masses of GCs must have been $\\sim$ 8--10 times larger than the current (observed) mass, when no second generation stars are lost, in agreement with the earlier results of \\citet{PrantzosCC06} and \\citet{Carretta10}. If all 2$^{\\rm nd}$ generation halo stars originate from the present population of GCs, the initial cluster masses must have been $\\approx 25$ times larger than the current mass. \\item The mass in low-mass stars ejected from GCs must be $\\sim 2.5-3.2$ times their observed, stellar mass if all 2$^{\\rm nd}$ generation stars were retained, or $\\sim$ 5--10 times the present day mass if $\\esll \\sim$ 0.43--0.65. These numbers translate to a contribution of 5--8\\% or 10--20\\% respectively of the ejected low-mass stars to the Galactic stellar halo mass. We have compared our estimate with earlier values obtained from various methods (cf.\\ \\S \\ref{contribhalo}). \\item The observations of 2$^{\\rm nd}$ generation stars in the Galactic halo can constrain the initial mass function of the GC population (GCIMF). In particular we have shown that a power-law with a slope $\\beta \\approx -2$, as often assumed, is in contradiction with recent determinations of the fraction of 2$^{\\rm nd}$ generation stars in the halo, whereas a log-normal GCIMF is compatible with these observations. This finding revives the question about a common mass function and about the physical processes leading to a distinction between globular clusters with multiple stellar populations and other clusters. \\item Due to their high initial masses, the amount of Lyman continuum photons emitted by GCs during their youth must have been substantial. Indeed, we find that their output corresponds to a total number of ionising photons emitted per baryon, $\\eta^\\prime \\approx 10^{4.8-4.9}$ for $\\esll=0$, or $\\sim$ 1.7--3.5 times more if $\\esll \\sim$ 0.4-0.6. Our results reinforce the conclusion of \\citet{Ricotti02} that GCs should contribute significantly to reionise the IGM at very high redshift ($z \\ga 6$). Individual, young proto-GCs with typical masses few times $\\sim 10^6 \\msun$ could just be detectable at high redshift in ultra-deep images with the HST, and are certainly within the reach of the JWST. \\end{itemize} The dependence of the initial and ejected masses on the IMF slope has been illustrated in Fig.\\ \\ref{fig_m}. The ionising photon production is found to be quite insensitive to the high mass IMF, since both the ejecta ``polluting'' the 2$^{\\rm nd}$ generation stars and the Lyman continuum flux originates from massive stars. Our main results should also be valid for the massive AGB scenario, at least qualitatively." }, "1101/1101.1590_arXiv.txt": { "abstract": "Direct imaging of the HR8799 system was a major achievement in the study of exoplanets. HR8799 is a $\\gamma$\\,Doradus variable and asteroseismology can provide an independent constraint on the inclination. Using 650 high signal-to-noise, high resolution, full visual wavelength spectroscopic observations obtained over two weeks at Observatoire de Haute Provence (OHP) with the SOPHIE spectrograph we find that the main frequency in the radial velocity data is 1.9875 d$^{-1}$. This frequency corresponds to the main frequency as found in previous photometric observations. Using the FAMIAS software to identify the pulsation modes, we find this frequency is a prograde $\\ell$=1 sectoral mode and obtain the constraint that inclination $i\\gtrsim$40$^{\\circ}$. ", "introduction": "The imaging discovery of the three \\citep{Ma08}, and now four \\citep{Mr10} planets around HR~8799 is a significant achievement in the search for and study of planets orbiting other stars. For the first time, the thermal emission of planets in orbit around another star has been unambiguously detected. The dynamical evolution of a planetary system is complex. From the basic planet formation assumption that planets form by the core accretion or disk instability scenario in a disk along the star's equatorial plane, systems can suffer drastic changes; planet-planet perturbations, interactions with a disk or stellar encounters can change a planet's orbital inclination, its semi-major axis, its orbital eccentricity or even eject it \\citep{Ra10}. In the case of HR~8799, it is not impossible that a close encounter occurred prior to planet formation, since its relatively high galactic velocity compared to the Columba association and its far distance away from the other association members \\citep{Hi10} are suggesting that it may have been kicked out and is probably moving quickly away from its birth place. Such an encounter could have tilted a disk relative to the star's equatorial plane and induce perturbations that may have led to planet formation where planets would have a non-negligible orbital inclination relative to the star. Radial velocity searches have confirmed such chaotic behaviors by detecting systems where planets are orbiting well away from the star's equatorial plane (\\citealt{Tr10} and references therein), although for close-in extrasolar planets this misalignment could also be caused by the Kozai mechanism \\citep[e.g.][]{Wu07,Fa07,Wi09}. For the wide HR8799 planets the Kozai mechanism is not operational (even if the system had a stellar companion in the past); finding a misalignment between the star and the planet's orbital plane would be a sign of a significant dynamical interaction in the system past. \\citet{Ma08} have suggested that the HR~8799 planets are in a similar orbital plane with a low inclination and have mostly circular orbits. This is because the detected orbital motions are close to their expected face-on circular orbit values, the orbital motion is mainly in azimuth and the star is known to be a slow rotator (thus it would be viewed mainly pole-on). Dynamical analyses \\citep[e.g. ][]{Re09,Fa10,Mo10,Ma10} have confirmed that the planets are mostly in the same plane with small eccentricities, although such fits are still very uncertain due to the limited amount of orbital coverage available. In addition, the planets also all orbit in the same counter-clockwise orientation, further supporting the idea that they formed in a disk, similar to the Solar system planets. Assuming a circular orbit for b, \\citet{La09} found $i\\sim$13$^\\circ$--23$^\\circ$; while attempting a coherent analysis of various portions of observational data on known components of the system, \\citet{Re09} concluded that $i$ should range between 20$^\\circ$ and 30$^\\circ$. Also, Spitzer observations of HR~8799's complex debris disk suggest that any inclination angle larger than $\\sim$25$^\\circ$ should be excluded \\citep{Su09}. Using a statistical distribution of star's rotation speed \\citep{Ro07}, HR~8799 with its 37.5$\\pm$2 km\\,s$^{-1}$ V\\,sin\\,$i$ \\citep{Ka98} would be consistent with an inclination of $\\sim 23.5^\\circ$ if it is an A5 star or $\\sim 18.5^\\circ$ if it is an F0 star (HR~8799 spectral classification is uncertain mainly due to its low metallicity that is affecting it's broad band colors). Such a determination is of course statistical and a direct star's inclination determination is required for a meaningful comparison with the estimated planet orbital plane inclination. HR~8799 is an intrinsic photometric and spectroscopic variable \\citep{Ro95,Ma04}. It has also been confirmed as a $\\gamma$\\,Doradus (Dor) variable \\citep{Ze99}. The $\\gamma$\\,Dor stars are late A to early F stars whose pulsations are driven by a flux-blocking mechanism at the base of their convective envelope \\citep[e.g.][]{Du04}. The $\\gamma$\\,Dor nature of HR~8799 offers a unique opportunity to estimate its inclination via an asteroseismic analysis of the observed $g$-modes. In a previous asteroseismic analysis using photometric frequencies \\citet{My10} has shown that an age determination for the system, which allows to discriminate between planets and brown dwarfs, is a difficult task with the current information and discussed the importance of an inclination determination since the equatorial rotational velocity can be used in constraining the age of the system. In this letter, we present a spectroscopic asteroseismic analysis and also obtain limits for HR~8799's stellar inclination. The data have been acquired from an extensive multi-site ground-based high-resolution spectroscopy campaign. Section~2 describes that data used in this letter. Section~3 discusses the pulsation mode identification and the determination of the stellar inclination and Section~4 describes the conclusions. ", "conclusions": "We conclude that the stellar rotational inclination axis has a value i$\\gtrsim$40$^\\circ$ based on identification of the 1.98\\,d$^{-1}$ frequency as an $\\ell$=1 $m$=1 mode. This is the strongest pulsation in both photometry \\citep{Ze99,Cu06} and spectroscopic radial velocities. Through dynamical analyses it is suspected that the planets are mostly in the same plane with small eccentricities and that the planets orbit inclination axis is $\\sim$20$^\\circ \\pm$10$^\\circ$ \\citep{Re09,La09}. The current data suggests a misalignment of $\\Delta i \\gtrsim$20$^\\circ$ between the stellar rotational inclination and planetary orbit axes, though more detailed pulsational analyses and better orbital fits are needed before this can be confirmed." }, "1101/1101.1559_arXiv.txt": { "abstract": "The Sloan Digital Sky Survey (SDSS) started a new phase in August 2008, with new instrumentation and new surveys focused on Galactic structure and chemical evolution, measurements of the baryon oscillation feature in the clustering of galaxies and the quasar Ly$\\alpha$ forest, and a radial velocity search for planets around $\\sim 8000$ stars. This paper describes the first data release of SDSS-III (and the eighth counting from the beginning of the SDSS). The release includes five-band imaging of roughly 5200 deg$^2$ in the Southern Galactic Cap, bringing the total footprint of the SDSS imaging to 14,555 deg$^2$, or over a third of the Celestial Sphere. All the imaging data have been reprocessed with an improved sky-subtraction algorithm and a final, self-consistent photometric recalibration and flat-field determination. This release also includes all data from the second phase of the Sloan Extension for Galactic Understanding and Exploration (SEGUE-2), consisting of spectroscopy of approximately 118,000 stars at both high and low Galactic latitudes. All the more than half a million stellar spectra obtained with the SDSS spectrograph have been reprocessed through an improved stellar parameters pipeline, which has better determination of metallicity for high metallicity stars. ", "introduction": "The Sloan Digital Sky Survey (SDSS; York \\etal\\ 2000) saw first light in May 1998, and has been in routine survey operation mode since April 2000. It uses a 2.5m telescope with an unvignetted $3^\\circ$ field of view (Gunn \\etal\\ 2006) at Apache Point Observatory (APO) in Southern New Mexico, which is dedicated to wide-angle surveys of the sky. The first and second phases of the survey (SDSS-I and SDSS-II) were carried out with two instruments: a drift-scan imaging camera (Gunn \\etal\\ 1998) with 30 CCDs imaging in five filters ($ugriz$, Fukugita \\etal\\ 1996), and a pair of double spectrographs, fed by 640 optical fibers. The imaging data, essentially all of which have been taken under photometric and good-seeing conditions (Ivezi\\'c \\etal\\ 2004; Padmanabhan \\etal\\ 2008; see also Hogg \\etal\\ 2001), now cover more than 14,500 deg$^2$ in five filters (of which about 11,600 deg$^2$ was observed as part of SDSS-I/II), or roughly one third of the Celestial Sphere. The 50\\% completeness limit for point sources is $r=22.5$. The data have been analyzed with a sophisticated pipeline (Lupton \\etal\\ 2001) and have been photometrically (Tucker \\etal\\ 2006, Padmanabhan \\etal\\ 2008; see also Smith \\etal\\ 2002) and astrometrically (Pier \\etal\\ 2003) calibrated; the resulting catalog contains almost half a billion distinct detected objects. Well-defined samples of galaxies (Strauss \\etal\\ 2002; Eisenstein \\etal\\ 2001), quasars (Richards \\etal\\ 2002a), stars (Yanny \\etal\\ 2009) and other objects are selected for spectroscopy; the survey has obtained roughly 1.8 million spectra of galaxies, stars, and quasars as of Summer 2009. The principal scientific goal of SDSS-I (2000--2005) and much of SDSS-II (2005--2008) was to create a well-calibrated and contiguous imaging and spectroscopic survey of the Northern Galactic Cap at high Galactic latitudes, with the spectroscopy primarily focused on extragalactic targets. We refer to this project in what follows as the Legacy Survey. SDSS-II carried out two additional surveys. The Sloan Extension for Galactic Understanding and Exploration (SEGUE; Yanny \\etal\\ 2009) imaged a series of stripes sampling low Galactic latitudes (each 2.5$^\\circ$ wide and tens to hundreds of degrees long), together with spectroscopy of roughly 250,000 stars, to study Galactic structure, dynamics, and chemical composition. The SDSS Supernova Survey (Frieman \\etal\\ 2008) used approximately 80 repeat scans of a $2.5^\\circ \\times 100^\\circ$ stripe centered on the Celestial Equator in the Southern Galactic Cap to identify Type Ia supernovae with redshifts less than about 0.4, and to use them as cosmological probes (Kessler \\etal\\ 2009); almost 500 objects were spectroscopically confirmed as Type Ia supernovae. These data have been made public in a series of yearly data releases (Stoughton \\etal\\ 2002; Abazajian \\etal\\ 2003, 2004, 2005, 2006; Adelman-McCarthy \\etal\\ 2007, 2008; Abazajian \\etal\\ 2009; hereafter the EDR, DR1, DR2, DR3, DR4, DR5, DR6, and DR7 papers, respectively). These data have been used in over 3500 refereed papers to date for studies ranging from asteroids in the Solar System to the discovery of the most distant quasars. It was clear, as SDSS-II was nearing completion, that the wide-field spectroscopic capability of the SDSS telescope and system remained state-of-the-art, and a new collaboration was established to carry out further surveys with this telescope. This new phase, called SDSS-III, consists of four interlocking surveys; it is described in detail in a companion paper (Eisenstein \\etal\\ 2011). In brief, these surveys are: \\begin{itemize} \\item SEGUE-2. This survey is an extension of the spectroscopic component of the SEGUE survey of SDSS-II, extending the survey footprint in area and using revised target selection to increase the number of spectra in the distant halo of the Milky Way. SEGUE-2 used the SDSS-I/II spectrograph and ran from August 2008 through July 2009. \\item The Baryon Oscillation Spectroscopic Survey (BOSS). This survey will measure the baryon oscillation signature in the correlation function of galaxies and the quasar Lyman $\\alpha$ forest. BOSS started operations in Fall 2009, and consists of a redshift survey over 10,000 deg$^2$ of 1.5 million luminous red galaxies to $z \\sim 0.7$, together with spectroscopy of 150,000 quasars with $z > 2.2$. This has required increasing the imaging footprint of the survey, and we have obtained an additional $\\sim 2500$ deg$^2$ of imaging data in the Southern Galactic Cap using the SDSS imaging camera. In addition, in Summer 2009 the SDSS spectrographs underwent a major upgrade (new gratings, new CCDs, and new fibers) to improve their throughput and to increase the number of fibers from 640 to 1000. \\item The Multi-object APO Radial Velocity Exoplanet Large-area Survey (MARVELS) uses a fiber-fed interferometric spectrograph that can observe sixty objects simultaneously to obtain radial velocities accurate to 10--40 m s$^{-1}$ for stars with $9 < V< 12$. Each star will be observed roughly 24 times in a search for extrasolar planets. The instrument has been in operation since Fall 2008. \\item The Apache Point Observatory Galactic Evolution Experiment (APOGEE) will use a fiber-fed H-band spectrograph with a resolution of 30,000, capable of observing 300 objects at a time. The spectrograph will see first light in 2011, and will obtain high signal-to-noise ratio (S/N) spectra of roughly 100,000 stars in a variety of Galactic environments, selected from the Two-Micron All-Sky Survey (2MASS; Skrutskie \\etal\\ 2006). \\end{itemize} SDSS-III started operations in August 2008 and will continue through July 2014. As with SDSS-I/II, the data will periodically be released publicly; this paper describes the first of these releases. For continuity with the previous data releases of SDSS-I/II, we refer to it as the eighth data release, DR8. DR8 includes two significant items of new data relative to DR7: \\begin{itemize} \\item Roughly 2500 deg$^2$ of imaging data in the Southern Galactic Cap, taken as part of BOSS. \\item SEGUE-2 spectroscopy, consisting of 204 unique plates with spectra of roughly 118,000 stars. \\end{itemize} As with previous data releases, DR8 is cumulative, and includes essentially all data from the previous releases. However, this is not just a repeat of previous data releases, but also an enhancement. In particular, we have re-processed all SDSS-I/II imaging data using a new version of the imaging pipeline with a more sophisticated sky subtraction algorithm, and all stellar spectra have been re-processed with an improved stellar parameters pipeline. This paper provides an overview of DR8. Section~\\ref{sec:scope} describes the scope of the imaging and spectroscopic data. More details on the changes to the photometric pipeline and photometric calibration may be found in \\S\\ref{sec:photo}, while the spectroscopy, including SEGUE-2 target selection, is described in \\S\\ref{sec:spectro}. Methods for accessing these data are presented in \\S\\ref{sec:data}. We conclude, and outline the plan for future SDSS-III data releases, in \\S\\ref{sec:summary}. The data, and portals to access them, are described in greater detail at the DR8 website\\footnote{\\tt http://www.sdss3.org/dr8/}. ", "conclusions": "\\label{sec:summary} This paper describes the eighth data release of the Sloan Digital Sky Survey, consisting of all the SDSS data taken through Summer 2009, together with the final imaging of the Southern Galactic Cap completed in 2010 January. The images cover a footprint of over 14,500 deg$^2$; including repeat observations, the total quantity of imaging data is more than twice this value. All these data have been reprocessed with an updated version of the photometric pipeline, which gives modest improvements to the photometry of bright galaxies and fainter galaxies near them. In addition, DR8 contains the spectra of over 1.6 million galaxies, quasars, and stars, including 118,000 new stellar spectra from the SEGUE-2 survey, as well as 108 plates of data not previously released. With the completion of the imaging survey, the SDSS camera has been retired. SDSS-III is described in detail in Eisenstein \\etal\\ (2011); it will continue through 2014. This release contains data from two of its four surveys: SEGUE-2, and the imaging component of BOSS. BOSS spectroscopy has started, and its first year of data will be made available as part of the ninth data release. Plots showing the quality of those data may be found in Eisenstein \\etal\\ (2011) and White \\etal\\ (2011). In addition, the MARVELS survey is well underway, and the first scientific results have been published (Lee \\etal\\ 2011). Finally, APOGEE will probably have seen first light by the time this article is published, and data from that survey will first be released publicly in the tenth data release. We thank the referee, Andrew West, for comments that improved the paper. Funding for SDSS-III has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, and the U.S. Department of Energy. The SDSS-III web site is http://www.sdss3.org/. SDSS-III is managed by the Astrophysical Research Consortium for the Participating Institutions of the SDSS-III Collaboration including the University of Arizona, the Brazilian Participation Group, Brookhaven National Laboratory, University of Cambridge, University of Florida, the French Participation Group, the German Participation Group, the Instituto de Astrofisica de Canarias, the Michigan State/Notre Dame/JINA Participation Group, Johns Hopkins University, Lawrence Berkeley National Laboratory, Max Planck Institute for Astrophysics, New Mexico State University, New York University, Ohio State University, Pennsylvania State University, University of Portsmouth, Princeton University, the Spanish Participation Group, University of Tokyo, University of Utah, Vanderbilt University, University of Virginia, University of Washington, and Yale University." }, "1101/1101.3847_arXiv.txt": { "abstract": "In this paper, we study large-scale structures from numerical simulations, paying particular attention to supercluster-like structures. A grid-density-contour based algorithm is adopted to locate connected groups. With the increase of the linking density threshold from the cosmic average density, the foam-like cosmic web is subsequently broken into individual supercluster-like groups and further halos. To be in accordance with normal FOF halos with the linking length of $0.2$ in unit of the average separation of particles, halos in this paper are defined as groups with the linking density threshold $\\rho/\\bar \\rho=1+\\delta=80$, where $\\rho$ is the grid density, $\\bar \\rho$ is the average mass density of the universe. Groups with lower linking densities are then generally referred to as supercluster-like groups. By analyzing sets of cosmological simulations with varying cosmological parameters, we find that a universal mass function exists not only for halos but also for low-density supercluster-like groups until the linking density threshold decreases to $1+\\delta\\sim 8$ where the global percolation of large-scale structures occurs. We further show that the mass functions of different groups can be well described by the Jenkins form with the parameters being dependent on the linking density threshold. On the other hand, these low-density supercluster-like groups cannot be directly associated with the predictions from the excursion set theory with effective barriers obtained from dynamical collapse models, and the peak exclusion effect must be taken into account. Including such an effect, the mass function of groups with the linking density threshold $1+\\delta=16$ is in good agreements with that from the excursion set theory with a nearly flat effective barrier. A simplified analysis of the ellipsoidal collapse model indicates that the barrier for collapses along two axes to form filaments is approximately flat in scales. Thus in our analyses, we define groups identified with $1+\\delta=16$ as filaments. We then further study the halo-filament conditional mass function and the filament-halo conditional mass function, and compare them with the predictions from the two-barrier excursion set theory. The shape statistics for filaments are also presented. ", "introduction": "One of the key issues in cosmological studies is to understand the physical processes related to the structure formation in the universe. In the cold dark matter scenario, gravitational effects play essential roles in amplifying small density fluctuations generated in the early universe to shape the large-scale structures seen today. Being directly associated with galaxies and clusters of galaxies, virialized dark matter halos have been widely studied theoretically and observationally. Their mass function, which describes statistically the formation and evolution of dark matter halos, is shown by numerical simulations to follow a functional form universally valid for a wide range of cosmological models \\citep[e.g.,][]{she01, jen01}. Such a universality can be largely explained in the context of halo model which links initial density fluctuations to nonlinear dark matter halos through gravitational collapse models \\citep[e.g.,][]{pre74, coo02}. Considerable efforts have been made to improve the spherical collapse model to include more realistic characteristics in the modeling. It has long been realized that the anisotropic features contained in the initial density fluctuations can be magnified by nonlinear gravity \\citep{zel70,zel82}. It is expected that the collapse of a region first happens along the direction with the largest eigenvalue of the linear deformation tensor, thus leading to a sheet-like structure. Subsequent collapse along the direction of the second largest eigenvalue contracts the sheet structure to a filament. A halo can eventually form once further collapse occurs in the remaining direction. An ellipsoidal collapse model is developed to extend these considerations to the nonlinear regime \\citep[e.g.,][]{ick73, whi79}. The peak-patch scenario further includes the external tidal force self-consistently into consideration and improves the modeling of gravitational collapse around initial density peaks \\citep{bon96, bon96b}. \\citet{she01} and \\citet{she02} incorporate the peak-patch scenario into the excursion set approach in an averaged way. They first obtain statistically the averaged shape parameters of the initial tidal field. These averaged parameters are then used in the peak-patch ellipsoidal collapse model to derive the collapse criterion. It is noticed that on average, the halo formation is delayed due to the anisotropy of the gravitational effects. The predicted halo mass function (MF hereafter) is then in good agreements with that from numerical simulations. \\begin{figure*}% \\epsscale{1} \\plotone{f1} \\caption{Isodensity contours with different linking density thresholds in simulation JS12. Only the largest groups (top 100 ranking in mass) are shown. The global cosmic web can already be seen in average density contours ($1+\\delta=1$) in the top left panel. It gets sharper as $1+\\delta$ increases to 2 and 4 (top middle and top right panels, respectively). At $1+\\delta=8$ (bottom left), the cosmic web starts to break out, and large tree structures are seen. At $1+\\delta=16$, the web structure breaks into individual supercluster-like groups (bottom middle). Finally at $1+\\delta=80$, virialized halos are idenitified (bottom right).} \\end{figure*} Being very important in the hierarchy of large-scale structures, virialized dark matter halos of galaxy scale and above contain only $\\sim 40\\%$ of the total mass in the universe. Majority of the mass is distributed outside these large halos. In the language of halo model, the dominant fraction of the mass in the universe is contained in numerous small halos down to very low mass depending on the physical properties of dark matter particles. These small halos present anisotropic clustering patterns in space, and form, together with the massive halos, cosmic web structures. From the view point of the large halos, their formation and evolution are affected mainly by the clustering properties of the surrounding small halos as a whole. Thus to the zeroth order, the mass distribution around a large halo can be described by a smooth component without considering the individuality of small halos. This approach is clearly stated in the peak-patch scenario \\citep{bon96, bon96b}. In the framework of the excursion set theory, \\citet{she06} introduce filaments and sheets to model the large-scale mass distribution within which virialized halos are embedded. In their analyses, filaments are treated as an intermediate state of the ellipsoidal collapse when the collapse finishes along two directions. Then these filaments represent the smoothed version of the anisotropic mass distribution around fully collapsed halos. Various approaches have been proposed to geometrically define filamentary structures in cosmological simulations and observations. For example tessellation method is introduced to reconstruct the density field, and the edge between tessellations naturally constitutes a segment of filaments \\citep[e.g.,][]{Ick87,sch00,pla07,dtfe07}. The second order derivatives, namely, Hessian matrix, of the tidal field or the density field, is also widely used to classify halos, filaments, sheets and voids according to the signs of the eigenvalues of the matrix \\citep[e.g.,][]{hah07a,hah07b,sou08,pogosyan09,bon10a,bon10b,ara10}. \\citet{sto04} propose the so called Candy model for filament finding, in which a marked point process with a set of chosen parameters is used to reject points at disfavored directions and to locate elongated filamentary segments. These geometrically defined filaments, however, cannot be directly associated with the excursion-set-based filaments in \\citet{she06}. The dynamics of long geometrically defined filaments may not be dominated by the local field. Therefore they may break at the saddle point and accrete into the two ends separately during the late evolution. Furthermore, many geometrical definitions of filaments concentrate on the features of their spatial distribution rather than give rise to countable filamentary objects. To emphasize their dynamical structures and to compare with the results of the excursion set theory, in this paper, we mainly consider supercluster-like filamentary structures. We adopt a simple but natural definition of filaments by connectivity. Specifically, we first obtain the density field on a set of grids from particle positions in a simulation. Then the site percolation algorithm is applied to link cells together into groups by specifying a linking density threshold. At a high enough density threshold, only virialized halos are expected to be identified. At lower thresholds, filamentary superclusters surrounding virialized halos are located. The global percolation of the cosmic web occurs when the linking density threshold reaches a lower critical value. This is illustrated in Figure 1. From top left to bottom right, the linking density thresholds are $\\rho/\\bar\\rho=1+\\delta=1, 2, 4, 8, 16, 80$, respectively. As we will discuss later, the groups identified at $1+\\delta=80$ correspond to virialized halos. At $1+\\delta=16$, the individual structures seen in the plot are related to filaments defined in \\citet{she06}. At $1+\\delta=8$, we see the global percolation, and a large structure with a scale comparable to the size of the simulation box appears. At lower linking thresholds, the global cosmic web gets smoother. It should be noted however, that even for $1+\\delta=1$, the global web structure can still be seen clearly. The paper is organized as follows. \\S 2 presents our method in detail. In \\S 3, we analyze the mass function and the occupation statistics of the identified groups with different linking density thresholds. In \\S 4, we compare our results from simulations with predictions of the excursion set theory. Shape statistics are given in \\S 5. \\S 6 contains summaries and discussions. ", "conclusions": "Applying a grid-based site percolation method to numerical simulations, we study groups identified with different linking density thresholds $1+\\delta$. Groups with $1+\\delta=80$ correspond well to FOF dark matter halos. Lowering $1+\\delta$ allows us to find supercluster-like groups beyond virialized dark matter halos. As the linking density threshold approaches the average density of the universe, the global cosmic web structure can be naturally found. In the studies presented in this paper, we focus on supercluster-like groups, which are expected to be dynamically bound, although not virialized yet. These groups provide immediate environments to dark matter halos therein. Therefore understanding their properties is an important step towards understanding the environmental effects on the formation and evolution of galaxies. Our analyses reveal that similar to dark matter halos, the mass functions of supercluster-like groups for different simulations listed in Table 2 also follow a universal behavior. This universality is consistent with the consideration that these groups are gravitationally bound systems, and form mainly through their own gravitational interactions. In other words, the universality found for supercluster-like groups and that for dark matter halos should arise from the same origin. We further find that the Jenkins functional form can describe well the mass functions for not only halos, but also supercluster-like groups. An extended Jenkins mass function applicable to both halos and supercluster groups is then explicitly presented, in which the parameters $a$, $b$, and $c$ depend on the linking density threshold $1+\\delta$. As expected, the universality of the mass functions breaks down for groups with the linking density $1+\\delta\\le 8$ where the global web structures occur. We also compare the mass functions from simulations with those from the excursion set theory with effective barriers derived from the ellipsoidal collapse model. For halos with $1+\\delta=80$, consistent with other studies, the two agree very well with the parameter $a$ adjusted to be $a=0.707$ for the moving barrier. For supercluster-like objects, the ellipsoidal collapse model gives rise to a nearly flat barrier for filaments defined as two-axis collapse objects \\citep[e.g.][]{she06}. However, incorporating this barrier into the excursion set theory predicts a mass function that cannot fit to any mass function of supercluster-like groups identified in simulations with the linking density threshold $1+\\delta<80$. The off-center problem in the excursion set theory leads to a significant over prediction for the mass function at low mass end. Taking into account this problem in the comparison, we find that the mass function of the groups identified with $1+\\delta=16$ is in good agreement with that from the excursion set theory for two-axis collapse filaments. Defining these groups as filaments, we further study the halo-filament and filament-halo conditional mass functions. Deviations from the predictions of the two-barrier excursion set theory are seen, which are especially significant for filament-halo conditional mass function. The studies carried out in this paper can have important cosmological applications. The universality of the mass functions found for supercluster-like groups raises a possibility for us to probe cosmologies with supercluster abundances. It can also be applied to model statistically how the projection effects affect clusters' weak-lensing signals. In the very recent paper by \\citet{mur10}, they identify filamentary galaxy groups from the 2dfGRS survey using galaxy FOF method, and compare the properties of the groups with those of mock surveys constructed from numerical simulations. This study shows that it is becoming feasible observationally to analyze filamentary galaxy groups statistically, and they are in turn can be used as cosmological probes. Physically, we expect that these filamentary galaxy groups should be closely associated with the supercluster-like dark matter groups in our studies. To quantitatively understand the relation between the two, detailed FOD modeling for galaxies in supercluster-like dark matter groups is necessary. We discuss the FOD for subhalos in \\S 3.2. For galaxies, however, the FOD can be much more complicated, and thorough investigations are highly desired. It is further noted that analyses on real galaxy groups can only be done in redshift space. Redshift distortions from peculiar velocities of galaxies can affect group identifications, and further their mass functions and shape statistics considerably. For supercluster-like groups, their ambient member galaxies tend to be in the stage of coherent infall, and thus their distribution suffers oblate distortions in redshift space. On the other hand, for their virialized inner regions, the distortion can generate finger-of-God structures in redshift space. The detailed impacts of redshift distortions on entire supercluster-like groups will be explored in our future studies." }, "1101/1101.4183_arXiv.txt": { "abstract": "{From the point of view of X-ray astronomers, galaxy clusters are usually divided into two classes: ``cool core'' (CC) and ``non-cool core'' (NCC) objects. The origin of this dichotomy has been subject of debate in recent years, between ``evolutionary'' models (where clusters can evolve from CC to NCC, mainly through mergers) and ``primordial'' models (where the state of the cluster is fixed ``ab initio'' by early mergers or pre-heating). We found that in a representative sample (clusters in the GMRT Radio halo survey with available X-ray data), none of the objects hosting a giant radio halo can be classified as a cool core. This result suggests that the main mechanisms which can produce the ingredients to start a large scale synchrotron emission (most likely mergers) are the same that can destroy CC and therefore strongly supports ``evolutionary'' models of the CC-NCC dichotomy. } ", "introduction": "Galaxy clusters are often divided by X-ray astronomers into two classes: ``cool core''(CC) and ``non-cool core'' (NCC) objects on the basis of the observational properties of their central regions. One of the open questions in the study of galaxy clusters concerns the origin of this distribution. The original model which prevailed for a long time assumed that the CC state was a sort of ``natural state'' for the clusters, and the observational features were explained with the old ``cooling flow'' model: radiation losses cause the gas in the centers of these clusters to cool and to flow inward. Clusters were supposed to live in this state until disturbed by a ``merger''. Indeed, mergers are very energetic events that can shock-heat \\citep{burns97} and mix the ICM \\citep{gomez02}: through these processes they were supposed to efficiently destroy cooling flows. After the mergers, clusters were supposed to relax and go back to the cooling flow state in a sort of cyclical evolution (e.\\,g.\\, \\citealt{buote02}). With the fall of the ``cooling flow'' brought about by the \\xmmn and \\chandra observations (e.\\,g.\\, \\citealt{peterson01}), doubts were cast also on the interpretation of mergers as the dominant mechanism which could transform CC clusters into NCC. These doubts were also motivated by the difficulties of numerical simulations in destroying simulated cool cores with mergers (e.\\,g.\\, \\citealt{burns08} and references therein). More generally speaking, the question arose whether the observed distribution of clusters was due to a primordial division into the two classes or rather to evolutionary differences during the history of the clusters. \\\\ For instance \\citet{mccarthy04, mccarthy08} suggested that early episodes of non-gravitational pre-heating in the redshift range $1 0$, circles when $\\xi(r) < 0$) with the linear theory correlation function overplotted (solid colored lines). The vertical dotted line shows the initial mean interparticle spacing. {\\it Center panel}: The correlation function near the BAO scale. Vertical dotted lines show the expected shift from \\cite{Seo_etal2009} colored according to epoch. Also shown is the smooth $\\xi_{\\rm{nw}}(r)$ (black dot-dashed line) derived by fourier transforming $P_{\\rm{nw}}(k)$ from \\cite{Eisenstein_Hu1998}. {\\it Right panel}: The result of subtracting $\\xi_{\\rm{nw}}(r)$ from the $\\xi(r)$ measurements.} \\label{fig:xi_cdm} \\end{figure*} Having described and explained the non-linear evolution of the BAO-feature with our powerlaw setup in some detail, it is worth discussing the relevance of these results to the canonical $\\Lambda$CDM cosmology. We approach this task first by simply assessing the resemblance of our results to $\\Lambda$CDM. To aid in this comparison we performed a set of four simulations with an initial $\\Lambda$CDM spectrum ($\\Omega_m = 0.226, \\, \\Omega_\\Lambda = 0.774$) as in Fig.~\\ref{fig:pk_lcdm} but evolved with $\\Omega_m = 1, \\, \\Omega_\\Lambda = 0$ so that $\\sigma_8$ and $r_0 / r_{\\rm{bao}}$ in this case can avoid the freeze out limit and reach values comparable to the powerlaw setup. The $\\Lambda$CDM-like simulations presented here were performed with essentially identical parameters as the earlier fiducial simulations in terms of box size, force resolution and number of particles. We show the primary $\\xi(r)$ results in Fig.~\\ref{fig:xi_cdm}; the $r_0 / r_{\\rm{bao}}$ values for each output is shown in Table~\\ref{tab:cdm_sig8}. \\begin{table}[h] \\caption{$\\Lambda$CDM outputs} \\begin{center}\\label{tab:cdm_sig8} \\begin{tabular}{lcc} \\tableline\\tableline\\\\ \\multicolumn{1}{l}{} & \\multicolumn{1}{l}{$r_0 / r_{\\rm{bao}}$} & \\multicolumn{1}{c}{$\\sigma_8$} \\\\[2mm] \\tableline\\\\ & 0.003 & 0.25 \\\\ & 0.019 & 0.5 \\\\ & 0.040 & 0.75 \\\\ & 0.062 & 1.0 \\\\ & 0.106 & 1.5 \\\\ & 0.218 & 3.0 \\\\ \\tableline \\end{tabular} \\end{center} \\end{table} Fig.~\\ref{fig:xi_cdm} is fairly unremarkable except that it shows the non-linear evolution of the correlation function in $\\Lambda$CDM well past $z = 0$ and beyond the freeze out limit ($\\sigma_8 \\sim 1.3$). As in Fig.~\\ref{fig:xi_lcdm}, the overall amplitude of the BAO feature at fixed $r_0 / r_{\\rm{bao}}$ is more similar to the $n = -0.5$ case than to the cases with more large scale power. The models for the non-linear shift from \\cite{Seo_etal2009}, shown with vertical dotted lines in the center and right panels of Fig.~\\ref{fig:xi_lcdm}, predict shifts of $3-4$ \\% when extrapolated to the final output\\footnote{The prediction depends on whether one assumes their $\\alpha_{\\rm{shift}} - 1 \\propto D(z)^2$ formula, as expected from SPT, or instead uses their empirical fit where $\\alpha_{\\rm{shift}} -1 \\propto D(z)^{1.74}$. Fig.~\\ref{fig:xi_cdm} shows the predictions of the $D(z)^2$ model. The empirical model is similar.}. The center panel also shows the smooth $\\xi_{\\rm{nw}}(r)$ correlation function, computed from a fourier transform of $P_{\\rm{nw}}(k)$ from \\cite{Eisenstein_Hu1998}, and in the right panel $\\xi_{\\rm{nw}}(r)$ is subtracted from the simulation data. In the center panel the combination of strong damping of the BAO feature and noise in the $\\xi(r)$ measurement make any shift non-discernible. In the right panel the result of subtracting out $\\xi_{\\rm{nw}}(r)$ does visually resemble an attenuating gaussian (much more than $\\xi(r) / \\xi_{\\rm{nw}}(r)$, which is not shown), but it is unclear whether the apparent drift of the BAO peak towards smaller scales, especially by the last output, is truly from the non-linear shift or whether the effect is simply from the changing broadband shape of $\\xi(r)$. A plot of $(\\xi(r) - \\xi_{\\rm{pow}}(r))/D^2(z)$ versus $r$ from any of our fiducial simulations would show a similar trend. \\subsection{Perturbation Theory and Modeling} In\\S~\\ref{sec:interp} we showed that a phenomenological approach matched the results from our fiducial simulations rather well. Eq.~\\ref{eq:phenom} bears a close resemblance to the damped-exponential models often used in the literature \\citep[e.g.][]{Eisenstein_etal07,Seo_etal08}, and we emphasize our conclusion that the broadening (damping) of the bump (wiggles) depends on the {\\it pairwise} dispersion, $\\Sigma_{\\rm{pair}}^2$, rather than the rms displacement, $\\Sigma^2$, which is sensitive to bulk motions. In Fig.~\\ref{fig:damping}, we compare $\\Sigma_{\\rm{pair}}^2 / \\Sigma^2$ on a wide range of scales for a $\\Lambda$CDM spectrum (Fig.~\\ref{fig:pk_lcdm}). Although we expect the two formulae to converge to the same result as $r \\rightarrow \\infty$, it is nevertheless surprising that $\\Sigma_{\\rm{pair}}^2(r_{\\rm{bao}})$ differs by less than 2\\% from the $\\Sigma^2$ displacement. In the literature some groups use Eq.~\\ref{eq:Sigma} to predict the damping, while for others $\\Sigma^2$ is a free parameter that is fit to simulations \\citep[e.g.][]{Seo_etal08}. In our view, like that of \\cite{Eisenstein_etal07}, it is $\\Sigma_{\\rm{pair}}^2(r_{\\rm{bao}})$ that matters physically, and the success of models based on Eq.~\\ref{eq:Sigma} is a lucky coincidence that holds in $\\Lambda$CDM-like models but can fail, by an infinite factor, for powerlaw models. \\begin{figure} \\centerline{\\epsfig{file=eisdamp_vs_intPdq4.eps, angle=0, width=3.0in}} \\vspace{-0.2cm} \\caption{ The rms pairwise displacement (Eq.~\\ref{eq:Eis_etal07}; solid) at different scales using $\\Lambda$CDM initial conditions divided by the commonly used rms displacement formula (Eq.~\\ref{eq:Sigma}; dashed), which includes the contribution from bulk motions. Both quantities scale as the linear growth function squared, so the result shown is independent of epoch.} \\label{fig:damping} \\end{figure} Another widely-used phenomenological approach assumes a model for $P_{\\rm{NL}}(k)$ motivated by Renormalized Perturbation Theory \\citep[RPT;][]{Crocce_Scoccimarro2006}. In these models the non-linear shift comes directly from including $P_{22}(k)$ in the phenomenological form, or, in real space, from modeling the shift with the closely-related $\\xi_{\\rm{mc}}(r)$ ansatz and calibrating the amplitude of this term to N-body results (e.g. \\cite{Sanchez_etal2008,Montesano_etal2010,Crocce_etal2010}). Using our setup and a natural value for the amplitude of this term, in \\S~\\ref{sec:shift} we showed that this approach adequately captures the shift in real space for the first output of the $n = -1.5$ case (when $\\sigma_8 = 0.5$). By the second output (corresponding to $\\sigma_8 = 1$), however, it fails, and although not rigorously justified by the derivation of the term, we argue that the formula would more accurately predict the shift if the broadening of the bump could be incorporated into $\\xi_{\\rm{mc}}(r)$. This may have been previously unnoticed because the shift in $\\Lambda$CDM when $\\sigma_8 \\sim 1$ is smaller than the shift in the $n = -1.5$ case, and the amplitude of the bump, i.e., $\\xi(r_{\\rm{bao}})$, is significantly smaller in $\\Lambda$CDM than in the $n = -1.5$ setup. % Finally, the success of the ``coupling strength'' RGPT method \\citep{McDonald2007} in matching our simulation results, both in fourier space and in real space, may certainly be informative to ongoing efforts to model the BAO evolution with {\\it ab initio} predictions from PT. \\cite{Carlson_etal09} show that this scheme also does a reasonable job in predicting the non-linear power spectra of $\\Lambda$CDM and cCDM cosmologies. Except for SPT \\cite{Makino_etal1992} we ignored other PT schemes, but in principle the predictions from many other PT schemes could be compared to our simulation results and useful insights gained from the kind of comparisons presented in \\S~\\ref{sec:pt}. This would no doubt be useful for BAO studies, and, more broadly, \\cite{Valageas_Nishimichi2011} find that the largest deficiency of the halo model is in capturing the transition from the 1-halo to 2-halo term, precisely the scales where the perturbation theory predictions are most important. Motivated by the importance of accurate modeling of the BAO feature in large scale structure for interpreting the results of future dark energy experiments, we have used N-body simulations to investigate the evolution of a BAO-like feature in a simpler, alternative setting, where it modulates an underlying powerlaw initial power spectrum in an $\\Omega_m = 1$ universe. Specifically, our initial conditions have the correlation function defined by Eq.~\\ref{eq:powgaus}, with a gaussian multiplicative bump centered at scale $r_{\\rm{bao}}$ and the amplitude $A_{\\rm{bump}}$ and width $\\sigma_{\\rm{bao}}$ chosen in approximate agreement with $\\Lambda$CDM expectations. The corresponding initial power spectrum follows Eq.~\\ref{eq:analyticapprox} to an excellent approximation. For given values of $A_{\\rm{bump}}$, $\\sigma_{\\rm{bao}}$, and the powerlaw spectral index $n$, non-linear matter clustering statistics (including the correlation function and power spectrum) should depend only on the ratio $r_0 / r_{\\rm{bao}}$, where $r_0$ is the correlation length defined by $\\xi(r_0) = 1$. We evolve our simulations to values of $r_0 / r_{\\rm{bao}}$ much higher than traditional $\\Lambda$CDM models, with final outputs corresponding to $\\sigma_8 = 4.0 \\, (n = -1.5)$, $6.0 \\, (n = -1)$, and $12.0 \\, (n = -0.5)$ if we define a physical scale by setting $r_{\\rm{bao}} = 100 h^{-1}$Mpc. Our standard simulations have box side $L_{\\rm{box}} / r_{\\rm{bao}} = 20$ and $512^3$ particles. We use our simulations to develop physical intuition for BAO evolution and to test analytic descriptions in a regime far from that where they have been tested previously. In this respect, the spirit of our exercise is similar to the cCDM investigation of \\cite{Carlson_etal09} and \\cite{Padmanabhan_white09}.\\footnote{ Another notable study is \\cite{Ma2007} who investigated the non-linear evolution of a $\\Lambda$CDM spectrum plus a fourier-space spike on scales relevant to BAO. \\cite{Blake_Glazebrook2003} and \\cite{Smith_etal2008} have also discussed toy models for BAO but without investigating non-linear effects.} Consistent with $\\Lambda$CDM studies, we find that the strongest effect of non-linear evolution on the BAO feature in $\\xi(r)$ is to flatten and broaden the bump, with $A_{\\rm{bump}}$ decreasing and $\\sigma_{\\rm{bao}}$ increasing. To a good approximation, failing only at late times in the $n = -0.5$ model, the area of the gaussian bump, proportional to $A_{\\rm{bump}} \\times \\sigma_{\\rm{bao}}$, remains constant, which suggests that pairs are ``diffusing'' out of the shell corresponding to the initial BAO feature (see the physical description of \\cite{Eisenstein_etal07}). The evolution of the bump width is well described by a model in which the non-linear $\\sigma_{\\rm{bao}}$ is the quadrature sum of the initial width and a length proportional to $\\Sigma_{\\rm{pair}}$, the rms relative displacement (computed from linear theory) of pairs separated by $r = r_{\\rm{bao}}$. The constant of proportionality varies with $n$, but the same constant that describes our standard $n = -1$ model also describes the faster evolution of an $n = -1$ model with a ``skinny'' initial bump, supporting the validity of the diffusion interpretation. For $n = -1.5$ (where the relevant integral converges without a small scale UV cutoff) the diffusion constant computed {\\it ab initio} describes the bump evolution accurately. We emphasize that it is $\\Sigma_{\\rm{pair}}$ rather than the rms absolute displacement $\\Sigma$ that is relevant to analytic descriptions of our models. The latter quantity has an infrared divergence for $n \\leq -1$, but this divergence corresponds to bulk translations induced by very large scale modes, which cannot affect the BAO peak itself. We think that the appearance of $\\Sigma$ rather than $\\Sigma_{\\rm{pair}}$ in many analytic models of BAO evolution is at best an approximation restricted to CDM-like models with a turnover in $P(k)$; by coincidence, $\\Sigma \\approx \\Sigma_{\\rm{pair}}(r_{\\rm{bao}})$ for $\\Lambda$CDM. The location of the BAO peak, defined by the scale $r_{\\rm{peak}}$ of a gaussian fit to the non-linear $\\xi(r)$ divided by the linear theory powerlaw, stays constant within the statistical precision of our measurements for the $n = -0.5$ and $n = -1$ models, even when these are evolved to a highly non-linear stage where the bump amplitude has dropped by a factor of $\\sim 4-10$ from its initial value. For $n = -1.5$, on the other hand, the peak location shifts to smaller $r$, an effect that is already noticeable at the first output ($r_0 / r_{\\rm{bao}} = 0.024$, equivalent to $\\sigma_8 = 0.5$) and that grows to a 30\\% drop by $r_0 / r_{\\rm{bao}} = 0.386$ (equivalent to $\\sigma_8 = 4.0$). The analytic models of \\cite{Smith_etal2008} and \\cite{CrocceScoccimarro08} accurately predict that shifts should be much larger for $n = -1.5$ than for $n = -0.5$ and $n = -1$, and the \\cite{Smith_etal2008} model accurately describes the evolution of the peak location for $n = -1.5$. However, both models predict non-linear shifts in the $n = -0.5$ and $n = -1$ cases that are inconsistent with our simulation results at late times. We carried out a number of additional numerical tests varying either numerical parameters or the physical model. Our fiducial simulations have $L_{\\rm{box}} / r_{\\rm{bao}} = 20$ and an initial mean interparticle spacing smaller than $r_{\\rm{bao}}$ by a factor of $r_{\\rm{bao}} / n_p^{-1/3} = 25$. We found consistent results in simulations with $L_{\\rm{box}} / r_{\\rm{bao}} = 10$ and $r_{\\rm{bao}} / n_p^{-1/3} = 50$, indicating that a box size ten times the BAO scale is acceptable. We found marginal discrepancies for $256^3$ simulations with $r_{\\rm{bao}} / n_p^{-1/3} = 12.5$. Success of the box size test and other internal consistency tests is achieved only because we include the integral constraint corrections described in Appendix~\\ref{ap:xicorr}, which make a noticeable difference for $n = -1$ and an important difference for $n = -1.5$. In other tests, we show that BAO evolution is nearly identical in an $\\Omega_m = 1$ model and a model with $\\Omega_m = 0.3$, $\\Omega_\\Lambda = 0.7$ (and the same initial conditions) provided they are evaluated at the same value of $r_0 / r_{\\rm{bao}}$ (or, equivalently, the same value of the linear growth function). For more thorough tests of analytic models, we turned to a fourier space description using the non-linear matter power spectrum. A ``phenomenological'' model in which we combine numerical results for the non-linear power spectrum of a pure powerlaw model (Appendix~\\ref{ap:purepow} and references therein) with our gaussian fits to the evolution of the BAO bump in $\\xi(r)$ gives a remarkably accurate description of the full non-linear outputs of the $n = -1$ and $n = -1.5$ models. This model assumes no shift of the $\\xi(r)$ peak location for $n = -0.5$ and $n = -1$ and $r_{\\rm{peak}} / r_{\\rm{bao}} = 1 - 1.08 (r_0 / r_{\\rm{bao}})^{1.5}$ for $n = -1.5$. The success of this model suggests that the BAO bump has little effect on the non-linear evolution of the underlying ``smooth'' power spectrum. At least for $r_0 / r_{\\rm{bao}} < 0.2$, we expect that this model is a {\\it more} accurate description than our numerical $P(k)$ measurements themselves, since it draws on self-similar scaling results from pure powerlaw spectra that have wider dynamic range than our simulations. We compared our results to predictions of two {\\it ab initio} analytic approaches, ``standard'' 1-loop perturbation theory (SPT; e.g. \\cite{Vishniac1983,Makino_etal1992}) and the ``coupling strength'' RGPT scheme of \\cite{McDonald2007}. This scheme provides a quite accurate description of the low-$k$ evolution in all cases, including $n = -1.5$ where the peak location shifts significantly, and it produces good but not perfect agreement with the evolution of the $\\xi(r)$ bump in configuration space. For SPT, we break up the terms into distinct ``interactions'' between the powerlaw and ``wiggle'' components of the linear power spectrum, both to obtain physical insight and to allow us to define a more accurate ``SPT+'' scheme that uses numerical results for pure powerlaw evolution and perturbation theory to describe the interaction terms that involve the ``wiggle'' spectrum. SPT alone gives a reasonable description of the early $P(k)$ outputs for $n = -1.5$, but on the whole ``coupling strength'' RGPT is substantially more accurate and has a wider range of validity. The high statistical precision achievable with future BAO surveys --- with $\\sim 0.2 \\%$ cosmic variance distance scale errors for $z \\geq 1$ and redshift bins $\\Delta z \\approx 0.2$ \\citep{Seo_Eisenstein2007} --- puts stringent demands on theoretical models. Exploiting the power of these surveys will require large numerical simulations supplemented by the physical insight and modeling flexibility afforded by analytic methods. The simulation results presented here offer valuable ``stress tests'' of numerical and analytic approaches in regimes beyond those where they are usually applied, and they allow isolation of distinct physical effects. Two natural directions that we plan to explore in future work are the clustering of biased tracers --- in particular the massive halos expected to host luminous galaxies --- and the impact of redshift-space distortions on BAO measurement from galaxy clustering. We will also investigate the impact of the initial conditions algorithms, comparing the scheme advocated by \\cite{Sirko2005} for simulation ensembles to the traditional scheme of mean density boxes used here. The combination of future BAO surveys and improved theoretical models will lead, ultimately, to new insights on the energy and matter contents of the cosmos." }, "1101/1101.2304.txt": { "abstract": "{}{}{}{}{} % 5 {} token are mandatory \\abstract % context heading (optional) % {} leave it empty if necessary {The extended environment of galaxies contains a wealth of information about the formation and life cycle of galaxies which are regulated by accretion and feedback processes. Observations of neutral hydrogen are routinely used to image the high brightness disks of galaxies and to study their kinematics. Deeper observations will give more insight into the distribution of diffuse gas in the extended halo of the galaxies and the inter-galactic medium, where numerical simulations predict a cosmic web of extended structures and gaseous filaments.} % aims heading (mandatory) {To observe the extended environment of galaxies, column density sensitivities have to be achieved that probe the regime of Lyman limit systems. {\\HI} observations are typically limited to a brightness sensitivity of $N_{HI} \\sim 10^{19}$ cm$^{-2}$, but this must be improved upon by $\\sim2$ orders of magnitude.} % methods heading (mandatory) {In this paper we present the interferometric data of the Westerbork Virgo {\\HI} Filament Survey (WVFS) -- the total power product of this survey has been published in an earlier paper. By observing at extreme hour angles, a filled aperture is simulated of $300\\times25$ meters in size, that has the typical collecting power and sensitivity of a single dish telescope, but the well defined bandpass characteristics of an interferometer. With the very good surface brightness sensitivity of the data, we hope to make new {\\HI} detections of diffuse systems with moderate angular resolution.} % results heading (mandatory) {The survey maps 135 degrees in Right Ascension between 8 and 17 hours and 11 degrees in Declination between $-$1 and 10 degrees, including the galaxy filament connecting the Local Group with the Virgo Cluster. Only positive declinations could be completely processed and analysed due to projection effects. A typical flux sensitivity of 6~mJy~beam$^{-1}$ over 16 km~s$^{-1}$ is achieved, that corresponds to a brightness sensitivity of $N_{HI} \\sim 10^{18}$~cm$^{-2}$. An unbiased search has been done with a high significance threshold as well a search with a lower significance limit but requiring an optical counterpart. In total, 199 objects have been detected, of which 17 are new {\\HI} detections.} % conclusions heading (optional), leave it empty if necessary {By observing at extreme hour angles with the WSRT, a filled aperture can be simulated in projection, with a very good brightness sensitivity, comparable to that of a single dish telescope. Despite some technical challenges, the data provide valuable constraints on faint, circum-galactic {\\HI} features. } ", "introduction": "In the current epoch, numerical simulations predict that most of the baryons are not in galaxies, but in extended gaseous filaments, forming a Cosmic Web (e.g. \\citealp{1999ApJ...511..521D, 1999ApJ...514....1C}) Galaxies are just the brightest pearls in this web, as the baryons are almost equally distributed amongst three components: (1) galactic concentrations, (2) a warm-hot intergalactic medium (WHIM) and (3) a diffuse intergalactic medium. Direct detection of the intergalactic gas is very difficult at UV, EUV or X-ray wavelengths \\citep{1999ApJ...514....1C} and so far the clearest detections have been made in absorption (e.g. \\citealp{2007ApJ...658..680L, 2008ApJS..177...39T}). In this and previous papers in this series, we make an effort to detect traces of the intergalactic medium in emission, using the 21-cm line of neutral hydrogen. Most of the gas in the Cosmic Web will be highly ionised, due to the moderately high temperatures above $10^4$ Kelvin, resulting in a low neutral fraction and relatively low neutral column densities. A more detailed background and introduction on this topic is outlined in \\cite{2010PhDT..APOPPING} and \\cite{2010arXiv1012.3236P} To investigate column densities that probe the Lyman Limit System regime, very deep {\\HI} observations are required with a brightness sensitivity significantly better than $N_{HI} \\sim 10^{19}$ cm$^{-2}$. Reaching these column densities is important to learn more about the distribution of neutral hydrogen in the inter-galactic medium and to have a better understanding of feedback processes that fuel star formation in galaxies. In \\cite{2010arXiv1012.3236P} and \\cite{2010A&A...HIPASS} two {\\HI} surveys have been presented that reach these low column densities in a region of $\\sim 1500$ square degrees. The first data product described in \\cite{2010arXiv1012.3236P} is the total power data of the Westerbork Virgo Filament Survey, an {\\HI} survey mapping the galaxy filament connecting the Virgo Cluster with the Local Group. The survey spans 11 degrees in Declination from $-$1 to +10 degrees and 135 degrees in Right Ascension between 8 and 17 hours. This survey has a point source sensitivity of 16 mJy beam$^{-1}$ over 16 km s$^{-1}$ corresponding to a column density of $N_{HI} \\sim 3.5 \\cdot 10^{16}$ cm$^{-2}$. The second data product presented in \\cite{2010A&A...HIPASS} is reprocessed data, using original data that has been observed for the {\\HI} Parkes All Sky Survey \\citep{2001MNRAS.322..486B, 2006MNRAS.371.1855W}. The 1500 square degree region overlapping the WVFS was reprocessed to permit comparison between these data products and detections. The point source sensitivity of the reprocessed HIPASS data is 10 mJy beam$^{-1}$ over 26 km s$^{-1}$, corresponding to a column density of $N_{HI} \\sim 3.5 \\cdot 10^{17}$ cm$^{-2}$. In this paper, a third data product is presented: the cross-correlation data of the Westerbork Virgo Filament Survey. As explained in \\cite{2010arXiv1012.3236P}, the aim of the WVFS was to achieve very high brightness sensitivity in a large region of the sky, to permit detection of {\\HI} features that probe the neutral component of the Cosmic Web. The configuration of the array was chosen such that the dishes of the interferometer form a filled aperture of $\\sim 300$ meters in projection by observing at extreme hour angles. Because of the much smaller beam size compared to the WVFS total-power or HIPASS observations, we will be able to identify brighter clumps within diffuse features if these are present. The special observing configuration creates some technical challenges itself. This novel observing strategy requires non-standard data-reduction procedures, which will be explained in section 2. In section 3 we will present the results, starting with a list of detected features. Objects are sought both blindly, by using a high signal-to-noise threshold, and in conjunction with a known optical counterpart by using a lower threshold. New {\\HI} detections and diffuse structures are briefly discussed, however detailed analysis of these features will be presented in a follow up paper, also discussing new and tentative detections obtained from the WVFS total-power data and the re-processed HIPASS data as described in \\cite{2010arXiv1012.3236P} and \\cite{2010A&A...HIPASS}. We will end with a short discussion and conclusion, summarizing the main results. ", "conclusions": "The WSRT has been used in a very novel observing mode to simulate a filled aperture in projection of $300 \\times 25$ meters by observing at very extreme hour angles. Because of the very short observing times per pointing it is a technical challenge to observe and reduce this data, while still achieving useful results. In total 22,000 pointings have been observed that cover a total area of $\\sim 1500$ square degrees. Each pointing has been observed two times for a period of one minute. Normally an integration time of one minute with an interferometer is not sufficient to fill the {\\it u,v} plane, however as there are essentially no gaps between the individual antennas in projection, and the two scans have a complimentary orientation, a well defined synthesized beam could be formed. The observing method is limited to a narrow range in declination, but has been very successful.\\\\ In the reduced data we reach a flux sensitivity of $\\sim 6$~mJy~beam$^{-1}$ over 16 km~s$^{-1}$. The synthesised beam has an average size of $395 \\times 245$ arcsec, which results in a brightness sensitivity of $N_{HI} \\sim 10^{18}$ cm$^{-2}$. Such a brightness sensitivity can typically only be achieved with single-dish observations. Because the WSRT is an interferometer, the calibration and stability of the bandpass is significantly superior to that of a single dish. The drawback of interferometric observations is that a {\\it cleaning} step has to be applied to correct for the synthesised beam shape, especially for bright objects. Although the synthesised beam is well defined, the side-lobes in our observations are very strong, making the {\\it cleaning} step a critical one. Because of the large size of the survey, only one pass of {\\it cleaning} deconvolution could be applied. Improved deconvolution results would require detailed masking of real emission features during component identification. Because of the relatively high side-lobe level, any automatic masking procedure is unlikely to be reliable. Each emission peak would have to be inspected visually, to be able to distinguish a real emission feature from a side-lobe artifact. Because of the size of the survey this approach was not considered practical. As a result of these limitations, there are low level artifacts in the data, although the flux sensitivity in the reduced data is very good.\\\\ An extra complication in processing the data is that the north-south synthesized beamsize increases towards low Declinations since the WSRT is an east-west array. Furthermore, the natural image plane projection of an east-west array is the North Celestial Pole ({\\it NCP}) projection \\citep{1971PhDT.......153B} which becomes undefined at zero degrees. As a result, only positive declinations could be analysed in our mosaiced images. Although this did not dramatically affect our results, it is still a major point of concern. We found another serious complication in using an {\\it NCP} projection in a wide field survey close to a declination of zero degrees. This complication is not caused by a shortcoming in the projection, but more likely a shortcoming in the imaging software that has been used. When gridding the observed data or {\\it u,v} coordinates to the projected plane, the observed flux is not being conserved due to an incorrect wighting when combining pointings in a mosaic. At positive declinations,the imaged flux of data below the declination of the reference pixel is diluted while for data above the reference pixel the imaged flux is enhanced compared to the observed flux. This effect is very apparent when using a {\\it NCP} projection but probably also happens when using other projections. When observing small fields, or fields between $\\sim 20$ and $\\sim 70$ degrees in declination this effect is negligible, however in our case a significant correction to the derived fluxes had to be applied. In general, other image plane projections are required at Declinations near zero degrees to enable both positive and negative Declinations to be imaged simultaneously. However several major interferometers are East-West arrays, including the WSRT and the Australia Telescope Compact Array (ATCA). All these telescopes are being upgraded, partly to serve as a pathfinders for the SKA (Square Kilometre Array). A general aim of future telescopes, especially the ones that use a FPA (Focal Plane Array) is to conduct large surveys of the entire sky. Ideally these surveys will have significant overlap with deep optical surveys, however several major optical surveys are concentrated at Declinations near zero degrees.\\\\ Two search methods have been applied to the reduced WVFS data, both using the {\\it Duchamp} source finding algorithm. The first method is a blind search at a peak brightness limit of $8\\sigma$. The second method uses an initial peak brightness limit of $5\\sigma$, but has the additional requirement that all detected features need to have an optical counterpart. In the blind search 138 objects have been detected, while the second search resulted in 198 {\\HI} counterparts to cataloged optical galaxies. Of all the detections 16 are new {\\HI} detections and only 1 detection does not have an optical counterpart. On average, the interferometric total fluxes of detections are $\\sim 10$\\% lower than the catalogued fluxes in the HIPASS archive. There are many features in the cube with a peak of 3 or 4$\\sigma$ and an integrated flux that probably exceeds 8$\\sigma$. It is very likely that many of these features are real, however they cannot be identified reliably by automated source finders. In a subsequent paper, the WVFS cross-correlation data will be compared with the WVFS total-power data and the re-reduced HIPASS data as described in \\cite{2010arXiv1012.3236P} and \\cite{2010A&A...HIPASS}. Both surveys contain several new {\\HI} detections and tentative filaments. As this is a limited number of sources, we can do a targeted search in the WVFS cross-correlation data. Although the brightness sensitivity of each of the three surveys is almost an order of magnitude different, comparison of the data will be useful to distinguish extended and diffuse emission from denser {\\HI} clumps." }, "1101/1101.0978_arXiv.txt": { "abstract": "We present a series of numerical experiments that model the evolution of magnetic flux tubes with a different amount of initial twist. As a result of calculations, tightly twisted tubes reveal a rapid two-step emergence to the atmosphere with a slight slowdown at the surface, while weakly twisted tubes show a slow two-step emergence waiting longer the secondary instability to be triggered. This picture of the two-step emergence is highly consistent with recent observations. These tubes show multiple magnetic domes above the surface, indicating that the secondary emergence is caused by interchange mode of magnetic buoyancy instability. As for the weakest twist case, the tube exhibits an elongated photospheric structure and never rises into the corona. The formation of the photospheric structure is due to inward magnetic tension force of the azimuthal field component of the rising flux tube (i.e., tube's twist). When the twist is weak, azimuthal field cannot hold the tube's coherency, and the tube extends laterally at the subadiabatic surface. In addition, we newly find that the total magnetic energy measured above the surface depends on the initial twist. Strong twist tubes follow the initial relation between the twist and the magnetic energy, while weak twist tubes deviates from this relation, because these tubes store their magnetic energy in the photospheric structures. ", "introduction": "} Flux emergence is one of the key mechanisms in various solar activities. It is widely accepted that the emerging flux has a form of a twisted flux tube so as not to be collapsed by the convective motions during its ascent in the solar interior. Emerging flux transports magnetic energy and helicity from the convection zone to the atmosphere, which yields active regions including sunspots. Magnetic helicity in the corona stores free energy that can be released in the forms of flares and coronal mass ejections (CMEs) (e.g. \\cite{hey77}). Many numerical experiments have revealed the dynamics of the flux emergence. \\citet{sch79} carried out two-dimensional magnetohydrodynamic (MHD) simulations to study the cross-sectional evolution of the emerging flux tube (see also \\cite{mor96,emo98}). \\citet{shi89} calculated the two-dimensional evolution of the undular mode of magnetic buoyancy instability (Parker instability: \\cite{par66}) to reproduce the formation of an $\\Omega$-shaped coronal loops. \\citet{tor10a} and \\citet{tor10b} gave numerical studies of emerging fluxes from much deeper convection zone ($\\sim -20,000$\\ km) to the corona. The three-dimensionality also exerts an influence on emerging process of magnetic flux evolution. \\citet{mat93} produced the first three-dimensional work of the Parker instability using a magnetic flux sheet and a flux tube. \\citet{fan01} compared her numerical results of the twisted tube's emergence with observations of an active region. In this paper, we perform three-dimensional simulations of the twisted emerging flux tube from the uppermost convection zone to the corona. Our aim is to study the effect of the initial twist on the emergence process. A series of parametric studies on the flux tube's twist was done by \\citet{mur06}. Our work is dedicated to further detailed analyses, especially focusing on the effect of the initial twist on the resulting tube's structure (photospheric lateral expansion and multiple magnetic domes) and on the consequent coronal magnetic energy. For numerical experiments, we used the same conditions as those by \\citet{mur06}; we calculated ten cases of different twist parameters that cover their three runs. As a result of experiments, we found that the evolution depends on the initial twist. When the twist is strong enough, the evolution to the corona reveals two-step way, showing a deceleration and a lateral expansion near the solar surface, although the case with weaker twist spends more time waiting for the secondary emergence to occur \\citep{mag01,arc04,mur06}. This picture of the two-step emergence is highly consistent with recent observations by \\citet{ots10}, especially its horizontally expanding speed and the rising speed. In addition to the confirmation of the results by \\citet{mur06}, it is also found that multiple magnetic domes are built and plasma accumulates in between the domes when the secondary emergence starts. At this moment, the direction of the field lines is almost perpendicular to the alignment of the domes, indicating that the second-step emergence is due to the interchange-mode instability. If the initial twist is too weak, the tube extends widely near the surface and further evolution never takes place, because the magnetic tension force of the azimuthal component cannot hold the tube's coherency. Also, we newly found that the total magnetic energy measured above the surface relies on the initial twist. In the strong twist regime, the resulting magnetic energy follows the initial relation between the twist and the magnetic energy. In the weak twist regime, however, the magnetic energy deviates from the initial rule, because the tube with weak twist stores magnetic energy around the photosphere. The rest of the paper is organized as follows. In Section \\ref{sec:setup}, we describe the numerical model. The simulation results are shown in Section \\ref{sec:results}. Summary and discussion are given in Section \\ref{sec:summary} and \\ref{sec:discussion}, respectively. ", "conclusions": "} In this paper, we carried out three-dimensional MHD simulations to investigate the effect of the initial twist on the flux tube evolution. Here, we summarize the results: \\begin{itemize} \\item Initially, the flux tube rises through the convection zone due to its magnetic buoyancy. Reaching the surface, the tube expands laterally to make a ``photospheric tongue.'' The secondary emergence occurs after sufficient flux accumulates within the photosphere. Due to the interchange mode instability, the tube builds multiple domes above the surface, between which the fluids piles up. Finally, the flux tube arrives at $z/H_{0}\\sim 60$ as a single dome. The overall emergence is consistent with the recent observations (e.g. \\cite{ots10}). \\item We run ten twist cases to investigate the effect of the initial twist. Nine out of ten reach the coronal height ($z/H_{0}\\ge20$) showing two-step emergence, while the weakest twist case fails to rise further above the surface ($qH_{0}=0.05$). In the two-step emergence regime, the rise time becomes shorter with increasing initial twist, which is consistent with the previous calculations by \\citet{mur06}. The photospheric tongue is more noticeable in weaker twist case. \\item We study the force components at the solar surface for different twist cases at the time $t/\\tau_{0}=40$. The stronger the initial twist is, the larger the inward magnetic tension is, resulting the tube keeps its coherency. In the weak twist case, the magnetic tension is much less effective, causing the tube distorted. At the same time, the strong twist tube rises further into the atmosphere mainly by the magnetic pressure gradient. \\item We found that the photospheric fields of the middle twist case ($qH_{0}=0.15$) undulate in the later phase of the emergence. The field lines gradually rise into the corona as longer loops. The photospheric and coronal fields are almost parallel to the axis of the initial flux tube. These features remind us of the resistive emergence model by \\citet{par04}. \\item We measure the magnetic energy $E_{\\rm mag}$ above the surface. The energy plot follows the initial $q^{2}$ law when the twist is strong ($qH_{0}\\ge 0.2$), while, for weaker twist cases ($qH_{0}\\le 0.2$), the energy depends on the time difference between reaching the surface and the corona. That is, weakly twisted tube takes more time for magnetic flux to accumulate near the surface and the secondary instability to be triggered. \\end{itemize} } In Section \\ref{sec:results}, we showed the time-evolution of the twisted flux tube. When the second-step emergence starts, multiple domes are observed above the surface and fluid is trapped between the expanding magnetic structures (see Figure \\ref{fig:ro-mag}). At this time, field lines are directed perpendicular to the alignment of the magnetic domes. For middle twist tube, we also found undulating fields near the surface, emerging into the corona (see Figure \\ref{fig:bvct}). In this section, we discuss these features in connection with future observations. \\subsection{Twist Intensity and the Interchange Instability \\label{sec:interchange}} In Section \\ref{sec:overview}, we saw that, as the twist decreases, the interchange mode instability becomes more noticeable. However, this is contrary to the expectation that the azimuthal field should be less pronounced in a weaker twist case. It may be because, in a weak twist case, the tube extends laterally near the photosphere and thus the twist increases. As the tube develops the interchange instability, field lines perpendicular to the alignment of the magnetic domes become more pronounced (see Sections \\ref{sec:param} and \\ref{sec:compare}). \\subsection{Twist of the Actual Flux Tube in the Sun \\label{sec:twist}} Multiple magnetic structures and the density accumulation between them are also found in previous observations and calculations. \\citet{par04} found that photospheric fields are undulating at its earlier phase of the flux emergence event, and proposed a resistive emergence model that undulating multiple loops reconnect with each other to make larger coronal fields. \\citet{iso07} carried out two-dimensional MHD simulation to study the evolution of the serpentine magnetic loops (resistive emergence model), finding that density accumulates in between the magnetic loops; their elongated vertical plasma structures are similar to our results (for three-dimensional study, see \\citet{arc09}). However, in our model, field lines are directed almost perpendicular to the alignment of the domes, which is against the observations of the undular field lines (e.g. \\cite{par04}). The direction of the field lines are the consequence of the initial tube's twist. Therefore, the difference between the present calculations and the observations indicates that the actual twist of the flux tube beneath the surface may be much weaker than those assumed in our models (e.g. $qH_{0}=0.15$ for Figure \\ref{fig:ro-mag}). On the other hand, flux tube with insufficient twist was found to fail to rise through the convection zone \\citep{mor96,emo98,tor10b}. It is because the weak azimuthal field of the flux tube cannot hold its coherency during its ascent within the solar interior. Therefore, one of the important problems to be solved is the emergence of the flux tube with much weaker twist ($qH_{0}\\le 0.1$). We also found that the medium twist tube ($qH_{0}=0.15$) reveals the undulating fields at the surface, which gradually rise into the upper atmosphere as longer coronal loops (Section \\ref{fig:bvct}). These fields are directed parallel to the main axis of the initial flux tube. This picture seems well accorded with the resistive emergence model. However, it is in the later phase that this undulatory evolution is observed, and, in the earlier phase, the field lines are perpendicular to the original tube's axis (see Figure \\ref{fig:ro-mag}). Therefore, the reproduction of the undulating fields parallel to the original axis \\citep{par04} is not achieved. \\subsection{Formation of Undulating Photospheric Fields \\label{sec:undulation}} Recently, \\citet{che10} have conducted a radiative MHD simulation of the formation of an active region. They showed that the rising flux tube flattens to make a pancake-like structure near the surface, and that the convective flows create serpentine field lines. In the present study, which does not include the convection, we also found the sideway expansion of the rising field at the photosphere (tongue) and the undulation of the photospheric fields in the later phase of the emergence of the weaker twist tube. These features are also confirmed by \\citet{arc10}; their simulations do not take account of the convection effects, either. Therefore, we can see that other mechanisms, apart from the convection, could also explain the formation of serpentine fields at the photosphere. \\subsection{For Future Observations \\label{sec:future}} In this paper, we found some aspects of the flux emergence event. One is the photospheric tongue, i.e., the magnetic structure extending horizontally around the surface just before further evolution takes place. Temporally- and spacially-resolved spectroscopic observations of the earlier phase of the flux emergence are required to study this magnetic extension at the photosphere. At the same time, we found that the initial twist of $qH_{0}\\sim 0.1$ at $-1700\\ {\\rm km}$ is too strong to match the observations (as mentioned above in Section \\ref{sec:twist}). Local and global helioseismology are needed to reveal the flux emergence (especially on the twist evolution) within the convection zone. The key issue is how weakly twisted flux tubes manage to rise through the solar interior. \\bigskip Numerical computations were carried out on NEC SX-9 at the Center for Computational Astrophysics, CfCA, of the National Astronomical Observatory of Japan, and on M System (Fujitsu FX1) of JAXA Supercomputer System. The page charge of this paper is partly supported by CfCA. S. T. and T. Y. thank Dr. Y. Fan of the High Altitude Observatory, the National Center for Atmospheric Research. We thank the referee for helpful suggestions for improvements of this paper." }, "1101/1101.0463_arXiv.txt": { "abstract": "We report on the results from the radial-velocity follow-up program performed to establish the planetary nature and to characterize the transiting candidates discovered by the space mission \\textit{CoRoT}. We use the SOPHIE at OHP, HARPS at ESO and the HIRES at Keck spectrographs to collect spectra and high-precision radial velocity (RV) measurements for several dozens different candidates from \\textit{CoRoT}. We have measured the Rossiter-McLaughlin effect of several confirmed planets, especially CoRoT-1b which revealed that it is another highly inclined system. Such high-precision RV data are necessary for the discovery of new transiting planets. Furthermore, several low mass planet candidates have emerged from our Keck and HARPS data. ", "introduction": "Transiting exoplanets are unique targets for which we can measure the planetary radius by high accuracy photometry when the planet passes in front of its host star and its mass and orbital characteristics (eccentricity, alignment with the spin of the star) by Doppler measurements of the host star. Thus it is possible to compute the mean density of the planet and to model its internal structure or to explore the composition and characteristics of its atmosphere (albedo, temperature, and atmospheric composition) by photometry or spectroscopy observations during the transit or the eclipse.\\\\ \\textit{CoRoT} (Baglin et al. 2006; Deleuil et al., this book) is the first space-based photometric survey dedicated to finding transiting exoplanets. \\textit{CoRoT} finds about 250 objects per run whose light curves show transit-like events. Most of them are clear eclipsing binaries (EB), but when a target shows periodic single transits (i.e. no secondary transits), no ellipsoidal variations and a shape, duration, and depth compatible with a transiting exoplanet, we consider it as a transiting exoplanet candidate. But these planetary candidates could still be mimicked by a transiting low-mass star, a grazing EB, a main sequence star eclipsing a giant star, or by an EB diluted by a third star. These EB scenarii (about 50\\% of candidates) could be resolved by high-resolution spectroscopy observations in order to discard all binary scenarii (SB1, SB2, SB3, ...). For example, Fig. \\ref{transitscenarii} shows the result of the CCF\\footnote[1]{Cross-Correlation Function between the spectrum and a numeric mask used as template of a star spectrum. It corresponds to the mean spectral line.} of a transiting candidate followed-up with SOPHIE. When EB scenarii are discarded, precise RV observations are required to measure the mass of the transiting object and characterize its orbit. \\begin{figure}[h!] \\begin{center} \\begin{minipage}[b]{.55\\textwidth} \\centering \\epsfig{file=Santerne_fig1a.eps, width=\\textwidth} \\end{minipage} \\hfill \\begin{minipage}[b]{.3\\textwidth} \\centering \\epsfig{file=Santerne_fig1b.eps, width=1\\textwidth} \\end{minipage} \\label{transitscenarii} \\caption{(left) RV variations as function of time of the planetary candidate LRc01\\_E2\\_1245 (Cabrera et al. 2009) revealed as a SB1 by SOPHIE with only two measurements. RV are compatible with a companion of $m_c \\sim 0.2 M_\\odot$ with a period of 4.974 days, assuming $M_*=1~M_\\odot$ and $i=90\\deg$ and $e=0$. (right) CCF$^*$ of a candidate revealed as a SB3 by SOPHIE. The third star have a large $v\\sin i$} \\end{center} \\end{figure} ", "conclusions": "Transiting exoplanet surveys need RV follow-up in order to determine the nature and the characteristics of the exoplanets candidates. Using the facilities of an optimized network of 3 high-resolution spectrographs for follow-up (SOPHIE, HARPS and HIRES) with powerful diagnostics to discard false positives and secure detection, \\textit{CoRoT} is, so far, the photometric survey that has discovered more planets per square degree of observed sky. More than 1000 spectra with signal-to-noise of up to 100 on about 200 transit candidates were taken with SOPHIE, HARPS, and HIRES during the first 3 years of \\textit{CoRoT}. Fifteen new exoplanets and brown dwarves have been discovered and characterized by these high-resolution spectrographs so far. Currently, 6 of these \\textit{CoRoT} planets have been observed in order to measure their RM effect. One planet orbit is is clearly misaligned with the spin of its host star while another one shows strong evidence of a misalignment, but requires more measurements to confirm this." }, "1101/1101.2699_arXiv.txt": { "abstract": "The Decadal Survey of Astronomy and Astrophysics created five panels to identify the science themes that would define the field's research frontiers in the coming decade. I will describe the conclusions of one of these, the Panel on Cosmology and Fundamental Physics, and comment on their relevance to the discussions at this meeting of the NASA Laboratory Astrophysics community. ", "introduction": "The \\cite{panel} convened five panels to consider the science themes that would define the field's research frontiers in the next decade. One of these, the Panel on Cosmology and Astrophysics, had a particularly broad mandate that included topics of interdisciplinary interest. The panel was chaired by David Spergel and included David Weinberg (vice chair), Rachel Bean, Neil Cornish, Jonathan Feng, Alex Filippenko, Marc Kamionkowski, Lisa Randall, Eun-Suk Seo, David Tytler, Cliff Will, and myself. The organizers of this Laboratory Astrophysics Workshop have asked me to summarize the Panel's conclusions and comment on their relevance to laboratory astrophysics and future NASA missions. The context for Panel discussions was established by a set of recent discoveries that have strengthened the links between astrophysics/cosmology and fundamental physics conducted in terrestrial laboratories. These include \\begin{list}{\\labelitemi}{\\leftmargin=0.25cm} \\item The development of a relatively simple cosmological model fitting astronomical data, Lambda Cold Dark Matter, with parameters known to better than 10\\% and with immediate implications for beyond-the-standard-model physics. \\item Cosmic microwave background (CMB) and large-scale structure (LSS) studies that appear consistent with the predictions of inflation: a nearly flat universe with a matter distribution that is Gaussian with nearly scale-invariant initial fluctuations. \\item CMB confirmation of the Big Bang nucleosynthesis (BBN) conclusion that baryons comprise about 4\\% of the closure density $\\Omega_c$, so that dark matter must be primarily nonbaryonic. \\item Supernova data indicating that the expansion of the universe is accelerating, consistent with dark energy dominance of the universe's present energy density. \\item Astrophysical $\\nu$ discoveries, from the Sun and from cosmic rays (CRs) impinging on Earth, that show neutrinos have mass and undergo flavor oscillations, providing the first direct evidence of physics beyond the standard model (and the first identification of a component of the dark matter). \\item The identification of a cutoff in ultra-high-energy (UHE) CRs consistent with the expected GZK scattering off the CMB. Thus the universe may be opaque to us at cosmological distances and asymptotic energies, apart from UHE $\\nu$s. \\end{list} The Panel considered community input, generally provided as white papers, on a wide range of topics: the early universe; the CMB; probes of LSS through observations of galaxies, intergalactic gas, or their associated gravitational distortions; determinations of cosmological parameters; dark matter; dark energy; tests of gravity; astrophysical measurements of physical constants; and the fundamental physics that might be derived from astronomical messengers ($\\nu$s, $\\gamma$s, CRs). Among the white papers considered, several addressed either laboratory astrophysics or theory and computation. The Panel's response was formulated around four ``big questions:\" 1) How did the universe begin (the mechanism behind inflation)? 2) Why is the universe accelerating (the nature of the dark energy)? 3) What is dark matter? 4) What are the properties of neutrinos? Gravitational wave astronomy was designated as the discovery area. This meeting's organizers have asked me to summarize the Panel's conclusions, commenting on their connections to laboratory astrophysics and NASA missions. In this written version of my talk I will focus the last two of the four questions, in part because I know these areas best, but also because they may have substantial connections to laboratory astrophysics. Here ``laboratory astrophysics\" is defined quite broadly, given that the Panel's charge included the intersection of astrophysics and astronomy with the particle and nuclear physics programs of major accelerator facilities, and with a broad array of ground-based detectors for dark matter, $\\nu$s, CRs, and related studies. \\section {What is Dark Matter?} The majority of matter in the universe is dark, invisible to us apart from its gravitational effects on the structure we do see. In addition to its deep roots in cosmology and astrophysics, the dark matter (DM) problem is central to high-energy physics, where DM particles may be discovered in the debris from collisions between ordinary particles, and in underground science, where the recoil of detector nuclei may indicate interactions with dark matter particles. (For a recent review of the topics summarized here, see \\cite{feng}.) DM was first postulated to account for the anomalous velocity rotation curves of galaxies. DM particles must be stable or long-lived, cold or warm (sufficiently slow that they can seed structure formation), gravitationally active, but without strong couplings to themselves or to baryons. The DM/dark energy contributions to the universe's total energy density evolves with redshift, with the former dominant early and the latter dominant today. Two leading DM candidates are Weakly Interacting Massive Particles (WIMPs) and axions. WIMPS are intriguing because the properties necessary for astrophysics match expectations that new particles will be found at the mass generation scale of the standard model of 10 GeV - 10 TeV. The WIMP ``miracle\" is the observation that the annihilation cross section for massive, weakly interacting particles natural leads to the expectation that $\\Omega_\\mathrm{WIMP} \\sim 0.1.$ Figure \\ref{fig:Wimp} illustrates three avenues of attack on the DM problem: direct detection where a DM particle $\\chi_\\mathrm{DM}$ scatters off a standard-model particle f$_\\mathrm{SM}$, causing recoil; indirect detection, where DM particles annihilate or decay into ordinary particles that can then be detected; and particle collider experiments, where DM particles are produced from the scattering of ordinary particles and identified from the missing energy. DM particle interactions can be either independent (SI) or dependent (SD) on target spin, depending on parameter choices in the underlying model, and while their predicted cross sections span a wide range, $\\sigma_{SI} \\sim 10^{-45}$ cm$^2$ is a representative value. Current detectors in the $\\sim$ 10-100 kg range are probing DM particle-nucleon cross sections well below $10^{-43}$ cm$^2$ for DM particle masses of $\\sim$ 100 GeV. The international program is focused on developing new detectors in the 1-10 ton range, using media such as ultra-clean nobel-gas liquids, with sensitivities of a few events per year, or $\\sigma_{SI} \\sim 10^{-47}$ cm$^2$. There have been claims of detection, interpreted as low-mass WIMPs, but no consensus has been reached. This field has a number of laboratory and theory needs, including: \\begin{list}{\\labelitemi}{\\leftmargin=0.25cm} \\item support for direct searches, including detector R\\&D and the development of deep underground locations for detectors, as energetic neutrons produced by penetrating CR muons are an important background; \\item clean-room facilities to control environmental activities, as trace radionuclides within detectors or in surrounding materials are a second major background source; \\item nuclear theory for estimates of WIMP SI form factors and SD cross sections; \\item for direct production experiments, facilities like the LHC that can reach the energies necessary for $\\chi_\\mathrm{DM}$ creation; and \\item for indirect detection searches, astrophysical modeling that will allow observers to distinguish WIMP annihilation signals from other high-energy astrophysics phenomena. \\end{list} Ideally, multiple lines of investigation will lead to DM detection. An attractive scenario is the discovery of supersymmetry at the LHC and the identification of a lightest stable SUSY particle; direct detection of cosmic WIMPS with consistent properties; and consequently, improved constraints on the local DM density and its effects on structure at subgalactic scales, testing the paradigm of cold, collisionless, stable DM. \\begin{figure}[ht] \\centering \\includegraphics[angle=0,scale=.50]{Wimp} \\caption{Three strategies for detecting dark matter particles $\\chi_\\mathrm{DM}$ via their interactions with standard-model particles f$_\\mathrm{SM}$.\\label{fig:Wimp}} \\end{figure} ", "conclusions": "The laboratory astrophysics $\\leftrightarrow$ astrophysics/cosmology intersection described in this talk is sometimes conventional -- e.g., the nuclear cross section measurements done in support of solar $\\nu$ experiments or BBN modeling -- but more often unconventional. For example, in determining the overall scale of neutrino mass, the experiments done in the laboratory (tritium $\\beta$ decay, $\\beta \\beta$ decay, reactor and accelerator $\\nu$ oscillation searches) complement what can be done in cosmology. Each probes an important quantity -- tritium $\\beta$ decay probes the masses of $\\nu$s in proportion to their coupling probability to the electron, $\\beta \\beta$ decay probes the Majorana mass, oscillations test $m_\\nu^2$ differences, and cosmology responds to the sum of the $\\nu$ masses -- but the quantities are different. So the traditional relationship has given way to one where the universe has be viewed as another laboratory, one that is playing a very important role in pushing back the frontiers of precision particle physics. This is a theme that may have been best expressed in another NRC study, From Quarks to the Cosmos. In the coming decade the community hopes to build new $\\nu$ beamlines and underground detectors on the $10^2-10^3$ kiloton scale to determine the hierarchy through matter effects, measure $\\nu$ CP violation, and fix all mixing angles and mass differences to high precision. Cosmology could be an equal partner in this effort: if the effects of $\\nu$ mass on LLS can be determined to sufficient accuracy, both the absolute scale and hierarchy questions might be resolved in this way. In DM the situation is similar. One of the goals of the LHC is to find the new particles that are expected to accompany TeV-scale physics. We also have a new generation of massive DM detectors that will be mounted on Earth (or more precisely, within the Earth) that will be probing the cosmological flux of DM particles. One hopes that both endeavors succeed, and that we will be able to reconcile the cosmological properties of DM with those determined from collider experiments." }, "1101/1101.2650_arXiv.txt": { "abstract": "We use numerical hydrodynamic simulations to investigate prestellar core formation in the dynamic environment of giant molecular clouds (GMCs), focusing on planar post-shock layers produced by colliding turbulent flows. A key goal is to test how core evolution and properties depend on the velocity dispersion in the parent cloud; our simulation suite consists of 180 models with inflow Mach numbers ${\\cal M}\\equiv v/c_s =1.1-9$. At all Mach numbers, our models show that turbulence and self-gravity collect gas within post-shock regions into filaments at the same time as overdense areas within these filaments condense into cores. This morphology, together with the subsonic velocities we find inside cores, is similar to observations. We extend previous results showing that core collapse develops in an ``outside-in'' manner, with density and velocity approaching the Larson-Penston asymptotic solution. The time for the first core to collapse depends on Mach number as $t_{\\rm coll} \\propto {\\cal M}^{-1/2}\\rho_0^{-1/2}$, for $\\rho_0$ the mean pre-shock density, consistent with analytic estimates. Core building takes 10 times as long as core collapse, which lasts a few $\\times 10^5$ yrs, consistent with observed prestellar core lifetimes. Core shapes change from oblate to prolate as they evolve. To define cores, we use isosurfaces of the gravitational potential. We compare to cores defined using the potential computed from projected surface density, finding good agreement for core masses and sizes; this offers a new way to identify cores in observed maps. Cores with masses varying by three orders of magnitude ($\\sim 0.05 - 50 M_\\odot$) are identified in our high-$\\cal M$ simulations, with a much smaller mass range for models having low $\\cal M$. We halt each simulation when the first core collapses; at that point, only the more massive cores in each model are gravitationally bound, with $E_{\\rm th} + E_g <0$. Stability analysis of post-shock layers predicts that the first core to collapse will have mass $M \\propto v^{-1/2} \\rho_0^{-1/2} T^{7/4}$, and that the minimum mass for cores formed at late times will have $M\\propto v^{-1} \\rho_0^{-1/2} T^2$, for $T$ the temperature. From our simulations, the median mass lies between these two relations. At the time we halt the simulations, the $M$ vs. $v$ relation is shallower for bound cores than unbound cores; with further evolution the small cores may evolve to become bound, steeping the $M$ vs. $v$ relation. ", "introduction": "Star formation begins with the creation of dense molecular cores, and understanding how cores grow and evolve is essential to identifying the origin of stellar properties \\citep{shu87,mcke07,andr08}. Through the 1990s, the prevailing theoretical picture was of slow core formation and evolution mediated by ambipolar diffusion, followed by core collapse initiated from a quasistatic, centrally-concentrated state \\citep[e.g.,][]{mous87,mous99}. Current observations, however, indicate that magnetic field strengths are insufficient to provide the dominant support of molecular cores \\citep{trol08}. In addition, over the past decade, a conception of star formation has emerged in which supersonic turbulence drives structure and evolution within giant molecular clouds (GMCs) on a wide range of scales \\citep[e.g.,][]{ball07, mcke07}. Because supersonic turbulence can compress gas to densities at which gravitational collapse can rapidly occur, it is likely to be important in the initiation of prestellar cores. Ultimately, models of core formation and evolution must take into account both moderate magnetic fields (with diffusion) and strong turbulence \\citep{kudo08,naka08}. In order to gain insight into the physics involved, however, it is informative to focus on individual limiting cases and explore dependence on parameters. Here, following \\citet{gong09} but generalizing to three dimensions, we consider core building and evolution in the turbulence-dominated, unmagnetized limit. Observations of dense cores in GMCs have provided detailed information on individual core properties as well as statistics of core populations \\citep[see e.g., the reviews of][]{difr07,ward07,berg07,andr08}. These properties, including internal structure and kinematics, durations of different evolutionary stages, and distribution of core masses, constrain core formation theories. In terms of structure, cores are observed to be centrally concentrated at all stages, with the specific profile fits differing depending on the stage of evolution. Cores can generally be fit with a uniform-density inner region surrounded by a power law $\\propto r^{-2}$ \\citep[e.g.,][]{shir00,bacm00, alve01,kand05,kirk05}; this shape is consistent with expectations for both static Bonnor-Ebert (BE) pressure-supported isothermal equilibria \\citep{bonn56,eber55}, and for collapsing isothermal spheres \\citep{ bode68,lars69,pens69}. The center-to-edge density contrast is frequently larger than the maximum possible for a stable BE sphere, however, and the inferred temperatures based on static BE fits are also often larger than observed temperatures. Although in principle some support could be provided by magnetic fields \\citep[e.g.,][]{ciol94}, another possibility is that these ``supercritical'' cores are in fact collapsing rather than static \\citep{dapp09,gong09}. In terms of kinematics, dense, low-mass cores generally have subsonic internal velocity dispersions, whether for isolated cores or for cores found in clusters \\citep[e.g.,][]{myer83,good98,case02,tafa04, kirk07,andr07,lada08}. Some prestellar cores also show indications of subsonic inward motions throughout their interiors based on asymmetry of molecular lines that trace dense gas \\citep[e.g.,][]{lee99,lee01,sohn07}. For cores containing protostars, signatures of supersonic inward motions on small scales ($\\sim 0.01 - 0.1 \\pc$) have been observed \\citep[e.g.,][]{greg97,difr01}; these are believed to be indicative of gravitationally-induced infall. In very recent work, \\citet{pine10} have used $\\mathrm{NH}_3$ observations to identify a sharp transition from supersonic to subsonic velocity dispersion from outer to inner regions in the core B5 in Perseus. Several recent statistical studies have reached similar conclusions regarding the durations of successive stages of core evolution \\citep[e.g.,][]{ward07, enoc08,evan09}, with prestellar and protostellar (class 0) stages having comparable lifetimes. The typical duration for each of these stages is a few times the gravitational free-fall time \\begin{equation}\\label{t_ff} t_{ff} = \\left(\\frac{3 \\pi}{32 G \\bar{\\rho}} \\right)^{1/2} = 4.3 \\times 10^5 \\yr \\left(\\frac{\\bar{n}_H}{10^4 \\pcc}\\right)^{-1/2} \\end{equation} at the mean core density $\\bar{\\rho} = 1.4 m_H \\bar{n}_H$, amounting to $\\sim$ 1 -- 5 $\\times 10^5$ yr for typical conditions. With prestellar lifetimes considerably below the ambipolar diffusion time for strong magnetic field $t_{AD} \\approx 10 t_{ff}$ \\citep[e.g.][]{mous99}, this suggests that observed cores are trans-critical or supercritical \\citep[see][]{ciol01} with respect to the magnetic field.\\footnote{ The critical mass-to-magnetic-flux defines the minimum that permits gravitational collapse in the field-freezing limit \\citep[e.g.][]{mest56,mous76,naka78}.} This conclusion is also supported by magnetic field Zeeman observations \\citep{trol08}, indicating that cores have mean mass-to-magnetic-flux ratios twice the critical value. Thus, magnetic field effects appear to be sub-dominant in terms of supporting cores against collapse, and ambipolar diffusion does not appear to control the dynamics of core formation and evolution. As magnetic fields are non-negligible, however, magnetohydrodynamic (MHD) stresses may still affect GMC and core dynamics. Empirical measurements of core mass functions (CMFs) \\citep[e.g.,][]{mott98,test98,john00, john01,mott01,onis02,beut04,reid05,reid06,stan06,enoc06,alve07,iked07,iked09b, iked09a,nutt07,simp08,kony10} show that CMFs have a remarkable similarity in shape to stellar initial mass functions (IMFs, see e.g. \\citealt{krou01, chab05}), with a shift toward lower mass by a factor of 3 -- 4 \\citep[see e.g.,][]{alve07,rath09}. The characteristic/turnover mass of observed CMFs ranges from 0.1 -- 3 $\\Msun$, although there are uncertainties in this associated with lack of spatial resolution at the low mass end. Many theoretical efforts have contributed to interpreting the observed properties of cores. The classic work of \\citet{bonn56} and \\citet{eber55} provided the foundation of later studies, by determining the maximum mass of a static isothermal sphere that is dynamically stable. In terms of the boundary pressure $P_\\mathrm{edge} = \\rho_\\mathrm{edge} c_s^2$ or mean internal density $\\bar{\\rho} = 2.5 \\rho_\\mathrm{edge}$, this maximum stable mass is \\begin{equation}\\label{m_be} M_{BE} = 1.2 \\frac{c_s^4}{(G^3 P_{\\mathrm{edge}})^{1/2}} = 1.9 \\frac{c_s^3}{(G^3\\bar{\\rho})^{1/2}} = 2.3\\Msun \\left(\\frac{\\bar{n}_H}{10^4 \\pcc} \\right)^{-1/2} \\left(\\frac{T}{10K}\\right)^{3/2}. \\end{equation} Here, $c_s = (kT/\\mu)^{1/2}$ is the internal sound speed in the core. Over many years, numerical simulations have been used to investigate isothermal collapse of individual, pre-existing cores \\citep{bode68,lars69,pens69,hunt77, fost93, ogin99,henn03,moto03,voro05,gome07,burk09}. These simulations include initiation from static configurations that are unstable, and initiation from static, stable configurations that are subjected to imposed compression, either from enhanced external pressure or a converging velocity field, or a core-core collision. A common feature of the results is that the collapse generally starts from outside and propagates in as the central density increases. At the time of singularity formation, the density profile approaches the ``Larson-Penston'' asymptotic solution $\\rho = 8.86 c_s^2/(4 \\pi G r^2)$ and the central velocity is comparable to the value $-3.28 c_s$ derived by \\citet{lars69} and \\citet{pens69}. However, these previous studies have not considered core evolution within the larger context, in particular including the process of \\emph{core formation}. Since the formation process may affect later evolution, it is important to develop unified models. At GMC scales, a number of groups have investigated the CMFs that result from numerical simulations of turbulent, self-gravitating systems \\citep[see e.g.,][]{kles01,gamm03,bonn03,li04,till04,heit08,clar08,offn08,basu09,smit09}. These models have shown -- for certain parts of parameter space -- features that are in accord with observed CMFs: mass functions dominated by the low end with a peak and turnover near $1 \\Msun$, and a high-mass power-law slope (at least marginally) consistent with the Salpeter value. These simulations have not, however, had sufficient resolution to investigate the internal properties of individual cores that form. In addition, these studies have not quantified how the core masses depend on the large-scale properties of the turbulent medium (see below). Taking the previous numerical simulations of individual cores one step further, \\citet{gong09} initiated a study of dynamically induced core formation and evolution in supersonic converging flows, focusing on the spherical case. In these simulations, the density is initially uniform everywhere: no initial core structure is assumed. Instead, dense cores form inside a spherical shock that propagates outward within the converging flow. Over time, cores become increasingly stratified as their masses grow. Eventually, the core collapses to create a protostar following the same ``outside-in'' pattern as in models initiated from static conditions. Subsequently, the dense envelope falls into the center via an inside-out rarefaction wave \\citep{shu77,hunt77}; this is followed by a stage of late accretion if the converging flow on large scales continues to be maintained. The unified formation and evolution model of \\citet{gong09} explains many observed core properties, including BE-sphere-like density profiles, subsonic internal velocities within cores, and short core lifetimes with comparable prestellar and protostellar durations. \\citet{gong09} also found that the inflow velocity of the converging flows affects core lifetimes, masses, sizes and accretion histories. Realistic supersonic inflows in clouds are not spherical, however, while mass inflow rates are affected by geometry. Thus, the quantitative results for masses, lifetimes, etc., as a function of Mach number and ambient density may differ for more realistic geometry. Numerical results on core formation have not reached consensus on how the characteristic mass in the CMF, $M_c$, depends on the bulk properties of the cloud -- its mean density $\\rho_0 = \\langle \\rho \\rangle$, sound speed $c_s$, and turbulent velocity dispersion $v_{\\mathrm{turb}}$. Some have suggested that the Jeans mass of the cloud at its mean density ($M_J = c_s^3 \\pi^{3/2} (G^3 \\rho_0)^{-1/2}$) determines $M_c$ in the CMF (e.g,. \\citealt{kles01,bonn06}), while others have found values of $M_c$ well below $M_J$ \\citep[see e.g.,][]{gamm03,li04}. As noted by \\citet{mcke07}, the difference between these conclusions is likely related to the Mach number of turbulence: the value found for $M_c/M_J$ is lower in simulations where the Mach number $\\mathcal{M} \\equiv v_{\\mathrm{turb}}/c_s$ is higher. Indeed, more recent simulations by \\citet{clar08} provide some indication that increasing $\\mathcal{M}$ lowers the value of $M_c$ in the CMF; they did not, however, conduct a full parameter study. Supersonic turbulence makes the density in a GMC highly non-uniform, creating a log-normal probability distribution function (PDF) in which most of the volume is at densities below $\\rho_0$ and most of the mass is at densities above $\\rho_0$ \\citep[e.g.,][]{vazq94,pado97,ostr99}. Given that the log-normal PDF allows for a range of Jeans masses (or Bonnor-Ebert masses; $M_{BE} \\propto M_J$), \\citet{pado02,pado04} proposed that the CMF is set by dividing the total available gas mass at each density into unstable cores. \\citet{pado07} propose that the peak mass in the CMF is given by $M_c = 3 M_{BE,0}/M_A^{1.1}$ for $M_A \\equiv v_{\\mathrm{turb}}/v_A$ the Alfv$\\acute{\\mathrm{e}}$n Mach number in a cloud, and $M_{BE,0}$ the Bonnor-Ebert mass evaluated at the mean cloud density $n_0$. Here, $v_A \\equiv B/(4 \\pi \\rho)^{1/2}$ is the Alfv$\\acute{\\mathrm{e}}$n speed. For realistic mean GMC density $n_0 \\sim 100 \\pcc$ and $\\mathcal{M}_A \\sim 1 -4$, from Equation (\\ref{m_be}) the Padoan et al formula in fact yields $M_c > 15 \\Msun$; only if one chooses a much higher reference density does this agree with observations. For the unmagnetized case, \\citet{pado07} propose that $M_c = 4 M_{BE,0}/\\mathcal{M}^{1.7}$. \\citet{henn08} point out that shock compression is underestimated in the magnetized case by \\citet{pado07}, and advocate a formula similar to their unmagnetized one: $M_c \\sim M_{BE,0}/\\mathcal{M}^{3/2}$. Since $\\mathcal{M} \\gtrsim 10$ in massive GMCs, these formulae yield more realistic values $M_c \\sim\\Msun$. Neither the \\citet{pado07} or the \\citet{henn08} proposal has, however, been tested directly using self-gravitating numerical simulations. In this contribution, we present results on core formation and evolution based on a large suite of 3-dimensional numerical simulations. Each simulation models a localized region of a turbulent cloud in which there is an overall convergence in the velocity field. Under the assumption that there is a dominant convergence direction locally, we choose inflow along a single axis, so that convergence is planar. With the more realistic geometry afforded by the current simulations, we are able to check the results obtained by \\citet{gong09} for core building and collapse in supersonic flows. We are also able to explore how the characteristic core mass is related to the velocity of the converging flows. Since the speed of converging flow is assumed to reflect the amplitude of the largest-scale (dominant) motions in a GMC, this relates the characteristic core mass to the turbulent Mach number in its parent GMC. Although a number of previous studies of core formation have been conducted, the present investigation is distinguished by our systematic study of Mach number dependence, together with our focus on internal structure and kinematics of the cores that form. The plan of this paper is as follows: In Section 2 we provide a physical discussion of self-gravitating core formation in the post-shock dense layers, identifying the mass, size, and time scales expected to be important. In Section 3, we summarize the governing equations and methods used in our numerical simulations. Section 4 describes the development of core structure and evolution in our models, paying particular attention to the influence of Mach number $\\mathcal{M}$ on the evolution, and comparing collapse of individual cores with \\citet{gong09}. Section 5 describes our method of core-finding, in which the largest closed contour of the gravitational potential determines the core size. We demonstrate that this method can be used for both three dimensional and two dimensional data with similar results, and can thus be applied to find cores in observed clouds. Section 6 describes the relations between core properties (core mass, core radius and core collapse time) and the large-scale Mach number of the converging flow, relating to the expectations from gravitational instability discussed in Section 2. In Section 6, we also quantify core shapes, and explore the relationship between core structure and kinematics. Section 7 summarizes our new results and discusses our findings in the context of previous theories and observations. ", "conclusions": "Stars form in GMCs pervaded by supersonic turbulence, and core formation theory must take these supersonic turbulent flows into account. In this work, we explore the physics of core formation in a dynamic environment, focusing on post-shock layers generated by collisions of supersonic flows. The framework we adopt -- three-dimensional planar converging flows containing multi-scale turbulence -- enables us to analyze the internal structure and kinematics of cores, and to investigate the relation between core properties and the inflow Mach number $\\mathcal{M}$. We consider a range $\\mathcal{M} = $ 1.1 -- 9, and conduct 180 simulations with different realizations of the initial turbulent power spectrum, in order to obtain a sizable statistical sample. In addition to core masses and sizes, we measure aspect ratios. To define cores, we introduce a new method based on the gravitational potential, and compare properties of cores identified using $\\Phi$ (from the volume density) and $\\Phi_\\mathrm{2D}$ (from the plane-of sky projected surface density). Unlike previous studies of core evolution that begin with pre-existing cores, the present models include formation stages. Our initial density is uniform everywhere, and cores grow, via self-gravity, from turbulence-induced perturbations within the post-shock layer; when the Mach number is high, initial growth of density perturbations is aided by shock-driven hydrodynamic instabilities. Based on a set of spherically-symmetric numerical simulations, \\citet{gong09} proposed four stages for core evolution in dynamic environments: core building, core collapse, envelope infall, and late accretion. The key features during core building and collapse described in \\citet{gong09} are verified here, for more realistic geometry. As the supersonic flows converge in a plane, two reversed shocks propagate outwards. With its high mean density, the stagnation layer between these two shock fronts becomes an incubator for self-gravitating cores. When these cores become sufficient stratified, they collapse. We halt the simulations at the instant of singularity formation in the most evolved core, because the time step becomes very short. Based on the analysis of our simulations, our chief conclusions are as follows: 1. Cores with realistic properties are able to form in post-shock dense layers within turbulent GMCs. Core building to become supercritical takes $\\sim$ 10 times as long as the subsequent ``outside-in'' collapse stage, which lasts a few $\\times 10^5 \\yr$. The duration of the supercritical stage is consistent with observations of prestellar core lifetimes \\citep{ward07, enoc08, evan09}. 2. At the time of singularity formation, the radial density profile within cores approaches the Larson-Penston asymptotic solution $\\rho = 8.86 c_s^2/(4 \\pi G r^2)$ and the velocity approaches the Larson-Penston limit $-3.28 c_s$. This is consistent with previous studies of spherical core collapse (see Section 1 for references). \\citet{till04} also found that $\\rho \\propto r^{-2}$ in their most massive cores, for turbulent simulations. As in \\citet{gong09}, we therefore conclude that the Larson-Penston asymptotic solution is an ``attractor'' for core collapse, no matter how the collapse is initiated. 3. Prior to collapse, the velocities within dense cores remain subsonic, in spite of the highly-supersonic flows that create them. This is true both for the ordered inflow, and for the mean internal velocity dispersion. This result is consistent with observations that most cores have subsonic non-thermal velocity dispersions \\citep{myer83,good98,case02,tafa04, kirk07,andr07,lada08}. The velocity dispersion can increase quite sharply at the edge of the core in our models (see Fig. \\ref{fig:v_los}), intriguingly similar to a sharp transition seen in $\\mathrm{NH}_3$ observations by \\citet{pine10} for the B5 core in Perseus. From some orientations, velocity dispersions in filaments containing cores may also be lower than in the surrounding gas (cf. Fig. \\ref{fig:v_los}). 4. At sub-pc scales, turbulent velocity perturbations (whether super- or subsonic) induce density perturbations that can grow strongly if the density is high enough for self-gravity to be important. In post-shock layers, turbulence and self-gravity collect gas into long, thin filamentary structures at the same time as the highest density regions within the filaments grow to become centrally-condensed cores. These filamentary structures containing embedded cores are similar to the structures in the Aquila rift and Polaris Flare clouds observed by {\\it Herschel} \\citep{andr10,mens10}. 5. Using the gravitational potential to identify cores is advantageous because it enables a core definition based on dynamical principles. For numerical simulations, the gravitational potential may be computed from the volume density (yielding $\\Phi$) or from the projected surface density (yielding $\\Phi_\\mathrm{2D}$). We show for our models that cores defined using $\\Phi$ and $\\Phi_\\mathrm{2D}$ are nearly the same, both for GRID-cores (defined by the largest closed potential isosurfaces) and bound GRID-cores (which additionally require $E_{\\rm th} + E_g <0$). Since $\\Phi_\\mathrm{2D}$ can be computed for observed clouds, using potential contours offers a promising new core identification method for application to high-resolution molecular cloud maps. IDL code implementing our GRID-core algorithm, suitable for application to observed data, is available from the authors. 6. We find that the range of core masses that forms increases as the Mach number $\\mathcal{M}$ increases. Physically, this is because a larger range of spatial scales has significant perturbations when the turbulence amplitude is higher, and because the minimum mass to be gravitationally unstable decreases as the density in the shocked layer increases. \\citet{basu09} also found broader mass distributions when the turbulent amplitude is increased. At high Mach number, GRID-core masses range between $\\sim 10^{-3}$ -- $1 M_J$, corresponding to $\\sim 0.05$ -- $50 \\Msun$ for typical GMC conditions. 7. Analytical arguments (see Section 2) suggest that the first core to collapse will have mass $M \\propto \\mathcal{M}^{-1/2}$, and that at late times, the minimum mass core will vary as $M \\propto \\mathcal{M}^{-1}$. Our numerical results for median core masses as a function of $\\mathcal{M}$ lie between these two relations. When the core definition includes the condition that $E_{\\rm th} + E_g < 0$, the median mass increases at the largest Mach number. This may be due to the nonlinear ``head start'' of massive cores, such that lower mass cores have not yet become concentrated when the first core collapses (and the simulation is stopped). 8. Analytical arguments (see Section 2) suggest that the effective core radius will decline with increasing Mach number, with powers between $r_\\mathrm{eff} \\propto \\mathcal{M}^{-1/2}$ and $r_\\mathrm{eff} \\propto \\mathcal{M}^{-1}$. Our numerical results show a decrease of $r_\\mathrm{eff}$ with $\\mathcal{M}$ in this range. For bound GRID-cores ($E_{\\rm th} + E_g < 0$), the relation is shallower than for GRID-cores defined by gravitational potential alone. 9. The time for the first core to collapse in our simulations depends on Mach number, with $t_\\mathrm{coll} \\propto \\mathcal{M}^{-1/2}$, and a slightly smaller coefficient for high-amplitude initial perturbations (see Fig. \\ref{fig:time_comp}). This scaling is consistent with analytic predictions for gravitational instability in a shocked converging flow (see eq. \\ref{t_mgrn}). For high $\\mathcal{M}$, as is observed in GMCs, the first cores could collapse within a few Myr of cloud formation. For high $\\mathcal{M}$, the first cores collapse when the shocked layer containing them is only barely self-gravitating; this suggests that collections of stars can begin to form individually before they collapse together to create a cluster. 10. A very small portion of cores are oblate, while most cores are prolate or triaxial. Large cores are preferentially prolate. The triaxiality of most cores is consistent with previous results from turbulent hydrodynamic and MHD simulations \\citep{gamm03,li04,naka08,offn08}. We also find that core shapes change as they evolve, from more oblate during early stages to more prolate during collapse. For high initial perturbation amplitudes, the distributions have a higher proportion of oblate cores because small cores are less evolved (at the time the first core collapses), compared to those in models with low initial perturbation amplitudes. As noted above, the current models have provided evidence that the masses of cores that form depend not just on the mean Jeans mass in a cloud, but also on the cloud's level of internal turbulence at large scales, $\\sigma_v$. Equations (\\ref{m_crit_sg1}) and (\\ref{m_crit_sg3}) suggest that at late times, the characteristic core mass will follow $M_c \\propto \\sigma_v^{-1}\\rho_0^{-1/2}T^2$, where $\\rho_0$ is the mean density in the cloud. For the current simulations, however, we halt at the instant when the most evolved core collapses (because the time step becomes very short). This limits the condensation of small cores; they are present, but not yet strongly bound. In order to fully test the dependence of $M_c$ on cloud parameters, it is necessary to implement sink particles \\citep[e.g.][]{krum04,fede10} so that the simulation can run until all the ``eligible'' cores in the post-shock region have had the opportunity to collapse. Including sink particles, as well as studying shocked converging flows within larger turbulent clouds via mesh-refined simulations, represent important avenues for future research." }, "1101/1101.0641_arXiv.txt": { "abstract": "How to properly understand coronal mass ejections (CMEs) viewed in white-light coronagraphs is crucial to many relative researches in solar and space physics. The issue is now particularly addressed in this paper through studying the source locations of all the 1078 LASCO CMEs listed in CDAW CME catalog during 1997 -- 1998 and their correlation with CMEs' apparent parameters. By manually checking LASCO and EIT movies of these CMEs, we find that, except 231 CMEs whose source locations can not be identified due to poor data, there are 288 CMEs with location identified on the front-side solar disk, 234 CMEs appearing above solar limb, and 325 CMEs without evident eruptive signatures in the field of view of EIT. Based on the statistical results of CMEs' source locations, four physical issues, including (1) the missing rate of CMEs by SOHO LASCO and EIT, (2) the mass of CMEs, (3) the causes of halo CMEs and (4) the deflections of CMEs in the corona, are exhaustively analyzed. It is found that (1) about 32\\% of front-side CMEs can not be recognized by SOHO, (2) the brightness of a CME at any heliocentric distance is roughly positively correlated with its speed, and the CME mass derived from the brightness is probably overestimated, (3) both projection effect and violent eruption are the major causes of halo CMEs, and especially for limb halo CMEs, the latter is the primary one, (4) most CMEs deflected towards equator near the solar minimum, and these deflections can be classified into three types, the asymmetrical expansion, non-radial ejection, and the deflected propagation. ", "introduction": "Coronal mass ejections (CMEs) are recognized as transient bright features in the field of view (FOV) of white-light coronagraphs. However, their apparent properties/behaviors manifested in coronagraphs may not reflect what the CMEs actually should be, as observations of coronagraphs have at least three intrinsic limitations. The first one comes from the projection effect. All the three-dimensional information is embedded in two-dimensional images. Thus the position or speed of a CME measured in coronagraphs is only the projection of real position or speed on the plane of the sky, the shape of a CME depends on the angle of view, and the brightness recorded is an integral of the photons scattered by free electrons along the line-of-sight. The second one, we called occulting effect, is due to the occulting disk, which is used by coronagraphs to block the photons directly emitted from the photosphere. It was clearly pointed out by \\citet{Howard_etal_1982} that two identical CMEs originating from the solar limb and disk-center, respectively, will look much different. The time and heliocentric distance of the disk-center CME entering the FOV of a coronagraph will be later and farther than those of the limb CME. It will further cause the disk-center CME fainter and diffuser than the limb CME. The third one is because of the Thomson scattering effect \\citep[e.g.,][]{Hundhausen_1993, Andrews_2002, Vourlidas_Howard_2006}. This effect results in a so-called Thomson sphere, on which the plasma material is the most visible. Moreover, in most popular coronagraph images, the inner corona is hidden behind the occulting disk. For example, the occulting disk size of the coronagraph LASCO/C2 onboard the SOHO spacecraft is 2 $R_S$, and it is 1.4 $R_S$ for the coronagraph COR1 onboard the STEREO twin spacecraft. Thus, we are blind to the CME behavior in the region covered by the occulting disk, where the CME propagation trajectory may change significantly. Here, we use the term `deflection' for the behavior of CME's non-radial ejection and/or propagation. It is an important factor for space weather. As early as \\citeyear{MacQueen_etal_1986}, \\citeauthor{MacQueen_etal_1986} had found the CME deflections in latitudinal direction by measuring 29 CMEs observed by the Skylab. \\citet{Gopalswamy_etal_2000a} discussed the non-radial propagation of the 1997 December 14 CME and pointed out that such a phenomenon clearly implied the constraint of the complex multi-polar structures surrounding the CME \\citep{Webb_etal_1997, Gopalswamy_etal_2004}. With more CME events detected by LASCO during 1996 to 2002, {\\it Cremades} and coworkers carried out a statistical study on their defined `structured' CMEs. They found that many CMEs do not propagate radially with respect to their source locations, and the neighboring and/or polar coronal holes played a major role in causing the deflections of CMEs \\citep{Cremades_Bothmer_2004, Cremades_etal_2006}. The presence of these effects requires us to be very careful when we interpret the observed bright features in coronagraphs. Only white-light images from coronagraphs are not enough. The information of the solar source locations of all CMEs is necessary. There have been some efforts except for the previously mentioned work about CME deflections. \\citet{Yashiro_etal_2005} investigated 1301 X-ray flares with intensity larger than C3 and their associations with CMEs, and found that about 14\\% of white-light CMEs were missed by LASCO. The statistically study of 9224 LASCO CMEs from 1996 to 2004 by \\citet{Lara_etal_2006} suggested that halo CMEs are different from normal CMEs, which can not be merely explained by projection effect, and the brightness of halo CMEs probably includes their driven (shock) waves. We acknowledge that these previous studies have advanced our understanding of the white-light CMEs observed by coronagraphs, but it is not comprehensive. We also realize that there are few works identifying the source locations of all CMEs no matter whether the CME is halo or narrow, strong or faint. Most studies involving the information of source locations considered halo CMEs only \\citep[e.g.,][]{Wang_etal_2002a, Zhou_etal_2003, Zhao_Webb_2003}. Some others set certain criteria in the selection of CMEs. For example, the study by \\citet{Subramanian_Dere_2001} only included the 32 CMEs with very clear EUV signatures on the solar surface. \\citet{Cremades_Bothmer_2004} selected so called `structured' CMEs, in which halo, narrow or faint CMEs are all excluded. \\citet{Yashiro_etal_2005} work involved the CMEs associated with flares above C3 level. To our knowledge, the study by \\citet{Plunkett_etal_2001}, might be the only statistical work, in which all the CMEs during the period of interest, which is from 1997 April to December, were identified for their source locations. An incomplete or biased sample may lead to unreliable or one-sided results, particularly, based on observations of coronagraphs, which have some intrinsic limitations. In this paper, we will identify the source locations of all the 1078 CMEs from 1997 to 1998 listed in the CDAW CME catalog\\footnote{A widely-used manually-compiled catalog, refer to \\url{http://cdaw.gsfc.nasa.gov/CME_list/}} \\citep{Yashiro_etal_2004}, and try to better understand CMEs viewed in white-light coronagraphs. Except for the statistical results of CMEs' source locations, our investigation will address the following four issues. \\begin{enumerate} \\item {\\it Missing rate of CMEs.} How many CMEs were missed by LASCO and EIT, and how many front-side CMEs were unnoticed by SOHO? \\item {\\it Mass of CMEs.} Whether or not can the enhanced brightness in coronagraphs reflect the CME mass? \\item {\\it Causes of halo CMEs.} Why do some CMEs manifest a halo appearance? \\item {\\it Deflections of CMEs.} How often and significant are CMEs deflected in the corona and why? \\end{enumerate} The period of 1997 -- 1998 is the beginning of the ascending phase of solar cycle 23, during which the solar condition is relatively simple and the solar activity level is low. Thus the source locations of CMEs are relatively easy to be identified with small ambiguity. This paper is organized as follows. In the next section, we present our data source, and particularly focus on the identification and classification of the CME source locations. The statistical results of the CME source locations are shown in Sec.\\ref{sec_source}. In Sec.\\ref{sec_implications}, the four issues mention above are extensively discussed. A summary and conclusions are given in Sec.\\ref{sec_conclusions}. ", "conclusions": "By manually checking the LASCO and EIT movies of all the 1078 CMEs listed in CDAW CME catalog during 1997 -- 1998, the solar surface sources of these CMEs are identified, and a web-based on-line list of them with the information of their source locations is established at \\url{http://space.ustc.edu.cn/dreams/cme_sources}. The source locations and apparent properties of CMEs have the following features. The distribution of CME source locations in latitude manifests a clear bimodal appearance with two most probable peaks in $\\pm(15^\\circ-30^\\circ)$, which is consistent with the location of active region belt. No CMEs came from polar regions (outside of $\\pm75^\\circ$). About 56\\% of detected CMEs occurred near the solar limb (refer to Sec.\\ref{sec_dis}). The average apparent speed of CMEs is about 435 km s$^{-1}$, and there is no evident difference between the apparent speeds of on-disk and limb CMEs. According to the analysis of limb CMEs, the average value of angular widths of CMEs is about $59^\\circ$, and about 65\\% of them have a width from $30^\\circ$ to $90^\\circ$. Generally, on-disk CMEs are twice wider than limb CMEs, which suggests a significant projection effect (refer to Sec.\\ref{sec_apparent}). Further, through the analysis based on the source locations of all CMEs, we infer many interesting results. \\begin{enumerate} \\item About 16\\% of LASCO CMEs probably originated from front-side solar disk but left no evident eruptive signatures in the FOV of EIT, and a lower cut-off for the CME visibility in EIT, which corresponds to the apparent speed range of 100 -- 200 km s$^{-1}$, probably exists. (Sec.\\ref{sec_eit}) \\item About 19\\% of CMEs were not detected by LASCO, and the missing rate has a trend to monotonomically increase as the CME source location moves from limb to disk center. (Sec.\\ref{sec_lasco}) \\item About 32\\% of front-side CMEs can not be recognized by SOHO, which becomes a natural explanation of high rate of missing alert of geomagnetic storms and is also in agreement with the previous results that about 18 -- 44\\% of ICMEs do not have the corresponding CMEs. (Sec.\\ref{sec_invisible}) \\item The brightness of a white-light CME at any heliocentric distance is roughly positively-correlated with its speed, which implies that (1) a bright transient recorded in white-light coronagraphs is contributed by both a CME and the compressed solar wind plasma surrounding the CME, and (2) the CME mass derived from the brightness is probably overestimated. (Sec.\\ref{sec_dis2}) \\item Both projection effect and violent eruption are the major causes of halo CMEs, but for limb halo CMEs, the latter should be the primary one. Overall, there are about 25\\% of halo CMEs stronger than the average level of CMEs. (Sec.\\ref{sec_halo}) \\item Most CMEs manifest deflection behaviors near the solar minimum. About 62\\% of CMEs underwent an equator-ward deflection with the average deflection angle of $\\sim22^\\circ$, and about 5\\% of CMEs had a significant pole-ward deflection with the average angle of $\\sim16^\\circ$. At high latitude regions (outside of $\\pm45^\\circ$), most CMEs deflected towards equator. (Sec.\\ref{sec_deflection}) \\item The CME deflections can be classified into three types. One is due to the asymmetrical expansion, one is the non-radial ejection, and the other is the deflected propagation caused by the interaction of the CME with other neighboring magnetic field structures. (Sec.\\ref{sec_dclass}) \\end{enumerate} These findings help people understanding the CMEs viewed in white-light coronagraphs more properly and precisely. We believe that some new and deeper questions have emerged from these results. This paper presents our first work established on the information of CME source locations, and gives us the overview of white-light CMEs. In our follow-up works, we will continue to address issues, e.g., the CME deflections, the role of active regions in producing CMEs, the relationship of CMEs with flares, etc." }, "1101/1101.0707_arXiv.txt": { "abstract": "\\boldmath We aim to present a tutorial on the detection, parameter estimation and statistical analysis of compact sources (far galaxies, galaxy clusters and Galactic dense emission regions) in cosmic microwave background observations. The topic is of great relevance for current and future cosmic microwave background missions because the presence of compact sources in the data introduces very significant biases in the determination of the cosmological parameters that determine the energy contain, origin and evolution of the universe and because compact sources themselves provide us with important information about the large scale structure of the universe. ", "introduction": "CMB anisotropies are extremely weak. Moreover, we cannot observe them directly; we observe instead a mixture of CMB and other astrophysical sources of radiation (usually referred to as \\emph{foregrounds} or just \\emph{contaminants}) that lie along the line of sight of CMB photons. The foregrounds at microwave wavelengths can have Galactic (synchrotron radiation, \\emph{brehmsstrahlung}, thermal and electromagnetic dust emission) or extragalactic (galaxies, galaxy clusters) origin. Besides for ground-based experiments the atmosphere must be accounted as another contaminant. In addition, obviously, all observations are affected by instrumental noise and convolution due to the finite resolution of the optical devices. Therefore a good source separation is an indispensable prerequisite for any serious attempt to do CMB science. Moreover, the microwave window of the electromagnetic spectrum has been opened for observation only very recently. As a result, the microwave sky is poorly known. In particular, the present knowledge about extragalactic sources in frequencies that range from 20 to $\\sim 1000$ GHz is almost completely an extrapolation from observations at lower or higher frequencies, with only a few relevant and very recent direct observations \\cite{SCUBA02_short,hinshaw07_short,NEWPS07}. Present-day CMB telescopes have relatively poor angular resolutions (from a few arcminutes to one degree). Extragalactic foregrounds show angular sizes smaller than, or at most comparable to, such scales\\footnote{In astronomy, the position and size of celestial objects are described in spherical coordinates, with center in the observer. The angular size or angular scale of any object is its visual diameter measured as an angle. Throughout this paper we will use the term \\emph{scale} as a synonym of angular scale.}. Therefore, they can be considered as point-like (compact) objects from the point of view of CMB experiments. The detection of faint compact objects is a very important task that is common to many branches of astronomy. Compact sources constitute a special case among CMB foregrounds. They are the main contaminant at small angular scales~\\cite{zotti99} and they have a dramatic impact on the estimation of the cosmological parametersfrom the CMB signal, on the separation of other components and on the study of the possible non-standard cosmological scenarios (by probing the Gaussianity and isotropy of the CMB). Compact sources come in three main varieties: far galaxies (including radio galaxies, quasars, blazars, 'ordinary' galaxies, dusty galaxies with high stellar formation rate, protogalaxies, spheroids and many other kinds of objects), galaxy clusters (leaving an imprint on CMB radiation, mainly through inverse Compton scattering) and Galactic compact objects (cold cores, supernova remnants, dense knots of dust, etc). The two first classes of objects are almost uniformly distributed across the sky, whereas the third class is mainly concentrated near the Galactic plane. Practically all the objects belonging to the first class have angular sizes that are much smaller than the angular resolution of CMB experiments, while some galaxy clusters and many Galactic objects are observed as extended sources. The electromagnetic spectral behavior of galaxy clusters is pretty well known, while it is not for the other two classes of objects. The previous sentences serve to illustrate the diversity and heterogeneity of compact sources. Due to this heterogeneity, component separation techniques based on generative mixing models usually fail to achieve the separation of compact sources (the only exception to this rule are galaxy clusters). Specific detection (identification) and separation techniques must be tailored for compact sources. ", "conclusions": "" }, "1101/1101.5408_arXiv.txt": { "abstract": "Coincident observations with gravitational wave (GW) detectors and other astronomical instruments are in the focus of the experiments with the network of LIGO, Virgo and GEO detectors. They will become a necessary part of the future GW astronomy as the next generation of advanced detectors comes online. The success of such joint observations directly depends on the source localization capabilities of the GW detectors. In this paper we present studies of the sky localization of transient GW sources with the future advanced detector networks and describe their fundamental properties. By reconstructing sky coordinates of ad hoc signals injected into simulated detector noise we study the accuracy of the source localization and its dependence on the strength of injected signals, waveforms and network configurations. ", "introduction": "There has been a significant sensitivity improvement of the gravitational wave detectors since the Laser Interferometer Gravitational Wave Observatory (LIGO)~\\cite{LIGO} and Virgo observatory~\\cite{Virgo} started their operation. In 2007 LIGO and Virgo completed the two year run at sensitivity that allows detection of a merger of two neutron stars (NS-NS) as far as $\\sim{30}$~Mpc away~\\cite{S5range, cbc2009}. In the most recent run (May 2009 - October 2010) the binary neutron star horizon distance has been increased to $\\sim{40}$~Mpc. However, even at this impressive sensitivity, the anticipated detection rate with the initial LIGO and Virgo detectors is quite low. A detection may be possible in the case of a rare astrophysical transient event such as a supernova explosion in our Galaxy or a nearby merger of binary neutron stars. The signal is likely to be weak and it will be difficult to prove its astrophysical origin unless it is confirmed with a coincident observation of the electromagnetic or neutrino counterpart. For this reason the LIGO and Virgo collaborations are conducting a wide range of joint observations~\\cite{LOOKUP} with other astrophysical experiments including radio~\\cite{LOFAR, ARECIBO}, optical and x-ray telescopes~\\cite{ROTSE, QUEST, TAROT, SWIFT}, and neutrino detectors~\\cite{Antares, IceCube}. A more robust detection of gravitational waves from astrophysical sources is anticipated in the next five years as Advanced LIGO and Advanced Virgo come online. Numerous GW signals, expected to be observed by advanced detectors (likely $\\sim{40}$ NS-NS events per year~\\cite{ExpRates}), will begin our exploration of the gravitational-wave sky and start the era of the gravitational wave astronomy. Along with the advanced GW detectors, a new generation of optical telescopes will come online ~\\cite{LSST, PanStarrs, 30mTelescope}, which will enable a wide and deep survey of the electromagnetic sky. Joint observations with the advanced gravitational wave detectors and electromagnetic instruments will not only increase the confidence of detection but also bring fundamentally new information about the GW sources. They will reveal the physics and dynamic of sources, provide the identification of host galaxies and the associated redshifts, and in some cases determine luminosity distance to the source. One of the major challenges for such joint observations is to establish unambiguous association between a gravitational wave signal and a possible electromagnetic counterpart. It greatly depends on the ability of the GW networks to reconstruct sky coordinates of a detected GW source. Given an accurate sky location, a corresponding electromagnetic transient may be identified in a list of events obtained with the all-sky telescope surveys, or the EM instruments can be guided to take images of a small area in the sky. In the second case, it is important that the sky localization is performed by GW detectors in real time with low latency. The efficiency of the GW-EM association and the choice of a partner telescope is affected by the sky localization error which should be well within the instrument's field of view (typically less than few square degrees). Moreover, exploring smaller area in the sky will decrease the probability of the false association. The problem of the source localization with networks of GW detectors is in the focus of research in the gravitational wave data analysis. There are several analytical studies~\\cite{triang, triang2008, Feir, WenChen} of this problem considering geometrical reconstruction of source coordinates based on the triangulation, which requires a measurement of the arrival time of a GW signal at different detectors. However, the accurate timing of the GW signal is intimately related to the reconstruction of the signal waveforms. Due to the different detector sensitivities to the GW polarizations, the waveforms recorded by individual detectors may be different and they may not have a common timing reference (like a signal peak time) for a direct measurement of the differences in the arrival time. Therefore, the problem of the source localization is better addressed in the framework of the coherent network analysis~\\cite{GT,FH,PRD05}, which reconstructs the waveforms and the sky coordinates simultaneously. By using both these methods (triangulation and coherent network analysis), several practical source localization algorithms~\\cite{klimenko2008, Omega, MBTA} have been recently developed and used during the LIGO and Virgo data taking runs in 2009-2010. There have been a number of studies addressing benefits of individual detectors~\\cite{Schutz1987,Searl2006} and various detector networks~\\cite{Weiss2010,Schutz2010}. In this paper we present a simulation study of the source localization and the reconstruction of GW waveforms with the networks of advanced detectors. The study is performed with a coherent network method, called coherent WaveBurst~\\cite{klimenko2008} (cWB), based on the likelihood analysis . In cWB the data from all detectors in the network is processed simultaneously in order to reconstruct a common GW signal which is consistent with the recorded detector responses. The consistency is measured by the likelihood ratio, which is a function of the source parameters (waveforms and sky location). The most probable source parameters are obtained by maximizing the likelihood ratio over the signal waveforms and sky coordinates. The method performs reconstruction of unmodeled burst signals (arbitrary waveforms) and signals with a certain polarization state: elliptical, linear and circular. The paper is organized as follows. Possible networks of advanced detectors and their fundamental properties are descussed in section~\\ref{sec:networks}. In sections~\\ref{sec:algorithm} we describe the reconstruction algorithm. The simulation framework for this study is presented in section~\\ref{framework}. The results are reported in section~\\ref{Results}. In sections~\\ref{limitation} and~\\ref{conclusions} we describe main factors limiting the source reconstruction and discuss the results. ", "conclusions": "\\label{conclusions} In the paper we present the results of the source localization and reconstruction of GW waveforms with the networks of GW interferometers. For a general characterization of the detector networks we introduce few fundamental network parameters, including the effective noise, and the network antenna and alignment factors. The effective power spectral density of the network noise determines the average network SNR for a given population of GW signals. For each direction in the sky the network performance is characterized by its antenna and alignment factors. The antenna factor describes how uniform is the network response across the sky. The alignment factor, which strongly depends on the number of detectors and the orientation of their arms, determines the relative contribution of the two GW polarizations into the total network SNR. It requires several non-aligned detectors (preferably more than three) for a robust detection and reconstruction of both GW components. The coordinate reconstruction strongly depends on the signal waveforms, network SNR and the number of detector sites in the network. The reconstruction can be significantly improved when it is constrained by the signal model. Although a crude coordinate reconstruction (ring in the sky) is possible with the networks of two spatially separated sites, at least three detector sites are required to perform the source localization. The accuracy of the localization dramatically increases for networks with more than three sites, particularly for the low SNR events. For example, the LH\\~{H}V and LHVA networks are expected to have about the same detection rates, however, the 4-site LHVA network would have much better performance for the accurate reconstruction of GW signals. The pointing resolution required for joint observations with the electromagnetic telescopes is achievable with the networks consisting of four sites. The LHVAJ network demonstrates further improvements, both in the detection and reconstruction of GW signals, reaching a sub-degree angular resolution. In addition, due to the limited duty cycle of the detectors, both the LCGT and the Australian detectors will significantly increase the observation time when any of 4-site networks are operational. The advanced LIGO and Virgo detectors are very capable of the first direct detection of gravitational waves. However, for better reconstruction of the GW signals more detectors are required. Extra detectors introduce an important redundancy which lower the impact of limited duty cycle of the detectors, makes the coordinate reconstruction more accurate, and less dependent on the waveform morphology and calibration uncertainties. The construction of the LCGT and the detector in Australia will significantly enhance the advanced LIGO-Virgo network and these detectors will play a vital role in the future GW astronomy." }, "1101/1101.5122_arXiv.txt": { "abstract": "{The MOND paradigm posits a departure from standard Newtonian dynamics, and from General Relativity, in the limit of small accelerations. The resulting modified dynamics aim to account for the mass discrepancies in the universe without non-baryonic dark matter. I briefly review this paradigm with its basic tenets, and its underlying theories--nonrelativistic and relativistic--including a novel, bimetric MOND gravity theory. I also comment on MOND's possible connection to, and origin in, the cosmological state of the universe at large. Some of its main predictions, achievements, and remaining desiderata are listed. I then succinctly pit MOND against the competing paradigm of standard dynamics with cold, dark matter. Some of the complaints leveled at MOND are: (i) ``MOND was {\\it designed} to fit rotation curves; so no wonder it is so successful in predicting them''. This is both incorrect and quibbling: The first ever MOND rotation curve analysis was undertaken more then four years after the advent of MOND. And, even if MOND, epitomized by a very simple formula, could have been designed to predict hundreds of rotation curves, it would still be a great achievement. (ii) ``MOND outperforms CDM only on small, galactic scales, where formation physics is anyhow very messy, but falls behind in accounting for `simpler', large-scale phenomena''. Quite contrarily, all the salient MOND predictions on galactic scales follow as unavoidable, simple, and immediate corollaries of the theory--{\\it independent of any messy formation scenario}--just as Kepler's laws, obeyed by all planetary systems, follow from an underlying theory, not from complex formation scenarios. To think, as dark-matter advocates say they do, that the universal MOND regularities exhibited by galaxies will one day be shown to somehow follow from complex formation processes, is, to my mind, a delusion. What is left for MOND to explain on large scales is a little in comparison, and has to await a full fledged relativistic MOND theory. (iii) ``The `bullet cluster' shows that MOND still requires some matter that is dark''. Yes, it has long been known that MOND does not fully remove the mass discrepancy in the cores of galaxy clusters. Some additional still-dark matter is needed. But this need not be THE ``dark matter''; a small amount of the still-missing baryons, in some dark form (dead stars? cold gas clouds?), or perhaps (sterile?) neutrinos, could fit the bill.} \\FullConference{Quarks, Strings and the Cosmos - H\\'{e}ctor Rubinstein Memorial Symposium\\\\ August 9-11 ~2010\\\\ AlbaNova, Stockholm, Sweden} \\begin{document} ", "introduction": " ", "conclusions": "" }, "1101/1101.3018_arXiv.txt": { "abstract": "Stars form in dense cores of molecular clouds that are observed to be significantly magnetized. A dynamically important magnetic field presents a significant obstacle to the formation of protostellar disks. Recent studies have shown that magnetic braking is strong enough to suppress the formation of rotationally supported disks in the ideal MHD limit. Whether non-ideal MHD effects can enable disk formation remains unsettled. We carry out a first study on how disk formation in magnetic clouds is modified by the Hall effect, the least explored of the three non-ideal MHD effects in star formation (the other two being ambipolar diffusion and Ohmic dissipation). For illustrative purposes, we consider a simplified problem of a non-self-gravitating, magnetized envelope collapsing onto a central protostar of fixed mass. We find that the Hall effect can spin up the inner part of the collapsing flow to Keplerian speed, producing a rotationally supported disk. The disk is generated through a Hall-induced magnetic torque. Disk formation occurs even when the envelope is initially non-rotating, provided that the Hall coefficient is large enough. When the magnetic field orientation is flipped, the direction of disk rotation is reversed as well. The implication is that the Hall effect can in principle produce both regularly rotating and counter-rotating disks around protostars. We conclude that the Hall effect is an important factor to consider in studying the angular momentum evolution of magnetized star formation in general and disk formation in particular. ", "introduction": "\\label{intro} Disks are an integral part of star formation; they are the birthplace of planets. How they form is a long-standing, unresolved problem. A major difficulty is that their formation is greatly affected by magnetic braking, which has been hard to quantify until recently. There is now increasing theoretical evidence that magnetic braking may suppress the formation of rotationally supported disks (RSDs hereafter) in dense cores magnetized to a realistic level, with dimensionless mass-to-flux ratios $\\lambda$ of a few to several \\citep{tc2008}. \\citet{als2003} first demonstrated through 2D (axisymmetric) simulations that RSDs are suppressed by a moderately strong magnetic field in the ideal MHD limit. \\citet{glsa2006} showed analytically that the disk suppression is due to the formation of a split magnetic monopole, which is an unavoidable consequence of flux freezing. The efficient disk braking was confirmed numerically by \\citet{ml2008} and \\citet{hf2008} (see, however, \\citealt{mim2010} and \\citealt{hc2009} for a different view, and \\citealt{lks2011} for a more detailed discussion). The ideal MHD approximation must break down in order for RSD to form. Two non-ideal MHD effects have already been explored in the context of disk formation: ambipolar diffusion and Ohmic dissipation. \\citet{kk2002} found that ambipolar diffusion tends to make disk braking more efficient, because it enables the magnetic flux that would have gone into the central protostellar object (and form a split monopole) in the ideal MHD limit to pile up at small radii outside of the object instead, making the region strongly magnetized (\\citealt{lm1996}, \\citealt{ck1998}, \\citealt{cck1998}). \\citet{ml2009} showed that the enhanced braking is strong enough to suppress the formation of RSDs for realistic core conditions. The effect of Ohmic dissipation was examined by several groups, starting with \\citet{sglc2006}. They suggested that enhanced resistivity (well above the classical value) is needed for Ohmic dissipation to weaken the magnetic braking enough to form a relatively large RSD of tens of AUs or more. The suggestion was confirmed by \\citeauthor{kls2010} (\\citeyear{kls2010}; KLS10 hereafter), although \\citet{mim2010} were apparently able to form RSDs using the classical resistivity computed by \\citet{nnu2002} (see also \\citealt{db2010}). In any case, a third non-ideal MHD effect, the Hall effect, has not been explored in the context of disk formation (see, however, the independent work of \\citeauthor{b2011} in her unpublished PhD thesis); it is the focus of this paper. Our goal is to determine whether the Hall effect by itself can enable an RSD to form in the presence of a relatively strong magnetic field and, if yes, to estimate the magnitude of Hall coefficient needed for RSD formation. We find that the Hall effect can actively spin up the inner part of the protostellar collapsing flow, potentially to Keplerian speed, unlike the other two non-ideal MHD effects. The combined effect of all these three non-ideal MHD terms on disk formation will be explored in another investigation (\\citealt{lks2011}; LKS11 hereafter). ", "conclusions": "\\label{discussion} We have studied the collapse of a non-self-gravitating rotating, magnetized envelope onto a central stellar object of fixed mass to illustrate the influence of the Hall effect on disk formation. In this idealized problem, the formation of a rotationally supported disk (RSD) is completely suppressed by magnetic braking in the ideal MHD limit. Including a Hall term with a coefficient $Q \\gtrsim 3\\times 10^{20}\\Qunit$ in the induction equation enables an RSD of radius $\\gtrsim 10\\AU$ to form. The RSD is formed not because the Hall effect has reduced the efficiency of the magnetic braking. Rather, it is produced because the Hall effect actively spins up the inner part of the equatorial material to Keplerian speed. The spin-up comes about because the collapsing envelope drags the magnetic field into a highly pinched configuration near the equator, producing a large toroidal electric current density, which forces the field lines to rotate differentially due to the Hall effect. The resulting twist of field lines yields the magnetic torque that spins up the equatorial material, even in the absence of any initial rotation. The spin-up is most effective in the inner part of the accretion flow where radial Hall diffusion enables the magnetic flux that would have been dragged into the central object by the accreted material in the ideal MHD limit to stay behind; the resulting pileup of magnetic flux at small radii is similar to the cases with either ambipolar diffusion or Ohmic dissipation. When the field direction is flipped, the Hall-induced magnetic torque changes direction as well. An implication is that the direction of the angular momentum of the material close to a protostar, including RSD, may depend more on the magnetic field orientation than on the initial rotation of the dense core that the star is formed out of, at least in principle. In practice, it is uncertain whether the Hall-induced magnetic torque can produce a sizable RSD or not. First, the value of the coefficient $Q$ required to produce an RSD of tens of AUs in size in our idealized model appears to be larger than the microscopic values expected in dense cores, by roughly an order of magnitude. Whether the coefficient can be enhanced somehow, perhaps through anomalous processes, is unclear. More importantly, the Hall effect is only one of the non-ideal MHD effects that are present in the lightly ionized, magnetized, dense cores. It can be dominated by ambipolar diffusion at low densities and by Ohmic dissipation at high densities. It will be interesting to explore the interplay of all these three non-ideal MHD effects and how they affect the collapse of dense cores and the formation of protostellar disks. Nevertheless, we have demonstrated that the Hall effect is unique among the non-ideal effects in its ability to actively spin up a magnetized collapsing flow, potentially to Keplerian speed, providing a new mechanism for disk formation." }, "1101/1101.4915_arXiv.txt": { "abstract": "We investigate the wavenumber scale of Fe~I and Fe~II lines using new spectra recorded with Fourier transform spectroscopy and using a re-analysis of archival spectra. We find that standards in Ar~II, Mg~I, Mg~II and Ge~I give a consistent wavenumber calibration. We use the recalibrated spectra to derive accurate wavelengths for the a$^6$D-y$^6$P multiplet of Fe~II (UV 8) using both directly measured lines and Ritz wavelengths. Lines from this multiplet are important for astronomical tests of the invariance of the fine structure constant on a cosmological time scale. We recommend a wavelength of 1608.45081~\\AA\\ with a one standard deviation uncertainty of 0.00007~\\AA\\ for the $\\rm a^6D_{9/2}-y^6P_{7/2}$ transition. ", "introduction": "The universality and constancy of the laws of nature rely on the invariance of the fundamental constants. However, some recent measurements of quasar (Quasi-stellar objects - QSO) absorption line spectra suggest that the fine-structure constant, $\\alpha$, \\cite{fund_const} may have had a different value during the early universe \\cite{murphy_03}. Other measurements (e.g. \\cite{chand_06}) do not show any change. The attempt to resolve these discrepancies can probe deviations from the standard model of particle physics and thus provide tests of modern theories of fundamental interactions that are hard to attain in other ways. QSO absorption lines are used in these investigations to measure the wavelength separations of atomic lines in spectra of different elements - the many-multiplet method \\cite{Dzuba_99} - and compare their values at large redshifts with their values today. Any difference in the separations would suggest a change in $\\alpha$. Since this method uses many different species in the analysis that have differing sensitivities to changes in $\\alpha$, it can be much more sensitive than previous methods, such as the alkali-doublet method \\cite{Bahcall_67}, that use just one species. However, it requires very accurate laboratory wavelengths to be used successfully, since the observed changes in $\\alpha$ are only a few parts in 10$^5$, requiring laboratory wavelengths to 1:10$^7$ or better. This has led to several recent measurements of ultraviolet wavelengths using both Fourier transform (FT) spectroscopy \\cite{Pickering_98,Aldenius_06,Aldenius_09} and frequency comb metrology \\cite{Salumbides_06,Hannemann_06,Batteiger_09}. One spectral line of particular interest is the $\\rm 3d^6(^5D)4s\\,a^6D_{9/2} - 3d^5\\,(^6S)4s4p(^3P)\\,y^6P_{7/2} $ line of Fe~II at 1608.45~\\AA. This line is prominent in many QSO spectra and its variation with $\\alpha$ has the opposite sign from that of other nearby lines \\cite{Murphy_03}. However, measurement of its wavelength using frequency comb metrology, which is at present the most accurate method, is extremely difficult due to its short wavelength. Although this line is strong in many of the FT spectra of iron-neon hollow cathode lamps recorded at the National Institute of Standards and Technology (NIST) and Imperial College, London, UK (IC), these spectra display inconsistencies in the wavelength of the 1608~\\AA\\ line of around 1.5 parts in 10$^7$ - too great for use in the many-multiplet method to detect changes in $\\alpha$. We discussed some of these discrepancies in our previous paper \\cite{Nave_04} presenting reference wavelengths in the spectra of iron, germanium and platinum around 1935~\\AA . Here we present a re-analysis of spectra taken at NIST and IC in order to resolve these discrepancies and provide a better value for the wavelength of the 1608.45~\\AA\\ line of Fe~II. The papers involved in this re-analysis are listed in table \\ref{corrections}, together with the proposed corrections to the wavenumber scale. The proposed corrections are up to three times the previous total uncertainty, depending on the wavenumber. In section \\ref{previous} we discuss previous measurements of the a$^6$D - y$^6$P multiplet. Section \\ref{expt} describes the archival data we use to obtain improved wavelengths for this multiplet Additional spectra taken at NIST in order to re-evaluate the calibration of these archival data are described in section \\ref{cal}. The accuracy of this calibration in the visible and ultraviolet wavelength regions is also discussed in section \\ref{cal}. Section \\ref{a6D-y6P} describes three different methods for obtaining the wavelengths of the a$^6$D - y$^6$P multiplet. The first method uses intermediate levels determined using strong Fe~II lines in the visible and ultraviolet in order to obtain the values of the y$^6$P levels and Ritz wavenumbers for the a$^6$D - y$^6$P multiplet. The second method uses energy levels optimized by using a large number of spectral lines to derive Ritz wavenumbers for this multiplet. Although better accuracy is achieved using this method than the first method because of the increased redundancy, the way in which the y$^6$P levels are determined is less transparent. The third method uses experimental wavelengths determined in spectra that are recalibrated from spectra in which we have re-evaluated the wavenumber calibration. In section \\ref{ge_comp} we re-examine the Fe~II wavenumbers in our previous paper \\cite{Nave_04}. All uncertainties in this paper are reported at the one standard deviation level. ", "conclusions": "We investigated the wavenumber scale of published Fe~I and Fe~II lines using new spectra recorded with the NIST 2-m FT spectrometer and a re-analysis of archival spectra. Our new spectra confirm the wavenumber scale of visible-region iron lines calibrated using the Ar~II wavenumber standards of Whaling et al. \\cite{Whaling_95}. Having confirmed the wavenumber scale of iron lines in the visible and ultraviolet regions, we have used lines from these spectra to derive Ritz values for the wavenumbers and wavelengths of lines in the $\\rm a^6D - y^6P$ multiplet of Fe~II (UV 8). Ritz wavenumbers derived using two different methods agree with one another and with directly measured wavenumbers within the joint uncertainties. We recommend a value of 1608.45081$\\pm$0.00007~\\AA\\ for the wavelength of the $\\rm a^6D_{9/2} - y^6P_{7/2}$ line of Fe~II, which is an important line for detection of changes in the fine-structure constant during the history of the Universe using quasar absorption-line spectra. We find that the wavenumbers in Learner \\& Thorne \\cite{Learner_88} and Table 3 of Nave et al. \\cite{Nave_91} should be increased by 6.7 parts in 10$^8$ to put them on the scale of the Ar~II lines of Whaling et al. \\cite{Whaling_95}. The wavenumbers in Tables 4 and 5 of Ref. \\cite{Nave_91} should be increased by 10.6 parts in 10$^8$ to put them on the Ar~II scale of Ref. \\cite{Whaling_95} and to correct for an error in the transfer of this wavenumber scale to the ultraviolet. The Ge~I wavenumbers of Kaufman \\& Andrew \\cite{Kaufman_62} and all the wavenumbers in Nave \\& Sansonetti \\cite{Nave_04} should be decreased by 1.4 parts in 10$^8$ to put them on the scale of recent measurements of the $^{198}$Hg line at 5462~\\AA." }, "1101/1101.0850_arXiv.txt": { "abstract": "Pulsars are unique astrophysical laboratories because of their clock-like timing precision, providing new ways to test general relativity and detect gravitational waves. One impediment to high-precision pulsar timing experiments is timing noise. Recently \\citet{lyn10} showed that the timing noise in a number of pulsars is due to quasi-periodic fluctuations in the pulsars' spin-down rates and that some of the pulsars have associated changes in pulse profile shapes. Here we show that a non-radial oscillation model based on asteroseismological theory can explain these quasi-periodic fluctuations. Application of this model to neutron stars will increase our knowledge of neutron star emission and neutron star interiors and may improve pulsar timing precision. ", "introduction": "Pulsars have long been known as superb laboratories for fundamental physics due to their timing precision and stable pulse profiles. Timing observations of pulsars have resulted in the best tests of general relativity in the strong-field regime, the most precise stellar mass measurements, and the most constraining limits on gravitational wave sources in the nanoHertz regime through `pulsar timing array' experiments \\citep[e.g.][]{jen06,kra06a}. One significant limitation for all of these studies is timing noise, or non-random variations in the residuals between measured and model-predicted arrival times. \\citet{lyn10} recently showed that the timing noise in 17 pulsars is due to the pulsar's spin-down rate fluctuating between different values in a quasi-periodic manner. Furthermore, for six of these pulsars, the changes in the spin-down rate directly correlate to changes in the pulse shape. While the pulsars described in \\citet{lyn10} have long periods, some millisecond pulsars, such as PSR B1931+27, also exhibit timing noise and may show similar but less obvious effects. This behavior offers a possible prescription for removing timing noise from data and dramatically increasing timing stability, facilitating projects such as the direct detection of gravitational waves through pulsar timing. In this paper, we present a model for this behavior based on non-radial oscillations. We outline the model in \\S\\ref{model}, discuss underlying assumptions in \\S\\ref{theory}, describe the implications for the pulsars discussed in \\citet{lyn10} in \\S\\ref{predictions}, and present conclusions in \\S\\ref{conc}. ", "conclusions": "\\label{conc} In summary, we have shown that the fluctuations in spin-down rate described in \\citet{lyn10} can be explained by a non-radial oscillation model. Our model is not complete, as little is understood about pulsar magnetospheres and there is no clear mechanism for driving the oscillations. However, our oscillation model does not require invoking an external influence to explain the periodicity, as in \\citet{cor08}, and is based on observations of other well studied stars. This will lead to better understanding of pulsar emission physics and may enable optimal removal of effects due to non-radial oscillations in pulsar timing data, increasing the usefulness of pulsars as fundamental physics laboratories." }, "1101/1101.0582_arXiv.txt": { "abstract": "The observation of high energy cosmic neutrinos can shed light on the astrophysical sites and mechanisms involved in the acceleration of protons and nuclei to the high energies observed at Earth by cosmic ray detectors. More generally, high energy neutrinos can be a key instrument in the multimessenger study of the high energy sky. Several neutrino telescopes of different sizes and capabilities are presently taking data and projects to further increase their size and sensitivity are underway. In this contribution we review the present status of neutrino telescopes based on the cherenkov ligh detection technique, their recent results and the plans to increase their sensitivity. \\vspace{1pc} ", "introduction": "The detection of high energy cosmic neutrinos is a long and much sought-after scientific goal. In the low energy domain (few MeV to several GeV) the observation of extraterrestrial and atmospheric neutrinos gave rise to the discovery of neutrino oscillations and to the -arguably- first experimental test of our models of supernova explosion. In the high energy regime (several GeV to EeV), the advantages of neutrinos as cosmic messengers are well known, their importance to provide information on the particle acceleration mechanisms in astrophysical objects is also acknowledged and experimental methods to detect them exist and have been technologically proven. Therefore, at present there is a struggle to reach the required sensitivities to detect them, the hope being that nature is generous enough to provide fluxes at a level observable with our present or soon-to-come neutrino telescopes. Let us briefly summarize the advantages of neutrinos as cosmic messengers. They are neutral particles, they are not deflected by magnetic fields and therefore they point back to their sources. They are weakly interacting and thus they can escape from very dense astrophysical objects and can travel long distances without being absorbed by matter or background radiation. Moreover, in cosmic sites where hadrons are accelerated it is likely that neutrinos are generated in the decay of charged pions produced in the interaction of those hadrons with the surrounding matter or radiation, being therefore a smoking gun of hadronic acceleration mechanisms. The observation of neutrinos in a cherenkov neutrino telescope is based on the detection of the muons produced by the neutrino charged current interactions with the matter surrounding the telescope by means of the cherenkov light induced by the muons when crossing the detector medium, natural ice or water. The telescope consists in a three dimensional array of light sensors, usually photomultipliers, that record the position and time of the emitted cherenkov photons and therefore enable the reconstruction of the muon track. To avoid the huge background of muons produced in cosmic rays showers, the telescopes look at the other side of the Earth, using it as a shield. The increase in the range of muons at high energies (form kilometres to several kilometres) together with the increase with energy of the neutrino cross section give rise to an approximately exponential increase of the effective areas of these devices in the GeV to PeV energy range. Above a few TeV, the telescopes can determine the direction of the incoming neutrinos with angular resolutions better than 1$^\\circ$, hence the name ``telescope''. At energies above the PeV, the Earth becomes opaque to neutrinos, but the atmospheric muon flux decreases dramatically so that the neutrino telescopes can look for downgoing neutrinos in that energy regime. Other neutrino flavours can be observed through the detection of hadronic or electromagnetic showers or, in the case of tau neutrinos, via the observation of its interaction and the subsequent decay of the produced tau lepton. ", "conclusions": "" }, "1101/1101.2855_arXiv.txt": { "abstract": "We present \\textit{Spitzer} infrared spectra and ultra-violet to mid-infrared spectral energy distributions (SEDs) of 25 luminous type 1 quasars at z $\\sim$ 2. In general, the spectra show a bump peaking around 3 $\\mum$, and the 10 $\\mum$ silicate emission feature. The 3 $\\mum$ emission is identified with hot dust emission at its sublimation temperature. We explore two approaches to modeling the SED: (i) using the \\clumpy model SED from Nenkova \\etal (2008a), and (ii) the \\clumpy model SED, and an additional blackbody component to represent the 3 $\\mum$ emission. In the first case, a parameter search of $\\sim 1.25$ million \\clumpy models shows: (i) if we ignore the UV-to-near-IR SED, models fit the 2--8 $\\mum$ region well, but not the 10 $\\mum$ feature; (ii) if we include the UV-to-near-IR SED in the fit, models do not fit the 2--8 $\\mum$ region. The observed 10 $\\mum$ features are broader and shallower than those in the best-fit models in the first approach. In the second case, the shape of the 10 $\\mum$ feature is better reproduced by the \\clumpy models. The additional blackbody contribution in the 2--8 $\\mum$ range allows \\clumpy models dominated by cooler temperatures ($T < 800 {\\rm K}$) to better fit the 8--12$\\mum$ SED. A centrally concentrated distribution of a small number of torus clouds is required in the first case, while in the second case the clouds are more spread out radially. The temperature of the blackbody component is $\\sim 1200 \\kel$ as expected for graphite grains. ", "introduction": "In the unified model of active galactic nuclei (AGN) \\citep{1993ARA&A..31..473A,1995PASP..107..803U}, the dust torus is a region immediately outside the accretion disk where dusty clouds are no longer sublimated by the radiation from the central engine. The dust torus reprocesses the incident ultra-violet/optical radiation from the accretion disk and this energy emerges in the near- and mid-infrared bands. \\citet{2006ApJS..166..470R} presented panchromatic spectral energy distributions (SEDs) for 259 type 1 quasars selected from the Sloan Digital Sky Survey\\footnote{http://www.sdss.org/} \\citep[SDSS,][]{2000AJ....120.1579Y}. These quasar SEDs constructed from broad-band photometry are remarkably similar over a large range in both luminosity and redshift. However, \\citet{2006ApJS..166..470R} noted small differences in the 1.3--8 $\\mum$ range between optically luminous and optically dim quasars. \\citet{2007ApJ...661...30G} investigated this further, and found that the 1--8 $\\mum$ spectral index ($\\alpha_{\\nu}$) is strongly anti-correlated with infrared luminosity in type 1 quasars. More luminous quasars have bluer 1--8 $\\mum$ slopes. Further, they noted a tight linear correlation between the ultra-violet (UV) continuum luminosity and the infrared luminosity for these quasars. This suggested that the observed near-IR emission at $3\\ \\mum$ in the SED of many type 1 objects is driven by the dust reprocessing of the intrinsic optical/ultra-violet continuum from the accretion disk, as had been noted previously \\citep{1969Natur.223..788R,1979ApJ...230...79N,1986ApJ...308...59E,1987ApJ...320..537B,1989ApJ...347...29S}. As a recent example, the near-IR emission is clearly visible in the spectrum of Mrk~1239 \\citep{2006MNRAS.367L..57R}. Theoretical work on the response of accretion disks to radiation and hydromagnetic pressure suggests that outflow of matter is associated with all accretion disks in the form of a wind coming off the surface of the disk \\citep{1994ApJ...434..446K,1995ApJ...454L.105M,2000ApJ...543..686P}. The dusty torus itself may be the outermost part of this accretion disk wind close to the equator of the system \\citep{1994ApJ...434..446K,2006ApJ...648L.101E}. Disk-winds have a natural dependence on luminosity through radiation pressure, and this begs the question: ``is the structure of the dusty torus related to the physics of the accretion disk?''. The need for proper radiation transfer treatment of clumps in dusty tori was recognized in the pioneering early studies \\citep{1988ApJ...329..702K,1992ApJ...401...99P,1995MNRAS.272..737R}, and was fully developed by \\citet{2002ApJ...570L...9N}. More recently, \\citet{2008ApJ...685..147N} presented their model in detail (denoted by \\textsc{Clumpy} hereafter). Significant effort has been invested in understanding the torus dust distributions with various groups favoring both clumpy and smooth dust density distributions \\citep{2002ApJ...570L...9N,2005A&A...436...47D,2005A&A...437..861S,2006MNRAS.366..767F,2006A&A...452..459H,2008A&A...482...67S,2008ApJ...685..147N}. The primary difference between clumpy and smooth models is that of the dust temperature distributions \\citep[see Fig.~3 of][]{2008A&A...482...67S}. While in smooth density models, the temperature steadily declines with radius from the inner wall, clumpy models can show a range of temperatures at different distances from the central source. This effect occurs primarily due to the shadowing effect from the finite size of clouds. The inner faces of clouds are directly exposed to radiation from the central source, and are hence hotter, while their outer faces are much cooler. And because of clumpiness, clouds farther out in radius can still have their inner faces exposed directly to radiation from the central source. The effective optical depth in a clumpy torus is a function of the number density of clouds in the central regions of the torus. This important model construction has resulted in better fits to both low-resolution \\textit{Spitzer} spectra \\citep{2009ApJ...705..298M,2009ApJ...707.1550N}, and high-resolution interferometric observations of dust tori in NGC 1068 \\citep{2004Natur.429...47J} and Circinus \\citep{2007A&A...474..837T}. \\clumpy models appear to be the most promising set of models with a wide range of applications to both active galactic nuclei (AGN) \\citep[e.g.,][]{2006ApJ...640..612M} and merger-driven ultra-luminous infrared galaxies (ULIRG) \\citep{2007ApJ...654L..45L}. Other notable models that employ clumps arranged in a disk-like geometry include \\citet{2008A&A...482...67S} and \\citet{2006A&A...452..459H}. For example, \\citet{2008ApJ...675..960P} employed clumpy torus models from \\citet{2006A&A...452..459H} to fit their optically obscured but infrared-bright sources at high-redshift. \\clumpy models show changes in their near-IR continua based on the average number of clouds ($N_{0}$) encountered along a radial equatorial ray \\citep[see Fig.~6 in][]{2008ApJ...685..160N}. Using \\textit{Spitzer} mid-IR spectroscopy of high redshift quasars it is then possible to constrain the parameters of their dusty tori. While \\textit{Spitzer} archives are rich in observations of low-redshift Seyfert galaxies, they are deficient in high-redshift observations of radio-quiet quasars at the peak of the quasar activity in the Universe. In this paper, we present such observations as obtained with the Infrared Spectrograph (IRS) on board \\textit{Spitzer}. Our goals include: (1) presenting high-quality mid-IR quasar spectra covering rest-frame 2--12$\\mum$ for comparison to the low-redshift templates already available \\citep[e.g.,][]{2005ApJ...625L..75H,2005ApJ...633..706W,2006AJ....132..401B,2006ApJ...640..579G,2006ApJ...653..127S,2006ApJ...649...79S}; (2) testing the validity of \\clumpy torus models by fitting the observed spectra with model SEDs. Using good-quality IRS spectra we hope to model the 10 $\\mum$ region properly and constrain \\clumpy torus parameters for these luminous quasars. The properties of the sample and reduction process of the IRS spectra are presented in Section~\\ref{sec:data}. The IRS spectra and SEDs of the sample are discussed in Section~\\ref{sec:the-spectra}. \\clumpy torus models are summarized in Section~\\ref{sec:clumpy-torus-models}, and Section~\\ref{sec:model-fits} presents results of model fits to ultra-violet to mid-IR SEDs. Results are summarized in Section~\\ref{sec:summary}. In all calculations, we assume a standard cosmology with $H_{0} = 71 \\kmsmpc$, $\\Omega_{\\rm{M}} = 0.27$, and $\\Omega_{\\rm{Vac}} = 0.730$. ", "conclusions": "\\label{sec:summary} We present \\textit{Spitzer}/Infrared Spectrograph (IRS) observations of a sample of optically luminous type 1 quasars at z$\\sim$2. Their rest-frame 2--12 $\\mum$ infrared spectra show two prominent features peaking at $\\sim$ 3 and 10 $\\mum$. The 10 $\\mum$ feature is the 10 $\\mum$ silicate emission feature, commonly observed in \\textit{Spitzer} observations of other type 1 AGN \\citep{2005ApJ...625L..75H,2005AN....326R.556S,2005ApJ...629L..21S}. The 3 $\\mum$ bump is the expected signature of the hottest thermal dust emission from the inner region of the dust torus. There is a strong correlation between the optical/UV and infrared luminosities \\citep{2007ApJ...661...30G}, and the detection of this near-IR bump in a sample of optically luminous high redshift quasars, shows that the optical/UV continuum heats the dust in the inner torus, which then radiates in the thermal near- to mid-infrared. We fit the spectra and the UV-to-MIR SED with \\clumpy torus models \\citep{2008ApJ...685..147N}. This is the first time such fits have been attempted to spectroscopically confirmed high-z quasars with near-IR data. We considered two different approaches. In the first case, we use the \\clumpy model SED. These \\clumpy torus models provide good fits to the 2--8 $\\mum$ part of the spectrum, if we only fit data longward of 1 $\\mum$. Models with average number clouds along a radial equatorial ray ($N_{0}$) $\\sim 1$, optical depth through each cloud ($\\tau_{V}$) $\\lesssim 10$, and a radial distribution of clouds ($r^{-q}$) described by a power-law exponent ($q$) $\\sim 3$ fit IRS spectra (not complete SEDs) with a strong hot-dust bump very well. The $q \\sim 3$ values suggests that the hot dust component is more centrally concentrated as expected. However, the 10 $\\mum$ silicate emission features of these models show strongly peaked profiles, and the 10 $\\mum$ feature in the observed spectra are more broad and flat. This problem can be partially removed by fitting the entire SED from UV-to-MIR; using this long lever-arm, the \\clumpy model SED is consistently weaker than the observed SED in the 1--7 $\\mum$ range (see left panels of Figure~\\ref{fig:clumpy_fits}), highlighting the lack of additional near-IR contribution in the models, if both UV and IR data is fitted together. To accurately model the 10 $\\mum$ silicate emission features, and remove the above inconsistency, we considered the \\clumpybb model where we fit the spectra and the SED with a linear combination of a hot dust blackbody and a \\clumpy model. In these fits, the clumpy models provide good fits to the 10 $\\mum$ region, while the blackbody contributes more strongly to the region between 2--8 $\\mum$. Use of the additional blackbody leads to a stronger contribution of the \\clumpy model to the far-IR emission. Whether this is a real effect may be tested using far-IR facilities like \\textit{Herschel}. We compared the infrared properties of this sample to the low-redshift PG quasar sample ($z\\sim0.1$) from the \\textit{Spitzer} archives, and find that the primary difference in the 2--8 $\\mum$ range between low- and high redshift samples is the absolute luminosity. There are however significant object-to-object differences in the 10 $\\mum$ silicate emission features, which point to real differences in the dust structure of their tori. In few cases, such as SDSSJ142945, the 9.7$\\mum$ peak of the silicate feature appears shifted to longer wavelengths. Just as other observations have noted the presence of different dust species \\citep{2005ApJ...625L..75H,2005ApJ...629L..21S,2007ApJ...668L.107M}, we note a feature around ~11.3 $\\micron$ in some sources that may be due to crystalline silicates \\citep{2007ApJ...668L.107M}. The 10 $\\mum$ feature shapes in 14 out of 25 objects are well-reproduced by \\clumpy models, the agreement is weak in other cases mostly due to lack of a clear emission feature. Presence of additional dust species also seems to contribute to this issue. More work is necessary to connect the near-IR emission with the rest of the torus structure. The lack of near-IR contribution in the torus models with clumpy media (in general) appears to be rooted in not considering the balance of amounts of silicate and graphite grains as a function of distance from the source. However, we find that the near-to-mid IR SED analysis is a powerful tool to distinguish between different distributions of $q$, $N_{0}$ and $\\tau_{V}$ in \\clumpy models. Observing a blue 3--8 $\\mum$ continuum indicates that the source is compact ($q > 1$) with $N_{0} \\sim 1$. A redder continuum may require a more extended ($q < 1$) distribution of clumps with $N_{0} \\sim 10$ and $\\tau_{V} \\sim 30$. Further, improvements in fits using the complete UV-to-MIR SED suggests the importance of using UV/optical data if available. Further FIR data where the contribution from cold dust associated with star formation in the host galaxy of the quasar may be dominant \\citep{2007ApJ...666..806N}, is also important. The radial extent of the torus ($Y$) is constrained by the location of the FIR turn-over in the infrared SED; however contribution from cold dust in the host galaxy is also dominant in the same region, disentangling these contributions will be interesting \\citep[see for example][]{2010A&A...518L..33H}. In a \\clumpy torus, the probability of viewing the AGN as a type 1 object depends more strongly on $N_{0}$ and $\\tau_{V}$, than on the inclination to the line of sight $i$. Using multi-component models decreases this sensitivity of the model SED to parameters like $N_{0}$. This is observed in the number of accepted models in Table~\\ref{tab:torus_par_stats}; even for objects with S/N $\\sim 25$ (SDSS100401, SDSS151307), the number of accepted models is $\\lesssim 1000$. The argument in favor of \\clumpybb models is that they represent the complete data range better, and adding a blackbody component improves the fits to the 10 $\\mum$ region (right panels in Figure~\\ref{fig:clumpy_fits}), even in case of objects like SDSS142730 that should be dominated by the \\clumpy model alone. Addition of the blackbody component to represent the near-IR emission does not by itself represent a failure of \\clumpy models, but suggests that more detailed treatment of the origin of the near-IR emission is required. The composite grain approximation assumed in radiative transfer calculations \\citep[DUSTY][]{1999astro.ph.10475I} may lead to stronger 10 $\\mum$ features than would be generated in the actual dust sublimation transition region. This effect is also seen in models of \\citet{2005A&A...437..861S} that use the standard MRN dust grain mixture, and obtain strong 10 $\\mum$ emission features in their SEDs. As the models fits in this paper show, \\clumpy models can reproduce the 10 $\\mum$ shapes adequately. Differences in number density of dust grains of different sizes and compositions with distance from the continuum source likely contribute to the nature of near-IR emission. This dust sublimation region may also be spread out over an extended region rather than in a thin AGN-facing layer of the cloud as assumed in \\clumpy models. Future clumpy torus models should consider both these effects to properly model the near- to mid-IR SEDs of active galaxies." }, "1101/1101.2122_arXiv.txt": { "abstract": "{} {We investigate the molecular gas properties of a sample of 23 galaxies in order to find and test chemical signatures of galaxy evolution and to compare them to IR evolutionary tracers.} { Observation at 3 mm wavelengths were obtained with the EMIR broadband receiver, mounted on the IRAM 30~m telescope on Pico Veleta, Spain. We compare the emission of the main molecular species with existing models of chemical evolution by means of line intensity ratios diagrams and principal component analysis. } {We detect molecular emission in 19 galaxies in two 8 GHz-wide bands centred at 88 and 112 GHz. The main detected transitions are the $J$=1--0 lines of CO, $^{13}$CO, HCN, HNC, HCO$^+$, CN, and C$_2$H. We also detect HC$_3$N $J$=10--9 in the galaxies IRAS~17208, IC~860, NGC~4418, NGC~7771, and NGC~1068. The only HC$_3$N detections are in objects with HCO$^+$/HCN$<$1 and warm IRAS colours. Galaxies with the highest HC$_3$N/HCN ratios have warm IRAS colours (60/100 $\\mu$m$>$0.8). The brightest HC$_3$N emission is found in IC~860, where we also detect the molecule in its vibrationally excited state. We find low HNC/HCN line ratios ($<$0.5), that cannot be explained by existing PDR or XDR chemical models. The intensities of HCO+ and HNC appear anti-correlated, because galaxies with low HCO+/HCN intensity ratios have high HNC/HCN. No correlation is found between the HNC/HCN line ratio and dust temperature. All HNC-bright objects are either luminous IR galaxies (LIRG) or Seyferts. Galaxies with bright polycyclic aromatic hydrocarbons (PAH) emission show low HNC/HCO$^+$ ratios. The CO/$^{13}$CO ratio is positively correlated with the dust temperature and is generally higher than in our galaxy. The emission of CN and C$^{18}$O is correlated.} {Bright HC$_3$N emission in HCO$^+$-faint objects may imply that these are not dominated by X-ray chemistry. Thus the HCN/HCO$^+$ line ratio is not, by itself, a reliable tracer of XDRs. Bright HC$_3$N and faint HCO$^+$ could be signatures of embedded star-formation, instead of AGN activity. Mechanical heating caused by supernova explosions may be responsible for the low HNC/HCN and high HCO$^+$/HCN ratios in some starbursts. We cannot exclude, however, that the discussed trends are largely caused by optical depth effects or excitation. Chemical models alone cannot explain all properties of the observed molecular emission. Better constraints to the gas spacial distribution and excitation are needed to distinguish abundance and excitation effects.} ", "introduction": " ", "conclusions": "" }, "1101/1101.4018_arXiv.txt": { "abstract": "We present the results of spectroscopic observations in the GOODS-N field completed using DEIMOS on the Keck II telescope as part of the DEEP3 Galaxy Redshift Survey. Observations of $370$ unique targets down to a limiting magnitude of $R_{\\rm AB} = 24.4$ yielded $156$ secure redshifts. In addition to redshift information, we provide sky-subtracted one- and two-dimensional spectra of each target. Observations were conducted following the procedures of the Team Keck Redshift Survey (TKRS), thereby producing spectra that augment the TKRS sample while maintaining the uniformity of its spectral database. ", "introduction": "\\label{sec_intro} Due in large part to the Great Observatories Origins Deep Survey \\citep[GOODS,][]{giavalisco04}, the GOODS-N field ($\\alpha = 12^{\\rm h}36^{\\rm m}55^{\\rm s}$, $\\delta = +62^{\\circ}14^{\\rm m}15^{\\rm s}$) has become one of the most well-studied extragalactic fields in the sky with existing observations among the deepest at a broad range of wavelengths (e.g., \\citealt{alexander03}; \\citealt{morrison10}; Elbaz et al.\\ in prep). In the coming years, this status as one of the very deepest multiwavelength survey fields will be further cemented by the ongoing and upcoming extremely-deep observations with {\\it Spitzer}/IRAC and {\\it HST}/WFC3-IR as part of the Spitzer Extended Deep Survey (SEDS, PI G.\\ Fazio) and the Cosmic Assembly Near-IR Deep Extragalactic Legacy Survey (CANDELS, PIs S.\\ Faber \\& H.\\ Ferguson), respectively. Given the large commitment of telescope time from both space- and ground-based facilities devoted to imaging the GOODS-N field, spectroscopic observations in this field possess a significant legacy value. For instance, spectroscopic redshifts dramatically improve the constraints inferred from imaging alone, allowing rest-frame quantities to be derived with increased precision. Furthermore, only through spectroscopy can assorted spectral and dynamical properties (such as the strengths and velocity widths of emission and absorption lines) be measured. Recognizing the potential legacy value of spectroscopic observations in the GOODS-N field, the Team Keck Redshift Survey \\citep[TKRS,][]{wirth04} utilized the DEep Imaging Multi-Object Spectrograph \\citep[DEIMOS,][]{faber03} on the Keck II telescope to create a publicly-available redshift catalog and uniform spectral database across the entire area imaged with {\\it HST}/ACS by the GOODS Team. Altogether, the TKRS observed nearly $3000$ sources, yielding secure spectroscopic redshifts for $\\sim 1500$ objects and enabling numerous studies of galaxy evolution and cosmology \\citep[e.g.,][]{kk04, weiner06, riess07, juneau10}. In an effort to augment the value of the existing TKRS dataset, we present observations of $370$ unique sources in the GOODS-N field (a $> \\! 10\\%$ increase to the TKRS sample size), collected as part of the DEEP3 Galaxy Redshift Survey (Cooper et al.\\ 2011, in prep) and using the same instrument and observation methods as the TKRS. The DEEP3 survey is an ongoing spectroscopic effort designed to leverage the vast amounts of multiwavelength data in another prime deep extragalactic field, the Extended Groth Strip (EGS). Once completed, DEEP3 will yield Keck/DEIMOS spectra of $\\gtrsim 7500$ sources at $z < 2$, which when combined with TKRS and this work will create an extensive spectral database that is both uniform and publicly-available. In Sections \\ref{sec_design} and \\ref{sec_redux}, we describe the design, execution, and reduction of our Keck/DEIMOS observations in GOODS-N, with the data products presented in Section \\ref{sec_data}. ", "conclusions": "" }, "1101/1101.3312_arXiv.txt": { "abstract": "We present the first example of binary microlensing for which the parameter measurements can be verified (or contradicted) by future Doppler observations. This test is made possible by a confluence of two relatively unusual circumstances. First, the binary lens is bright enough ($I=15.6$) to permit Doppler measurements. Second, we measure not only the usual 7 binary-lens parameters, but also the ``microlens parallax'' (which yields the binary mass) and two components of the instantaneous orbital velocity. Thus we measure, effectively, 6 'Kepler+1' parameters (two instantaneous positions, two instantaneous velocities, the binary total mass, and the mass ratio). Since Doppler observations of the brighter binary component determine 5 Kepler parameters (period, velocity amplitude, eccentricity, phase, and position of periapsis), while the same spectroscopy yields the mass of the primary, the combined Doppler + microlensing observations would be overconstrained by $6 + (5 + 1) - (7 + 1) = 4$ degrees of freedom. This makes possible an extremely strong test of the microlensing solution. We also introduce a uniform microlensing notation for single and binary lenses, we define conventions, summarize all known microlensing degeneracies and extend a set of parameters to describe full Keplerian motion of the binary lenses. ", "introduction": "\\label{s:intro} Gravitational microlensing is nowadays a well established method for discovering binary and planetary systems \\cite[\\eg][]{gould09,gaudi10}. By analyzing flux variations in time, microlensing measures a wide range of system parameters including the distance and mass of the binary, its orbital and proper motion as well as the mass ratio and separation of its components \\cite[\\cf][]{dong09a, bennett10}. Previous studies of gravitational microlensing events have had little or no possibility for {\\it post factum} observational confirmation of derived system parameters, since usually the stars involved in this one-time event are too faint and too distant to be within reach of current astrometric or spectroscopic instruments. Some exceptions come from astrometric confirmation of the nature of a single lens event with direct imaging done with Hubble Space Telescope (HST) \\citep{alcock01, gould04b, kozlowski07}. Also in some cases it was possible to take spectra of the microlensed source that confirm the microlensing interpretation \\citep{gaudi08b}. But the fact of an event being caused by microlensing is the main subject of these tests rather than the values of previously derived parameters. This does not mean that the microlensing measurements have no tests. Indeed there are many self-consistency checks including the agreement of mass and distance with the amount of light coming to us \\cite[\\eg][]{gaudi08a,bennett10}, as well as testing of the derived values of the orbital parameters being consistent with bounded Keplerian orbits (\\eg\\ \\citealt{sumi10}, Batista \\etal\\ 2011 in prep., Yee \\etal\\ 2011 in prep.). The only issue is the shortage of possibilities to verify the results with independent or direct observations. In this work we present analysis of one microlensing event (OGLE-2009-BLG-020\\footnote{\\url{http://ogle.astrouw.edu.pl/ogle3/ews/2009/blg-020.html}}) caused by a $\\sim 1.1 \\, M_\\Sun$ binary system passing near the line of sight to an ordinary red giant star. We derive system parameters using the same, standard methods that are applied to other binary and planetary microlensing events. Because of a peculiarity of this event, \\ie, that it is caused by a relatively close by ($\\approx 1.1$ kpc) and bright ($I=15.6$) binary star, there is a possibility of direct verification of the derived parameters with follow-up spectroscopic measurements. The data gathered on this event are described in Section \\ref{s:data}. The fitting of the microlensing model to the light curve is presented in Section \\ref{s:modelling}, and the physical parameters of the system are calculated in Section \\ref{s:params}. In Section \\ref{s:discussion} we discuss the results and, in particular, in Section \\S\\ref{s:test} we present how the microlensing solutions can be tested by radial velocity measurements. In the Conclusions (\\S\\ref{s:conclusions}) we advocate for radial velocity follow-up observations to confirm the nature of the event and its parameters, which constitute a general test of the microlensing method, in particular the accuracy of the parameters currently being derived for microlensing planets. In Appendix \\ref{a:notation} we present microlensing parameters together with all required conventions, introduce a uniform microlensing notation, extend work of \\citet{gould00} by introducing parameters describing full Keplerian motion of the lens components, and review symmetries known in microlensing. In Appendix~\\ref{a:jacobian} we derive the transformation between the microlensing and Keplerian orbit parameters. ", "conclusions": "\\label{s:conclusions} The binary star that manifested itself in the microlensing event OGLE-2009-BLG-020 is the first case of a lens that is close enough and bright enough to allow ground-based spectroscopic follow-up observations. This makes it a unique tool to test the microlensing solution. We derive lens parameters using the same method by which the majority of planetary candidates discovered by microlensing are analyzed. We detect a signal from the orbital motion of the lens in the microlensing light curve. This signal, as well as our measurements of the orbital parameters of the binary lens, can be confirmed or contradicted by future observations. We propose a test in \\S\\ref{s:test}. Combining the microlensing solution with the radial velocity curve will yield a complete set of system parameters including 3-d Galactic velocity of the binary and all Keplerian orbit elements. This work undertakes an effort to establish a uniform microlensing notation, extending work of \\citet{gould00} by including the full set of orbital elements of the binary lens (Appendix \\ref{a:notation}). We also summarize all known microlensing symmetries and degeneracies. The method of deriving orbital elements from the 6 phase-space coordinates, used to parametrize microlensing event, is described in Appendix \\ref{a:jacobian}. The Fortran codes we use for transformation of the microlensing parameters to orbital elements and for deriving all quantities described in the Appendix \\ref{a:jacobian} will be attached to {\\it astro-ph} sources of this paper, and will be published on the author's web page\\footnote{\\url{http://www.astronomy.ohio-state.edu/\\textasciitilde{}jskowron/OGLE-2009-BLG-020/}}." }, "1101/1101.3581_arXiv.txt": { "abstract": "We present observations and interpretation of the Type IIn supernova SN~2008am discovered by the ROTSE Supernova Verification Project (RSVP). SN~2008am peaked at approximately -22.3~mag at a redshift of $z =$~0.2338, giving it a peak luminosity of $\\sim 3 \\times 10^{44}$~erg~s$^{-1}$ and making it one of the most luminous supernovae ever observed. The total radiated energy is $\\simeq 2 \\times 10^{51}$~erg. The host galaxy appears to be an SB1 of normal luminosity ($M_{r^{\\prime}} \\sim$~-20) with metallicity $Z \\sim $~0.4~$Z_{\\odot}$. ROTSE upper limits and detections constrain the rise time to be $\\sim$~34 days in the rest frame, significantly shorter than similar events, SN~2006gy and SN~2006tf. Photometric observations in the ultraviolet, optical and infrared bands ({\\it J,H,K$_{s}$}) constrain the SED evolution. We obtained six optical spectra of the supernova, five on the early decline from maximum light and a sixth nearly a year later plus a very late-time spectrum ($\\sim$~2 yr) of the host galaxy. The spectra show no evidence for broad supernova photospheric features in either absorption or emission at any phase. The spectra of SN~2008am show strong Balmer-line and He I $\\lambda$5876 \\AA\\ emission with intermediate widths ($\\sim$ 25 \\AA) in the first $\\sim$~40 days after optical maximum. The width formally corresponds to a velocity of $\\sim$~1000~km~s$^{-1}$. We examine a variety of models for the line wings and conclude that multiple scattering is most likely, implying that our spectra contain no specific information on the bulk flow velocity. We examine a variety of models for the ROTSE light curve subject to the rise time and the nature of the spectra, including radioactive decay, shocks in optically-thick and optically-thin circumstellar media (CSM) and a magnetar. The most successful model is one for which the CSM is optically-thick and in which diffusion of forward shock-deposited luminosity gives rise to the observed light curve. The model suggests strong mass loss and a greater contribution from the interaction of the forward shock with optically thick CSM than from the reverse shock. Diffusion of the shock-deposited energy from the forward shock is found to be important to account for the rising part of the light curve. Although there are differences in detail, SN~2008am appears to be closely related to other super-luminous Type IIn supernovae, SN~2006gy, SN~2006tf and perhaps SN~2008iy, that may represent the deaths of very massive LBV-type progenitors and for which the luminosity is powered by the interaction of the ejecta with a dense circumstellar medium. ", "introduction": "The Texas Supernova Search (TSS; Quimby et al. 2005) and its successor, the ROTSE-Supernova Verification Project (RSVP; Yuan et al. 2007), discovered a new class of super-luminous supernovae (SLSNe). The advantage of the TSS/RSVP project is that it is essentially free of selection bias and the limits of a targeted search. The automated wide field (3.4 square degree) ROTSE-III telescopes (Akerlof et al. 2003) scan the whole sky, looking for transients down to $\\sim$19~mag. They do not focus on pre-selected galaxies nor omit galaxy centers. The first TSS/RSVP discoveries in this new class of SLSNe were SN~2005ap (Quimby et al. 2007a), SN~2006gy (Quimby 2006; Smith et al. 2007), SN~2006tf (Quimby, Castro \\& Mondol 2007; Quimby et al. 2007b; Smith et al. 2008) and SN~2008es (Yuan et al. 2008b; Gezari et al. 2009; Miller et al. 2009). These exceptionally luminous supernovae are rare, with an estimated rate of $\\sim 2.6 \\times 10^{-7}$~events~Mpc$^{-3}$~yr$^{-1}$ (Quimby et al. 2009). The SLSNe introduced new modes of stellar death. Traditional ideas about the mechanisms that can power supernova luminosity were found to be inadequate to explain the observed properties of these events. The small, but growing, sample of SLSNe is heterogeneous. Some show strong emission lines of hydrogen in their spectra close to maximum light (SN~2006gy, SN~2006tf, SN~2008fz, SN~2008iy) and typically belong to the Type~IIn subclass; some show hydrogen in later phases and a linear decline of the light curve expressed in magnitudes (SN~2008es). Others may show no hydrogen at all (SN~2005ap, SCP06F6). For the super-luminous Type~IIn events, the energy generation mechanism is very likely the interaction between the ejecta and a circumstellar medium (CSM) that was shed by the progenitor star in the years prior to the explosion (Chevalier \\& Fransson 2003). SN~2006gy triggered discussions about the possibility of nearby pair-instability supernovae (Smith et al. 2007). Such models proved unsatisfactory for SN~2006gy and many other SLSNe, but may account for SN~2007bi (Gal-Yam et al. 2009; Young et al. 2010). Even for some events that do not show clear signs of CSM interaction, simple radioactive decay diffusion models (Arnett 1982; Valenti et al. 2008) have proven inconsistent with the observations (Quimby et al. 2007a; Gezari et al. 2009; Quimby et al. 2009). Other mechanisms that can account for the large luminosity have been proposed: interaction between expelled shells (Woosley, Blinnikov \\& Heger 2007); interaction between a GRB-like jet and the progenitor envelope (Young et al. 2005; Gezari et al. 2009); a buried magnetar (Kasen \\& Bildsten 2010; Woosley 2010); or a very energetic core-collapse explosion (Umeda \\& Nomoto 2008; Moriya et al. 2010). In addition, the possibility that many Type~IIn SNe (of normal or high luminosity) have been spectroscopically confused with radio-quiet low-luminosity blazars has been discussed (Filippenko 1989). All recent SLSNe candidates have shown spectroscopic features that are more consistent with SNe. In the present work, we report on SN~2008am discovered by RSVP (Yuan et al. 2008a). The paper is organized as follows. In \\S 2 we present the photometric and spectroscopic observations of SN~2008am and discuss the evolution of its spectral energy distribution (SED). In \\S 3 we consider the nature of the emission-line features and models to account for the line profiles, and in \\S 4 we discuss the applicability of various models to account for the light curve. Finally, in \\S 5 we summarize our conclusions. ", "conclusions": "\\label{disc} We presented an analysis of the available photometric and spectroscopic data of the SLSN 2008am. The spectroscopic signatures of intermediate width H and He emission lines ($\\sim$~25 \\AA) place this SN in the category of Type IIn. SN~2008am was an extremely luminous event, with a peak absolute R-magnitude of $M_{R} \\simeq -$22.3~mag corresponding to a luminosity of $\\sim 2\\times10^{44}$ erg s$^{-1}$, putting it in the ``hall of fame\" of the most luminous SNe ever observed. The host of SN~2008am is a faint extended galaxy with magnitude $\\sim$~20 in the SDSS catalog. At the redshift of $z =$~0.2338, the absolute r$^{\\prime}$-magnitude of the host is $M_{r^{\\prime}} \\sim -$20~mag, which is in the range typical for elliptical and spiral galaxies. The very late (+554d) Keck spectrum of SN~2008am is consistent with an SB1 template spectrum. From the line flux ratios the host has sub-solar metallicity, $Z \\sim $~0.4~$Z_{\\odot}$. We conclude that the host of SN~2008am is a metal-poor, but normal galaxy, not a subluminous dwarf as is the case for many SLSNe (Miller et al. 2009; Drake et al. 2010; Miller et al. 2010). SN~2008am was followed up photometrically from the IR to the UV. The ROTSE light curve provides a reasonably accurate estimate of the rise time to maximum light of 34 d in the rest frame, a significantly short time compared to other SLSNe. The photometric coverage allowed us to create broad-band SEDs of SN~2008am for 6 epochs. We fit single temperature black-body curves to the SEDs to study the evolution of the black-body temperature and radius as well as to estimate a pseudo-bolometric light curve. The derived black-body temperatures ($\\sim$~10,000-12,000~K) are consistent with the continua of contemporaneous spectra. These temperatures are very high compared with typical core-collapse supernova photospheres (5,000-6,000 K) as well as with the temperatures obtained for SN~2006gy and SN~2006tf (6,000-8,000 K; Smith et al. 2008, 2010). The single-temperature black-body fits for SN~2008am were imperfect, implying that the underlying emission mechanism is more complex in nature. Spectra obtained about 10 - 30 rest-frame days after maximum light showed intermediate-width emission lines of H$\\alpha$, H$\\beta$, H$\\gamma$, and a feature at the HeI/Na D blend. There is no sign in our data of broad P Cygni lines that might signify the photosphere of the underlying supernova nor of narrow P Cygni lines as displayed by SN~2006gy and SN~2008tf on decline that indicate absorption in the unshocked, but expanding circumstellar matter (Smith et al. 2008, 2010). A spectrum obtained 352 days after our estimated maximum showed only H$\\alpha$ that was significantly narrower than in the earlier spectra. Our failure to detect P Cygni features is most likely attributable to our lack of data at phases fainter than 2 magnitudes from maximum when SN~2006gy began to show such features, although we cannot rule out issues of S/N ratio and wavelength resolution. The overall observed spectral evolution of SN~2008am is similar to that of SN~2006gy at similar phases, and we conclude that they are closely related. We considered a variety of models for the emission line profiles: Gaussian as might typify thermal Doppler broadening, an exponential profile that might characterize single electron scattering in an optically-thin medium, and Lorentzian that might represent models of multiple electron scattering (Chugai 2001; Smith et al. 2010). We find that the latter provides the best fit to the overall line shape. An important implication is that the line broadening is probably dominated by electron scattering in our spectra and that the line width contains little or no information about the bulk kinetic expansion velocity of the matter in the circumstellar medium or the underlying supernova. An upper limit to the velocity in the line-forming region is about 1,000 km s$^{-1}$. Chugai made specific assumptions in his models, for instance that the velocity profile decreased outward as might result from radiative acceleration, that might not apply in general. The line profiles might contain information about the velocity structure even if scattering dominated (Fransson \\& Chevalier 1989). The radiative transfer that results in these line profiles is worthy of re-examination. We explored a number of light curve models based on a generalization of the models of Arnett (1980; 1982) that use a specified power input and the first law of thermodynamics coupled with the diffusion equation. The models are constrained by the rise time, the quasi-bolometric light curve and the general nature of the emission line spectra that show intermediate-width lines with little or no sign of P Cygni absorption on any scale. We examined models based on radioactive decay, an underlying supernova striking an optically thick circumstellar shell, a supernova shock in an optically-thin CSM, and a magnetar. The radioactive decay model is ruled out because the required nickel mass would exceed the deduced ejecta mass. The magnetar model provides a decent fit to the light curve, but no natural explanation for the emission-line structure. A shell-shock model similar to that of Smith \\& McCray (2007) for SN~2006gy in which the diffusion times on the rise and the decline are equal fails drastically for SN~2008am, primarily because of the rapid rise, 34 d in the rest frame. We conclude that the success of a model with a single diffusion time for SN~2006gy was a coincidence and not a general property of this class of events. We could generate a reasonable fit in the context of this model by including an input power source representing the collision of the supernova with an optically-thick shell that was a ``top hat\" function of constant luminosity during the rise that shut off at 34 days, the point of maximum light in this model. In this model, the timescale of the power source dictated the rise time and a separately determined diffusion time of 120d governed the decline. This model rose with $L \\propto t^2$ at very early times (as assumed for the full rise of SN~2006gy by Smith \\& McCray) and was overall concave upward in contrast to the observed light curve that appears to be concave downward. The shape of the rise depends on the input power profile. We will investigate more general models in a future paper. The optically-thin model is not an Arnett-like diffusion model, but assumes that the shock energy is rapidly radiated. This model did not provide a self-consistent fit to SN~2008am and gives no natural explanation for the rise nor for the failure to see any high-velocity features corresponding to the photosphere of the underlying supernova. The best fit to the ROTSE light curve data was obtained with an ejecta-CSM interaction model in which the supernova is not seen directly and the luminosity is produced by shocks and diffusion in a circumstellar medium that is optically thick. For fiducial parameters of $v = 1,000$~km~s$^{-1}$ and $\\kappa = 0.4$~cm$^{2}$~g$^{-1}$, this model gives a rather small initial radius for the shell, $R_{0} \\sim 1.0 \\times 10^{14}$ cm, with a mass of about 1~M$_{\\odot}$ and a total energy input from the underlying SN/CSM shock of $E_{s} = 5.5 \\times 10^{51}$ erg. The model suggests that the optically-thick component may dominate the luminosity and that the forward shock provides a greater contribution to the luminosity than the reverse shock. This model suggests a rather large mass loss rate for the progenitor, as perhaps would be consistent with an LBV-type mass loss process. As noted above, SN~2008am seems to be a close cousin of SN~2006gy, showing Lorentzian emission lines shortly after maximum and a narrower H$\\alpha$ line about a year later. The Lorentzian lines in both events show a slight redshift of about 100 km s$^{-1}$. The most notable difference is our failure to see the distinct broader (4,000 km s$^{-1}$) and narrow (200 km s$^{-1}$) P Cygni features that appeared in SN~2006gy 20 to 80 days after maximum light. Smith et al. (2010) attribute the first phase of pure emission to conditions where the shock is still beneath the photosphere of the dense CSM and the second phase where P Cygni features form to conditions where the shock has proceeded beyond the photosphere. The high-velocity absorption is presumably related to the SN ejecta, and the low-velocity absorption to the motion of the CSM that has not yet been hit by the shock, but is subject to radiative excitation and recombination. The most likely explanation of our failure to detect these P Cygni features is absence of data at the appropriate phase. The narrower H$\\alpha$ line in both events nearly a year after explosion is consistent with the emitting matter becoming more dilute with less broadening by multiple electron scattering. Other SLSNe seem to fall broadly in the category of SN~2006gy and SN~2008am. SN~2006tf shows nearly symmetric emission lines, especially of H$\\alpha$, up to 40 days after discovery (there is no data on the rise so the explosion date and date of maximum are uncertain), with the development of narrow P Cygni features by 66 days after discovery (Smith et al. 2008). At these later epochs, there are indications in both the emission and absorption for rapidly moving material, $\\sim 7,500$ km s$^{-1}$, presumably from the underlying SN ejecta. SN~2006tf is somewhat different from SN~2006gy and SN~2008am in the late phases, a year after explosion, where the H$\\alpha$ line seems to be formed by collisional rather than radiative excitation (Smith et al. 2008). The H$\\alpha$ line in SN~2006tf at this stage shows a prominent blue ``plateau\" extending to about 1,000~km~s$^{-1}$. SN~2008iy was an unusual SN IIn with the unprecedented slow rise time of 400 days (Miller et al. 2010). The post-maximum spectra of SN~2008iy are somewhat similar to those of SN~2008am, with strong intermediate-width H and He emission lines. A single temperature black body failed to provide a good fit to the SEDs of SN~2008iy, as we found for the SED of SN~2008am (\\S 2.3). Miller et al. identified three distinct components in the late-time H$\\alpha$ profile of SN~2008iy: broad ($\\sim$~4,500~km~s$^{-1}$), intermediate ($\\sim$~1,650~km~s$^{-1}$) and narrow ($\\sim$~75~km~s$^{-1}$). Miller et al. proposed a model of interaction of the SN ejecta with a clumpy circumstellar medium similar to that for SN~1988Z presented by Chugai \\& Danziger (1994). Miller et al. argued that the rise in the light curve resulted from an increase in the number of clumps with radius. Given the success of the shell-shock models of SN~2006gy and SN~2008am, it would be of interest to apply such a model to SN~2008iy. As for SN~2006gy and SN~2006tf, the suggestion of a relatively massive shell around SN~2008am and the estimated mass loss rates ($\\sim$~0.1-10~$M_{\\odot}$~yr$^{-1}$) imply that the progenitor star must have undergone substantial mass-loss in the years prior to the explosion. Episodic mass loss can occur around very massive LBV stars, similar to $\\eta$ Carinae. The reason for LBV mass-loss is not currently fully understood (Smith \\& Owocki 2006). Massive shell ejection can also be the product of pulsational pair-instability (Rakavy \\& Shaviv 1967; Barkat, Rakavy \\& Sack 1967). Models suggest that repetitive shell ejection takes place for progenitor main sequence masses in the range 95-130~$M_{\\odot}$ (Woosley, Blinnikov \\& Heger 2007). Supernova-like luminosity can be produced either during the ejection of each of these shells individually (Kasen et al. 2008), or during the collisions between shells ejected at different times (Woosley, Blinnikov \\& Heger 2007). Although there is some sign of high-velocity material in some of the SLSNe that otherwise resemble SN~2008am, the nature of the presumed underlying supernova in these Type IIn SLSNe remains obscure. As noted in the Introduction, other SLSNe show little or no evidence for hydrogen or interaction with a CSM. In the case of SN~2007bi, this is an important part of the argument that it is a pair-instability supernova (Gal-Yam et al. 2009). An important goal in the study of SLSNe remains the determination of the density distribution in the CSM that will give important clues to the mass-loss history. We are grateful to the anonymous referee for valuable guidance on style and science and to Andy Howell and Milos Milosavljevic for useful discussions. This research is supported in part by NSF Grant AST-0707669 and by the Texas Advanced Research Program grant ASTRO-ARP-0094. E. Chatzopoulos would like to thank the Propondis foundation of Piraeus, Greece for its support of his studies. J. Vinko received support from Hungarian OTKA Grant K76816." }, "1101/1101.1067_arXiv.txt": { "abstract": "We present new observations of the CB130 region, composed of three separate cores. Using the \\textit{Spitzer Space Telescope} we detected a Class 0 and a Class II object in one of these, CB130-1. The observed photometric data from \\textit{Spitzer} and ground-based telescopes are used to establish the physical parameters of the Class 0 object. SED fitting with a radiative transfer model shows that the luminosity of the Class 0 object is 0.14 $\\--$ 0.16 \\lsun, which is a low luminosity for a protostellar object. In order to constrain the chemical characteristics of the core having the low luminosity object, we compare our molecular line observations to models of lines including abundance variations. We tested both ad hoc step function abundance models and a series of self-consistent chemical evolution models. In the chemical evolution models, we consider a continuous accretion model and an episodic accretion model to explore how variable luminosity affects the chemistry. The step function abundance models can match observed lines reasonably well. The best fitting chemical evolution model requires episodic accretion and the formation of CO$_2$ ice from CO ice during the low luminosity periods. This process removes C from the gas phase, providing a much improved fit to the observed gas-phase molecular lines and the CO$_2$ ice absorption feature. Based on the chemical model result, the low luminosity of CB130-1 is explained better as a quiescent stage between episodic accretion bursts rather than being at the first hydrostatic core stage. ", "introduction": "The \\textit{Spitzer Space Telescope} Legacy Project, From Molecular Cores to Planet Forming Disks (c2d, \\citealt{2003PASP..115..965E}) has completed a survey of nearby star-forming regions. It covered five clouds, containing 1024 objects classified as young stellar Objects (YSOs) \\citep{2009ApJS..181..321E}. When a dense core forms a central protostar, it develops a luminosity source from the accretion of infalling mass. The accretion luminosity is given as $L_{acc} = G M_* \\dot{M}_{acc} / R_*$, where $M_*$ is the mass of the protostar, $\\dot{M}_{acc}$ is the mass accretion rate onto the protostar, and $R_*$ is the radius of the protostar. According to the standard model, the mass accretion rate from spherical infall of the envelope is $\\dot{M}_{acc} \\simeq 2 \\times 10^{-6}$ \\msun yr$^{-1}$ if the infall occurs at the thermal sound speed at 10 K \\citep{1977ApJ...214..488S}. With a typical protostellar radius of 3 R$_{\\odot}$, and a mass at the stellar/brown dwarf boundary of 0.08 \\msun, the resulting $L_{acc}$ is 1.6 \\lsun. Any luminosity from contraction of the forming star will add to this accretion luminosity, making it a minimum value in the standard model. In contrast to the predictions, the c2d survey showed that 59\\% of the 112 embedded protostars (Class 0/I) have luminosity lower than 1.6 \\lsun ~\\citep{2009ApJS..181..321E}, indicating that either $M_*$ is very small, or $\\dot{M}_{acc}$ is lower than expected, or both. This luminosity problem was aggravated by the discovery of Very Low Luminosity Objects (VeLLOs, \\citealt{2007prpl.conf...17D}). VeLLOs are defined as embedded protostars with internal luminosity lower than 0.1 \\lsun. The internal luminosity of an embedded protostar is the luminosity of the protostar and disk, excluding luminosity from external heating. By using the correlation between 70 $\\micron$ flux and the internal luminosity of protostars, \\citet{2008ApJS..179..249D} identify 15 VeLLO candidates in the full c2d sample. Several VeLLOs including L1014-IRS \\citep{2004ApJS..154..396Y}, L1148-IRS \\citep{2005AN....326..878K}, L1521F-IRS \\citep{2006ApJ...649L..37B}, IRAM 04191-IRS \\citep{1999ApJ...513L..57A, 2006ApJ...651..945D}, L328-IRS \\citep{2009ApJ...693.1290L}, and L673-7 \\citep{2010ApJ...721..995D}, have been studied in detail. Among those VeLLOs, IRAM 04191$+$1522 and L673-7 drive strong molecular outflows \\citep{1999ApJ...513L..57A, 2006ApJ...651..945D, 2010ApJ...721..995D}. A low accretion luminosity with a significant molecular outflow can imply higher $\\dot{M}_{acc}$ in the past \\citep{2010ApJ...721..995D}. One possible explanation for the sources with low luminosities combined with strong molecular outflows is that the mass accretion is not a constant process but rather episodic, and the VeLLOs are in quiescent phases between the mass accretion bursts \\citep{1990AJ.....99..869K, 2009ApJ...692..973E, 2010ApJ...710..470D}. \\citet{2010ApJ...710..470D} present a set of evolutionary models describing collapse including episodic accretion. The luminosity distribution of YSOs can be matched only when the model includes episodic accretion. While dust continuum emission can provide the physical structure and evolutionary stage of a core, molecular line observations trace the dynamics and chemistry of the core. The simple empirical models of step function \\citep{2003ApJ...583..789L} and drop function \\citep{2004A&A...416..603J} abundance profiles have been used to approximate the abundance profiles. The chemical evolution model during protostellar collapse developed by \\citet{2004ApJ...617..360L} has been tested for individual cores, like B335 \\citep{2005ApJ...626..919E}, and L43 \\citep{2009ApJ...705.1160C}. Using this model, \\citet{2007JKAS...40...85L} studied chemical evolution in VeLLOs. The low luminosity sources can provide chemical laboratories to test the astrochemistry of low luminosity environments. Because the chemical equilibrium time can be longer than the duration of an accretion burst, the abundance profile may provide a fossil record of previous luminosity bursts. This paper presents new observations of the CB130 region. With a detailed analysis of \\emph{Spitzer}, submm continuum, and molecular line data we will reveal the physical properties of this source and use CB130 as a laboratory to study the chemistry in low luminosity sources. In \\S \\ref{CB130-1} we give a general introduction to CB130. The observations are described in \\S \\ref{obs_sec}. \\S \\ref{result_sec} presents images, photometry, and molecular line data. \\S \\ref{radmodel_sec} presents dust continuum radiation transfer models used to determine physical parameters. The line modeling with an empirical step function is described in \\S \\ref{line_sec}. The chemical evolution models with different luminosity evolution are presented in \\S \\ref{chemical_evol_model}, and we summarize our findings in \\S \\ref{summary}. ", "conclusions": "\\label{summary} We presented a detailed study of a low luminosity object in the CB130-1 region. We performed radiative transfer modeling and chemical modeling to explain the core with a low luminosity protostar in it. The embedded protostar CB130-1-IRS1 has $L_{int} = 0.14$ \\lsun\\ (1-D model) to $L_{int} = 0.16$ \\lsun\\ (2-D model). This is a slightly higher luminosity than a VeLLO has, but still CB130-1-IRS1 is a low luminosity source at about 0.1 of the expected luminosity for steady accretion onto an object at the boundary between stars and brown dwarfs. We tested both a step function abundance model and self-consistent chemical evolution models. The step function abundance profile fits observed lines reasonably well, but is of course, ad hoc. For the self-consistent models, we use three different luminosity evolution scenarios to explain the low luminosity of the object. All three luminosity evolution models with the standard chemical network show that the modeled C$^{18}$O line is strong and N$_2$H$^+$ is weak compared to the observations. The step function model has more free parameters than chemical evolution model, so it is more feasible to fit the observed lines. The deep 15.2 $\\micron$ CO$_2$ ice feature indicates that CO$_2$ ice has formed from CO ice. We added a reaction to the chemical network, which turns CO ice into CO$_2$ ice. A model with that reaction decreases the C$^{18}$O abundance and increases the N$_2$H$^+$ abundance. With the episodic accretion model and the modified network, we can explain the low luminosity of the source, the strong CO$_2$ ice feature, and most of the gas phase emission lines at the same time. So CB130-1 is most likely neither a FHSC nor a very low mass object, but a more evolved protostar in a quiescent stage between accretion bursts. With a further study of CO$_2$, CO and N$_2$H$^+$ in low luminosity objects, we may find chemical imprints of episodic accretion in low luminosity objects. We thank the Lorentz Center in Leiden for hosting several meetings that contributed to this paper. Support for this work, part of the \\textit{Spitzer} Legacy Science Program, was provided by NASA through contracts 1224608 and 1288664 issued by the Jet Propulsion Laboratory, California Institute of Technology, under NASA contract 1407. Support was also provided by NASA Origins grant NNX07AJ72G and NSF grant AST-0607793 to the University of Texas at Austin. This research was also supported by the National Research Foundation of Korea (NRF) grant funded by the Korea government (MEST) (No. 2009-0062866) and by Basic Science Research Program through the NRF funded by the Ministry of Education, Science and Technology (No. 2010-0008704). TLB was partially supported by NASA through contracts 1279198, 1288806, and 1342425 issued by the Jet Propulsion Laboratory, California Institute of Technology to the Smithsonian Astrophysical Observatory, and by the NSF through grant AST-0708158." }, "1101/1101.1298_arXiv.txt": { "abstract": "The vast majority of Type II supernovae (SNe) are produced by red supergiants (RSGs), but SN\\,1987A revealed that blue supergiants (BSGs) can produce members of this class as well, albeit with some peculiar properties. This best studied event revolutionised our understanding of SNe, and linking it to the bulk of Type II events is essential. We present here optical photometry and spectroscopy gathered for SN\\,2000cb, which is clearly not a standard Type II SN and yet is not a SN\\,1987A analog. The light curve of \\cb\\ is reminiscent of that of SN\\,1987A in shape, with a slow rise to a late optical peak, but on substantially different time scales. Spectroscopically, SN\\,2000cb resembles a normal SN\\,II, but with ejecta velocities that far exceed those measured for SN\\,1987A or normal SNe\\,II, above 18,000\\,km\\,s$^{-1}$ for H$\\alpha$ at early times. The red colours, high velocities, late photometric peak, and our modeling of this object all point toward a scenario involving the high-energy explosion of a small-radius star, most likely a BSG, producing 0.1\\,$\\Msun$ of $^{56}$Ni. Adding a similar object to the sample, SN\\,2005ci, we derive a rate of $\\sim 2$\\% of the core-collapse rate for this loosely defined class of BSG explosions. ", "introduction": "Massive stars that retain their hydrogen envelope during their evolution end their lives as Type II supernovae (SNe\\,II). The vast majority of these are characterised by a fast (few days) rise to a flat light curve, most pronounced in the reddest optical bands, with a duration of 80--100\\,d. This ``plateau'' phase, for which they have been named SNe\\,II-P, is interpreted as the recession of the photosphere as the ejecta expand and cool \\citep[e.g.,][]{kirshner73, barbon79}. The spectra of SNe\\,II-P are typically dominated by strong P-Cygni profiles of hydrogen lines, as well as iron absorption features \\citep[e.g., see the review by][]{filippenko97}. Due to the relatively simple physics driving their\\citet{} optical evolution, SNe II-P were the first SNe to be suggested as distance indicators, via the so-called expanding photosphere method (EPM; \\citealt{kirshner74}). While this method and its descendants depend on modeling, \\citet{hamuy02} suggested an empirical relation between luminosity and ejecta velocity, as measured from iron absorption lines, that proved compelling when combined with dust-extinction correction \\citep{nugent06,olivares10,poznanski09,poznanski10}. Over the past decade, $\\sim 20$ pre-explosion locations of SNe\\,II-P have been directly imaged with the {\\it Hubble Space Telescope} or deep ground-based images, yielding five detections of progenitor stars, all of which were red supergiants (RSGs), and many limits on stars with masses in the range 7.5--15 $\\Msun$ \\citep[][ and references therein]{smartt09}. These masses are somewhat at odds with those derived from explosion modeling, which tend to be higher, closer to 15--25 $\\Msun$ \\citep[e.g.,][]{nadyozhin03,utrobin07,utrobin09}, though \\citet{dessart10b} push toward lower masses of less than $\\sim 20$\\,$\\Msun$. SNe\\,II-P show some diversity, but overall they tend to be a fairly homogeneous group. However, there are intriguing relatives to this class of objects, most notably SN\\,1987A, which exploded in the nearby Large Magellanic Cloud and whose progenitor was a compact blue supergiant (BSG) star \\citep[][ and references therein]{arnett89}. Given the unique contribution of SN\\,1987A to the study of core-collapse SNe, particular attention to objects that resemble it is warranted. The most similar published SN was SN\\,1998A \\citep{pastorello05}. While its light-curve shape was nearly identical to that of SN\\,1987A, SN\\,1998A was more luminous and bluer, and its spectra showed higher expansion velocities. \\citet{pastorello05} attribute these differences to a higher-energy BSG explosion. \\cb\\ was discovered on 27.4 April 2000 (UT dates are used throughout this paper) by \\citet{papenkova00} in the spiral galaxy IC\\,1158 at $\\alpha = 16^{\\rm h}01^{\\rm m}32.15^{\\rm s}$, $\\delta = +1^\\circ42\\arcmin23\\arcsec.0$ (J2000) as part of the Lick Observatory Supernova Search \\citep{li00,filippenko01}. An unfiltered image obtained with the Katzman Automatic Imaging Telescope (KAIT) on 24.4 April shows the supernova, while there is no detection to a limiting magnitude of 18.7 in an image taken on 9.5 April. The SN is at a projected distance of about 4.3\\,kpc from the host-galaxy centre, off one of the outer arms of this spiral galaxy. \\citet{jha00} classified it spectroscopically as a SN\\,II on 28.4 April and noted high expansion velocities measured from the hydrogen absorption features (up to 18300 km\\,s$^{-1}$ for H$\\alpha$), as confirmed on 29.3 April by \\citet{aldering00}. Some peculiarities of \\cb\\ and its similarity to SN\\,1987A were noted by \\citet{hamuy01b}, and as a result, both SNe were excluded from a sample of SNe\\,II-P used as standardizable candles \\citep{hamuy04}. In this paper, we present photometric and spectroscopic data on SN\\,2000cb, all taken within 160~d after explosion. We investigate its properties and compare it to the prototypical Type II SN\\,1999em, to SN\\,1987A, and to SN\\,1998A. As we show, while \\cb\\ shares some characteristics with SNe\\,1987A and 1998A, it does not appear to closely match any of our comparison objects, further expanding the range of known possible outcomes of massive stellar death. Various arguments, in addition to our model fit to the bolometric evolution, point to a BSG progenitor that produced a strong explosion with a small envelope and a significant amount of $^{56}$Ni. ", "conclusions": "We have presented data and analysis of \\cb, a peculiar Type II SN, with an atypical light curve and extreme photospheric velocities. \\cb\\ strongly resists our attempts to cast it into the standard SN\\,II-P category while also avoiding identity with SN\\,1987A. Photometrically it is intrinsically redder than SN\\,1999em and fainter by about 0.5\\,mag on the plateau. It rises to maximum brightness much more slowly than a normal SN\\,II-P, yet this rise is faster than that of SN\\,1987A and proceeds with a distinctively different colour and luminosity evolution. The main spectroscopic peculiarity lies in the exceptionally high velocities of its features at all times, most notably for the hydrogen absorption. Our modeling, as well as extensive comparisons to other SNe, favour a high-energy explosion of a relatively small-radius star, most probably a BSG. We derive a rate for BSG explosions that is on the order of 2\\% of the rate of all core-collapse SNe. SNe\\,II-P are the most abundant SNe in the Universe and originate from the explosion of RSGs, the most abundant evolved massive stars. SN\\,1987A is the keystone of modern Type II SN theory, despite being the explosion of a BSG. Linking peculiar SNe such as SN\\,1987A and \\cb\\ to regular SNe\\,II is critical for elucidating the processes that bring a star from formation to obliteration." }, "1101/1101.2453_arXiv.txt": { "abstract": "Many experiments in the near future will test dark energy through its effects on the linear growth of matter perturbations. In this paper we discuss the constraints that future large-scale redshift surveys can put on three different parameterizations of the linear growth factor and how these constraints will help ruling out different classes of dark energy and modified gravity models. We show that a scale-independent bias can be estimated to a few percent per redshift slice by combining redshift distortions with power spectrum amplitude, without the need of an external estimation. We find that the growth rate can be constrained to within 2-4\\% for each $\\Delta z=0.2$ redshift slice, while the equation of state $w$ and the index $\\gamma$ can be simultaneously estimated both to within 0.02. We also find that a constant dimensionless coupling between dark energy and dark matter can be constrained to be smaller than 0.14. ", "introduction": "The linear growth rate of matter perturbations is one of the most interesting observable quantities since it allows to explore the dynamical features related to the build-up of cosmic structures beyond the background expansion. For example it can be used to discriminate between cosmological models based on Einstein's gravity and alternative models like $f(R)$ modifications of gravity (see e.g.~\\cite{defelice10}) or multi-dimensional scenarios like in the Dvali-Gabadaze-Porrati (DGP)~\\cite{dvali00} theory (e.g.~\\cite{wang08} and references therein). In addition, the growth rate is sensitive to dark energy clustering or to dark energy-dark matter interaction. For instance, in models with scalar-tensor couplings or in $f(R)$ theories the growth rate at early epochs can be larger than in $\\Lambda$CDM models and can acquire a scale dependence \\cite{diporto08,tsujikawa09,gannouji09} (see for instance \\cite{amendola_book} for a review on dark energy). Simultaneous information on geometry and growth rate can be obtained by measuring the galaxy power spectrum or the 2-point correlation function and their anisotropies observed in redshift space. These redshift distortions arise from peculiar velocities that contribute, together with the recession velocities, to the observed redshift. The net effect is to induce a radial anisotropy in galaxy clustering that can be measured from standard two-point statistics like the power spectrum or the correlation function~\\cite{hamilton98}. The amplitude of the anisotropy is determined by the typical amplitude of peculiar velocities which, in linear theory, is set by the growth rate of perturbations\\footnote{In order to avoid confusion with the $f(R)$ models, we use the letter $s$, for slope, rather than the more popular $f$, to indicate the growth rate.}: \\begin{equation} s\\equiv\\frac{d\\log G}{d\\log a}\\,,\\end{equation} where $G(z)\\equiv\\delta(z)/\\delta(0)$ is the growth function, $\\delta(z)$ the matter density contrast and the scale factor $a$ is related to the redshift $z$ through $a=(1+z)^{-1}$. Since however we only observe the clustering of galaxies and not that of the matter, the quantity that is accessible to observations is actually \\begin{equation} \\beta\\equiv\\frac{s}{b}\\,,\\label{eq:beta_sb} \\end{equation} where the bias $b$ is the ratio of density fluctuations in galaxies and matter. The bias is in general a function of redshift and scale, but in the following we will consider it as a simple scale-independent function. Once the power spectrum is computed in $k$-space, the analysis proposed in~\\cite{seo03} can be exploited to constrain not only geometry but also the growth rate (as pointed out in~\\cite{amendola05}; see also~\\cite{sapone07,wang08}), provided that the power spectrum is not marginalized over its amplitude. In configuration space, the first analysis of the two-point correlation function explicitly aimed at discriminating models of modified gravity from the standard $\\Lambda$CDM scenario has been performed by~\\cite{guzzo08}. Currently, there are several experimental estimates of the growth factor derived from the analysis of the redshift space distortions~\\cite{hawkins03,verde02,tegmark06,ross07,guzzo08,daangela08,peacock01,blake10}, from the redshift evolution of the {\\it rms} mass fluctuation $\\sigma_8$ inferred from Ly$\\alpha$ absorbers~\\cite{mcdonald05} and from the power spectrum of density fluctuations measured from galaxies' peculiar velocities~\\cite{nusser11}. Current uncertainties are still too large to allow these measurements to discriminate among alternative cosmological scenarios. (e.g.~\\cite{nesseris08,dossett10}). On-going redshift surveys like VIPERS~\\cite{guzzo10} or BOSS~\\cite{BOSS} will certainly provide more stringent constraint and will be able to test those models that deviate most from the standard cosmological model. However, only next generation large-scale redshift surveys at $z\\approx1$ and beyond like EUCLID~\\cite{euclid} or BigBOSS~\\cite{BigBOSS} will provide an efficient way to discriminate competing dark energy models. The growth rate $s$ clearly depends on the cosmological model. It has been found in several works \\cite{peebles76,lahav91,polarski08,linder05,wang98} that a simple yet effective parameterization of $s$ captures the behavior of a large class of models. Putting \\begin{equation} s=\\Omega_{m}^{\\gamma}\\,,\\label{eq:standard} \\end{equation} where $\\Omega_{m}(z)$ is the matter density in units of the critical density as a function of redshift, a value $\\gamma\\approx0.545$ reproduces well the $\\Lambda$CDM behavior while departures from this value characterize different models. For instance the DGP is well approximated by $\\gamma\\approx0.68$~\\cite{linder07,wei08} while viable models of $f(R)$ are approximated by $\\gamma\\approx0.4$ for small scales and small redshifts~\\cite{tsujikawa09,gannouji09}. This simple parameterization is however not flexible enough to accommodate all cases. A constant $\\gamma$ cannot for instance reproduce a growth rate larger than $s=1$ in the past (as we have in $f(R)$ and scalar-tensor models) allowing at the same time $s<1$ at the present epoch if $\\Omega_{m}\\le1$. Even in standard cases, a better approximation requires a slowly-varying, but not strictly constant, $\\gamma$. In addition, the measures of the growth factor obtained from redshift distortions require an estimate of the galaxy bias, which can be obtained either independently, using higher order statistics (e.g.~\\cite{verde02,marinoni05}) or inversion techniques~\\cite{sigad00}, or self consistently, by assuming some reasonable form for the bias function {\\it a priori} (for instance, that the bias is independent of scale, as we will assume here). The goal of this paper is to forecast the constraints that future observations can put on the growth rate. In particular we use representative assumptions for the parameters of the EUCLID survey to provide a baseline for future experiments and we focus on the following issues. {\\it i}) We assess how well one can constrain the bias function from the analysis of the power spectrum itself and evaluate the impact that treating bias as a free parameter has on the estimates of the growth factor. We compare the results with those obtained under the more popular approach of fixing the bias factor (and its error) to some independently-determined value. {\\it ii}) We estimate how errors depend on the parameterization of the growth factor and on the number and type of degrees of freedom in the analysis. {\\it iii}) We explicitly explore the case of coupling between dark energy and dark matter and assess the ability of measuring the coupling constant. We do this in the context of the Fisher Matrix analysis. This is a common approach that has been adopted in several recent works, some of which exploring the case of a EUCLID-like survey as we do. We want to stress here that this work is, in fact, complementary to those of the other authors. Unlike most of these works, here we do not try to optimize the parameter of the EUCLID survey in order to improve the constraints on the relevant parameters, as in~\\cite{wang10}. Instead, we adopt a representative sets of parameters that describe the survey and derive the expected errors on the interesting quantities. In addition, unlike \\cite{simpson10} and ~\\cite{samushia10}, we do not explicitly aim to study the correlation between the parameters that describe the geometry of the system and the growth parameters, although in our approach we also take into account the degeneracy between geometry and growth. Finally, the main results of this paper are largely complementary to the work of~\\cite{majerotto11} that perform a more systematic error analysis that does not cover the main issues of our work. Although, as we mentioned, in general $s$ might depend on scale, we limit this paper to an exploration of time-dependent functions only. Forecasts for specific forms of scale-dependent growth factor motivated by scalar-tensor models are in progress and will be presented elsewhere. The layout of the paper is as follows. In the next section we will introduce the different parameterizations adopted for the growth rate and for the equation of state of dark energy, together with the models assumed for the biasing function, and describe the different cosmological models we aim to discriminate. In sec.~\\ref{sec:fm} we will briefly review the Fisher matrix method for the power spectrum and define the adopted fiducial model. In sec.~\\ref{sec:survey} we will describe the characteristics of the galaxy surveys considered in this work, while in sec.~\\ref{sec:results} we will report our results on the forecast errors on the parameters of interest. Finally, in sec.~\\ref{sec:conclusions} we will draw our conclusions and discuss the results. ", "conclusions": "\\label{sec:conclusions} In this paper we addressed the problem of determining the growth rate of density fluctuations from the estimate of the galaxy power spectrum at different epochs in future redshift survey. As a reference case we have considered the proposed EUCLID spectroscopic survey modeled according to the latest, publicly available survey characteristics~\\cite{euclid,geach10}. In this work we focused on a few issues that we regard as very relevant and that were not treated in previous, analogous Fisher Matrix analysis mainly aimed at optimizing the survey setup and the observational strategy. These issues are: {\\it i}) the ability in measuring self-consistently galaxy bias with no external information and the impact of treating the bias as an extra free parameter on the error budget; {\\it ii}) the impact of choosing a particular parameterization in determining the growth rate and in distinguishing dark energy models with very different physical origins (in particular we focus on the $\\Lambda$CDM, $f(R)$ and the DGP, models that are still degenerate with respect to present growth rate data); {\\it iii}) the estimate of how errors on the growth rate depend on the degrees of freedom in the Fisher matrix analysis; {\\it iv}) the ability of estimating a possible coupling between dark matter and dark energy. The main results of the analysis were already listed in the previous Section, here we recall the most relevant ones. \\begin{enumerate} \\item With the ``internal bias'' method we were able to estimate bias with 1\\% accuracy in a self consistent way using only galaxy positions in redshift-space. The precision in measuring the bias has a very little dependence on the functional form assumed for $b(z)$. Measuring $b$ with 1\\% accuracy will be a remarkable result also from an astrophysical point of view, since it will provide a strong, indirect constraint on the models of galaxy evolution. \\item We have demonstrated that measuring the amplitude and the slope of the power spectrum in different $z$-bin allows to constrain the growth rate with good accuracy, with no need to assume an external error for $b(z)$. In particular, we found that $s$ can be constrained at $1\\sigma$ to within 3\\% in each of the 8 redshift bin from $z=0.5$ to $2.1$. This result is robust to the choice of the biasing function $b(z)$. The accuracy in the measured $s$ will be good enough to discriminate among the most popular competing models of dark energy and modified gravity. \\item Taking into account the possibility of a coupling between dark matter and dark energy has the effect of loosening the constraints on the relevant parameters, decreasing the statistical significance in distinguishing models (from $\\gtrsim 2\\sigma$ to $\\lesssim 1.5\\sigma$). Yet, this is still a remarkable improvement over the present situation, as can be appreciated from Fig.~\\ref{fig:gamma_eta_new_past} where we compare the constraints expected by next generation data to the present ones. Moreover, the {\\it Reference} survey will be able to constrain the parameter $\\eta$ to within 0.04. Reminding that we can write $\\eta=2.1 \\beta_c^2$~\\cite{diporto08}, this means that the coupling parameter $\\beta_c$ between dark energy and dark matter can be constrained to within 0.14, solely employing the growth rate information. This is comparable to existing constraints from the CMB but is complementary since obviously it is obtained at much smaller redshifts. A variable coupling could therefore be detected by comparing the redshift survey results with the CMB ones. \\end{enumerate} It is worth pointing out that, whenever we have performed statistical tests similar to those already discussed by other authors in the context of a EUCLID-like survey, we did find consistent results. Examples of this are the values of FOM and errors for $w_0$, $w_1$, similar to those in~\\cite{wang10,majerotto11} and the errors on constant $\\gamma$ and $w$~\\cite{majerotto11}. However, let us notice that all these values strictly depend on the parameterizations adopted and on the numbers of parameters fixed or marginalized over. In particular, we also found that all these constraints can be improved if one uses additional information from e.g. CMB and other observations. We made a first step in this direction in Fig.~(\\ref{fig:histogram}), which shows how the errors on a constant $\\gamma$ decrease when progressively more parameters are fixed by external priors." }, "1101/1101.5273_arXiv.txt": { "abstract": "The anisotropy of cosmic rays (CRs) in the solar vicinity is generally attributed to the CR streaming due to the discrete distribution of CR sources or local magnetic field modulation. Recently, the two dimensional large scale CR anisotropy has been measured by many experiments in TeV-PeV energy range in both hemispheres. The tail-in excess along the tangential direction of the local spiral arm and the loss cone deficit pointing to the north Galactic pole direction agree with what have been obtained in tens to hundreds of GeV. The persistence of the two large scale anisotropy structures in such a wide range of energy suggests that the anisotropy might be due to a global streaming of the Galactic CRs (GCRs). This work tries to extend the observed CR anisotropy picture from solar system to the whole galaxy. In such a case, we can find a new interesting signature, a loop of GCR streaming, of the GCR propagation. We further calculate the overall GCR streaming induced magnetic field, and find a qualitative consistence with the observed structure of the halo magnetic field. ", "introduction": "Galactic magnetic field (GMF), Galactic cosmic ray (GCR) and the ordinary matter are the basic components of the interstellar medium (ISM). These three constituents have comparable pressure and are bounded together by the electromagnetic force. The GMF and the propagation of GCRs help to support the ordinary matter against the self-gravity. Conversely, the weight of the ordinary matter confines GMF and then GCRs in the Galaxy \\citep{fer01}. The dynamic balance process between these three constituents could give birth to new molecular-cloud complexes and ultimately trigger star formation \\citep{mou74,elm82}. The physical nature of and the interactions between them are studied for decades. However, the fundamental question, the origin of them, is still open. Nowadays, due to the development of the observation technology and methods, the information about these constituents is more adequate, and we should approach a better understanding of this question. The GMF is commonly believed to be produced by a dynamo process with a seed field (e.g., see review \\citep{bra05} and references therein). Based on the measurements of the Faraday rotation of the linearly polarized radiation from pulsars and extragalactic radio sources \\citep{sim81,con98,bro03,han09,tay09,law11}, two kinds of structures are found in the halo magnetic field: the poloidal fields of dipole structure perpendicular to the plane in the Galactic center (GC) region and the toroidal fields of opposite directions above and below the Galactic plane \\citep{han97,han99}. The \\textquotedblleft A0\\textquotedblright dynamo model \\citep{han97} could produce such a large scale magnetic field. Nevertheless, the origin of the GMF is still an open question. The GCR streaming, i.e. GCR propagation, could have contributions to the GMF. \\citet{par92,hana09} suggested that the GMF could be generated by the dynamo process driven by GCR streaming. \\citet{dol04} proposed a mechanism to produce the large-scale GMF by the electric current induced by GCRs. In their model, the current was theoretically calculated based on the diffusion law, and the result depends strongly on the assumption of the diffusion coefficient tensor. However, there is no direct measurement of the diffusion coefficient. It is better to use a direct method to estimate the GCR streaming. The anisotropy of GCRs should be a precise tool to probe the GCR streaming. Early in 1935, \\citet{com35} suggested that the relative movement between the GCR plasma and the observer due to the Galactic rotation could lead to a dipole anisotropy, known as the Compton-Getting (CG) effect. According to this effect, the direction and velocity of the CR streaming can be estimated by the phase and amplitude of the anisotropy respectively. \\citet{ame06} did precise measurement above 300 TeV and concluded that the GCR plasma might corotate with the solar environment around the Galactic center with a velocity ${\\sim}220 \\ km s^{-1}$. In such a manner, the CR streaming in the rest frame of the Galaxy could be deduced from the high-precision two-dimensional anisotropy map, which is observed by many experiments in a wide energy range. In this work, we attempt to extend the anisotropy pattern observed in the solar vicinity to the whole Galaxy. The GCR streaming related to the anisotropy should exist in Galactic scale, forming a new picture of the GCR propagation. Furthermore, the global streaming of the GCR can be used to explore the contribution of GCR to the GMF. This Letter is organized as follows: In Section 2, we introduce the observational anisotropy and the GCR streaming in the extended picture. In Section 3, we estimate the contribution of the GCR to the GMF in the halo. In Section 4, we present discussions and the conclusions of this extension. ", "conclusions": "The observational anisotropy has a stable large-scale structure in the energy range from tens of GeV to hundreds of TeV. These structures imply three directions of the GCR streaming: inward and outward along the tangential direction of the spiral arm and the perpendicular direction to the Galactic pole. In this work, we extend the anisotropy observed locally to the whole Galaxy, which means the GCR streamings might exist at the global scale. A loop of the GCR streaming gives a new signature of the propagation of GCRs, and also provides contributions to the GMF. Qualitatively, the GCR streamings inferred from the anisotropy of the GCRs can generate a large-scale magnetic field with poloidal and toroidal structures. These structures are consistent with the observations. This indicates that the extension of the local anisotropy to the whole Galaxy is probably reasonable. It will be helpful for further understanding of the origin of GCR anisotropy. Moreover, the anisotropy of the GCRs provides a new measure to study the Galactic electric current as well as a new window to understand the origin of the GMF. This analysis also indicates that the magnetic field in the Galactic halo may be partially contributed by the electric current induced by GCRs. The extension of the local anisotropy to the whole Galaxy is too simple in this work and the uncertainties introduced by the parameters used in the calculation are quite large, so the quantitative result only can be regarded as an order-of-magnitude estimate." }, "1101/1101.5759_arXiv.txt": { "abstract": "{On 17 January 2005 two fast coronal mass ejections were recorded in close succession during two distinct episodes of a 3B/X3.8 flare. Both were accompanied by metre-to-kilometre type-III groups tracing energetic electrons that escape into the interplanetary space and by decametre-to-hectometre type-II bursts attributed to CME-driven shock waves. {\\bl A peculiar type-III burst group was observed below 600 kHz 1.5 hours after the second type III group. It occurred without any simultaneous activity at % higher frequencies, around the time when the two CMEs were expected to interact.} We associate this emission with the interaction of the CMEs at heliocentric distances of about 25~\\RSUN. Near-relativistic electrons observed by the EPAM experiment onboard ACE near 1~AU revealed successive particle releases that can be associated with the two flare/CME events and the low-frequency type-III burst at the time of CME interaction. We compare the pros and cons of shock acceleration and acceleration in the course of magnetic reconnection for the escaping electron beams revealed by the type III bursts and for the electrons measured {\\it in situ}.} ", "introduction": "The acceleration of charged particles to high energies in the solar corona is related to flares, which reveal the dissipation of magnetically stored energy in complex magnetic field structures of the low corona, and to coronal mass ejections (CMEs), which are large-scale, complex magnetic-field-plasma structures ejected from the Sun. CMEs can drive bow shocks, and their perturbation of the coronal magnetic field can also give rise to magnetic reconnection, where energy can be released in a similar way as during flares. When several CMEs are launched along the same path, a faster CME may overtake a slower preceding one, and the two CMEs can merge into a single structure. For this phenomenon \\citet{Gopalswamy01} introduced the term {\\em{CME Cannibalism}}. The CME-CME interaction was found associated with a characteristic low-frequency continuum radio emission. \\citet{Gopalswamy02} interpreted this type of activity as the radio signature of non-thermal electrons originating either during reconnection between the two CMEs or as the shock of the second, faster CME travels through the body of the first \\citep[see ][for a numerical study of two interacting coronal mass ejections]{Schmidt04}. In this paper we use radio diagnostics to study electron acceleration during a complex solar event broadly consisting of two stages, {\\bl each associated with a distinct episode of a flare} and with a fast CME, which occurred in close temporal succession on 17 January 2005. The CMEs interacted at a few tens of {\\RSUN}. Both the flare/CME events and the CME interaction were accompanied by radio emission, which is used here to study electron acceleration scenarios. Energetic electrons in the corona and interplanetary space are traced by their dm-to-km-wave radio emission, mostly excited at or near the electron plasma frequency. The emission provides a diagnostic of the type of the exciter and its path from the low corona (cm-dm wavelengths) to 1~AU (km wavelengths). Radio emissions from exciters moving through the corona appear in dynamic spectra as structures exhibiting a drift in the time--frequency domain. The drift rate depends on their speed and path, resulting in a variety of bursts. Type~III bursts trace the path of supra--thermal electrons guided by magnetic structures. They appear, on dynamic spectra, as short (lasting from a fraction of a second at dm-waves to a few tens of minutes at km-waves) structures with fast negative drift, \\citep[$\\frac{1}{f} \\frac{df}{dt} \\approx 0.5 \\rm \\; sec^{-1}$;~see for example~][]{GuedelBenz88}. This corresponds to anti-sunward propagation of the electrons through regions of decreasing ambient density at speeds $\\approx c/3$ \\citep[e.g.,][]{Suz:Dul-85}. Sunward travelling beams produce reverse drift bursts (RS bursts), and beams propagating in closed loops emit type U or J bursts comprising a succession of an initial drift towards lower frequencies and a more or less pronounced RS burst. Type~II bursts are more slowly drifting bursts \\citep[$\\frac{1}{f} \\frac{df}{dt} \\approx 0.001-0.01 \\rm \\; sec^{-1}$;~see, for example, Table A.1 in~][]{Caroubalos04} excited by electrons accelerated at travelling shocks and emitting in their upstream region. Finally broadband dm-m wave continuum emission that may last over several minutes or even hours (type IV burst) is ascribed to electrons confined in closed coronal magnetic structures. The reader is referred to the reviews in \\cite{McLean85}, \\cite{Bas:al-98}, \\cite{Nindos08} and \\cite{Pick08} for more detailed accounts of the radio emission by non thermal electrons in the corona. \\begin{table}[t] \\centering \\caption{Overview of the 17 January 2005 Event and associated activity. } \\label{T} \\begin{tabular}{{lllll}} % \\hline\\hline \\textbf{Event}\t\t&\\textbf{Time} \t&\\textbf{Characteristics} & \\textbf{Remarks}\t \\\\ &\\textbf{UT} \t& \t\t \t\t\t&\t\t\t \\\\ \\hline SXR Start \t\t\t& \t06:59 \t\t& \t\t\t\t& AR10720 (N15$^\\circ$ W25$^\\circ$) \\\\ Type IV \t\t\t& \t08:40 \t\t& 3.0-630 MHz\t\t\t\t& AR10720 \t \t \\\\ CME$_1$ \t\t\t& \t09:00\t\t\t& \t\t\t& lift-off \t \\\\ \\hline \\textbf{SXR Stage 1} \t\t&\t09:05 \t\t& \t\t\t\t& \t\t\t \\\\ { First cm }\t\t\t&\t09:05 \t\t& \t\t\t\t& RSTN 15400 MHz \\\\ { burst start}\t\t\t&\t\t\t\t&\t\t\t\t\t&\t\t\t\\\\ Type III$_1$\t\t\t& 09:07-09:28 \t\t\t& 0.2-630 MHz \t \t\t\t& AR10720 \t \t \\\\ Type II$_1$ \t\t\t&\t09:11 \t\t& 0.2-5 MHz \t\t\t\t& AR10720 \t \t \\\\ \\Ha~Start \t\t\t&\t09:13 \t\t& 3B \t \t\t\t& KANZ, AR10720 \t \t \\\\ CME$_1$ \t\t\t& \t09:30 \t\t& 2094 km sec$^{-1}$ \t\t\t& On C2\t\t\t \\\\ HXR Start \t\t\t& 09:35:36\t\t\t&\t\t\t\t\t& RHESSI Number 5011710\t \\\\ CME$_2$ \t\t\t& \t09:38 \t\t& \t\t\t& lift-off \t \\\\ \\hline \\textbf{SXR Stage 2} \t\t&\t09:42\t\t\t& \t\t\t\t& End SXR Stage 1 \t \\\\ { Second cm }\t\t\t&\t09:43 \t\t& \t\t\t\t& RSTN 15400 MHz \\\\ { burst start}\t\t\t&\t\t\t\t&\t\t\t\t\t&\t\t\t\\\\ Type III$_2$\t\t\t& 09:43-09:59 \t\t\t& 0.2-630 \t \t\t\t& AR10720 \\\\ HXR peak \t\t\t& 09:49:42\t\t\t& 7865 counts sec$^{-1}$\t\t& \t\t \\\\ Type II$_2$ \t\t\t&\t09:48 \t\t& 0.2-8 MHz \t \t\t\t& AR10720 \\\\ SXR peak \t\t\t&\t09:52 \t\t& X3.8 \t\t\t\t& End SXR Stage 2 \t \\\\ \\hline CME$_2$ \t\t\t& \t09:54 \t\t& 2547 km sec$^{-1}$ \t\t& On C2\t\t\t \\\\ First rise\t\t\t&\t10:00\t\t\t& 38-315 keV\t\t\t\t& ACE/EPAM\t\t\\\\ Electron flux\t\t\t&\t\t\t\t&\t\t\t\t\t&\t\t\t\\\\ SXR End \t\t\t&\t10:07 \t\t& \t\t\t\t& AR720 \t \t \\\\ HXR End \t\t\t& 10:38:52\t\t\t& 53152112 total counts\t\t\t& RHESSI\t\t \\\\ Second rise\t\t\t& 12:00 \t\t\t& 38-315 keV\t\t\t\t& ACE/EPAM\t\t\\\\ Electron flux\t\t\t&\t\t\t\t&\t\t\t\t\t&\t\t\t\\\\ Type III$_3$\t\t\t&\t11:37\t\t\t& 0.5 MHz \t \t\t\t& CME$_1$, CME$_2$ merge at 37 \\RSUN \\\\ &\t\t\t\t& \t\t \t\t\t& type II$_2$ overtakes type II$_1$ \\\\ \\Ha~End\t \t\t\t&\t11:57 \t\t&\t \t\t\t& KANZ \t\t \\\\ Type IV End \t\t\t&\t15:24 \t\t& 3.0-630 MHz\t \t\t\t& AR10720 \\\\ \\hline\\hline \\end{tabular} \\end{table} \\begin{figure}[h!]\\centering \\includegraphics[scale=0.70]{newKANZ_20050117_panel_bw_2.eps} \\caption{Snapshots of active region NOAA 10720 on 17 January 2005 in \\Ha ~line centre (top left) and in the wing, observed at Kanzelh\\\"ohe Observatory (courtesy M.~Temmer). Solar north is at the top, west on the right. The two snapshots at the top show the active region before the flare under discussion, the two bottom images show two instants during the stages 1 and 2, respectively. These stages were associated with the disappearance of the filaments labelled `F1' and `F2'.} \\label{Fig_KANZ} \\end{figure} \\begin{figure}[h!]\\centering \\includegraphics[scale=0.95]{CME_Frames.eps} \\caption{The two LASCO CMEs in close succession; the images have been subjected to high-pass filtering. Top: Two frames of the 09:30:05 Halo CME with back-extrapolated lift off at 09:00:47 UT and plane-of-the-sky speed 2094 km sec$^{-1}$. Bottom: Two frames of the 09:54:05 Halo CME with back-extrapolated lift off at 09:38:25 UT and plane-of-the-sky speed 2547 km sec$^{-1}$. Solar north is at the top, west on the right.} \\label{CMES} \\end{figure} ", "conclusions": "" }, "1101/1101.3725_arXiv.txt": { "abstract": "{} {In spite of large overabundances of \\ion{Xe}{ii} observed in numerous mercury-manganese (HgMn) stars, \\ion{Xe}{ii} oscillator strengths are only available for a very limited number of transitions. As a consequence, several unidentified lines in the spectra of HgMn stars could be due to \\ion{Xe}{ii}. In addition, some predicted \\ion{Xe}{ii} lines are redshifted by about 0.1\\,\\AA\\ from stellar unidentified lines, raising the question about the wavelength accuracy of the \\ion{Xe}{ii} line data available in the literature. For these reasons we investigated the \\ion{Xe}{ii} lines lying in the 3900-4521\\,\\AA, 4769-7542\\,\\AA, and 7660-8000\\,\\AA\\ spectral ranges of four well-studied HgMn stars. } {We compared the \\ion{Xe}{ii} wavelengths listed in the NIST database with the position of the lines observed in the high-resolution UVES spectrum of the xenon-overabundant, slowly rotating HgMn stars HR\\,6000, and we modified them when needed. We derived astrophysical oscillator strengths for all the \\ion{Xe}{ii} observed lines and compared them with the literature values, when available. We checked the stellar atomic data derived from HR\\,6000 by using them to compute synthetic spectra for three other xenon-overabundant, slowly rotating HgMn stars, HD\\,71066, 46\\,Aql, and HD\\,175640. In this framework, we performed a complete abundance analysis of HD\\,71066, while we relied on our previous works for the other stars. } {We find that all the lines with wavelengths related to the 6d and 7s energy levels have a corresponding unidentified spectral line, blueshifted by the same quantity of about 0.1\\,\\AA\\ in all the four stars, so that we identified these lines as coming from \\ion{Xe}{ii} and modified their NIST wavelength value according to the observed stellar value. We find that the \\ion{Xe}{ii} stellar oscillator strengths may differ from one star to another from 0.0\\,dex to 0.3\\,dex. We adopted the average of the oscillator strengths derived from the four stars as final astrophysical oscillator strength. } {} ", "introduction": "Several studies of mercury-manganese (HgMn) stars have pointed out the presence of xenon with overabundances up to 5\\,dex relative to the solar value $\\log$(N$_{Xe}$/N$_{tot}$)=$-$9.87 (Grevesse \\& Sauval 1998). This is, for instance, the case of $\\kappa$\\,Cnc and 33\\,Gem, for which abundances equal to $-$4.87$\\pm$0.13\\,dex and $-$4.90$\\pm$0.07\\,dex were determined by Dworetsky et al. (2008). The xenon overabundance implies the presence of numerous \\ion{Xe}{ii} lines in the spectra of the HgMn stars, but the \\ion{Xe}{ii} transition probabilities are very incomplete, when we compare the large number of transitions listed in the NIST database and the small number of them with an associated $\\log\\,gf$-value. As a consequence, the computed spectra do not include numerous \\ion{Xe}{ii} lines, raising the doubt that some unidentified lines could just be due to \\ion{Xe}{ii}. In addition, we noticed that the wavelengths of several \\ion{Xe}{ii} lines are close, but not coincident with the wavelength of some unidentified stellar lines (Castelli \\& Hubrig 2007), so that we wondered about the accuracy of the wavelength determination from laboratory spectra. The most complete work on \\ion{Xe}{ii} is that of Hansen \\& Persson (1987), who analyzed all the published (Boyce 1936; Humphreys 1939) and unpublished \\ion{Xe}{ii} lines from 392\\,\\AA\\ to 10220\\,\\AA\\ obtained in laboratory by Humphreys and Boyce. In their discussion on the wavelength accuracy, Hansen \\& Persson (1987) pointed out that the wavelength accuracy for many lines is too low to be satisfactory, mostly owing to the widely varying quality of the experimental data they used. They announced new experimental work to improve the \\ion{Xe}{ii} atomic data. Unfortunately, this work has never been published up to now, all the more so that some preliminary results had indicated that, for the high 6d and 7s levels, there were shifts of about 0.5\\,cm$^{-1}$ between the energy levels determined from the Humphrey wavelengths and the energy levels determined from the new data. This energy difference corresponds to a difference of 0.1\\,\\AA\\ in wavelengths. Saloman (2004), who performed a critical compilation of all the work on energy levels and wavelengths of \\ion{Xe}{ii} made up to that time, adopted the data from Hansen \\& Persson (1987) for almost all the lines of the optical region. The Saloman (2004) critical compilation is the one adopted by the NIST database. To study wavelengths and $\\log\\,gf$-values of the \\ion{Xe}{ii} lines having intensities $\\ge$ 100 in the NIST line list, we used UVES spectra of the four xenon overabundant HgMn stars HR\\,6000, HD\\,71066, 46\\,Aql, and HD\\,175640. They are slowly rotating stars with v{\\it sini} 1.5\\,km\\,s$^{-1}$, 1.5\\,km\\,s$^{-1}$, 1.0\\,km\\,s$^{-1}$, and 2.5\\,km\\,s$^{-1}$, respectively. We already performed a complete abundance analysis for HD\\,175640 (Castelli \\& Hubrig 2004a)\\footnote{http://wwwuser.oat.ts.astro.it/castelli/hd175640/hd175640.html} and for HR\\,6000 and 46\\,Aql (Castelli et al. 2009)\\footnote{http://wwwuser.oat.ts.astro.it/castelli/hr6000new/hr6000.html}. To be consistent with the other papers, we present here an abundance analysis of HD\\,71066, which was studied with the same methods as adopted for the other stars. A previous work on HD\\,71066, related to vertical abundance stratification in HgMn stars, was performed by Thiam et al. (2010), who adopted the same observations as are used in this paper. We note, however, that no mention about \\ion{Xe}{ii} was made in their study. \\section {Observations and data reduction} All the stars were observed at the European Southern Observatory (ESO) using the Very Large Telescope Ultraviolet and Visible Echelle Spectrograph (UVES) with a resolving power ranging from 80000 to 110000. HD\\,175640 was observed on June 13, 2001 (Castelli \\& Hubrig 2004a). HR\\,6000, 46\\,Aql, and HD\\,71066 were part of the same observational run (ESO program 076.D-0169(A)). The spectra of HR\\,6000 were observed on September 19, 2005, those of 46 Aql on October 18, 2005 (Castelli et al., 2009), while the spectra of HD\\,71066 were taken on October 27, 2005. Because Nunez et al. (2010) found spectral variations in 19 HgMn stars out of a sample of 28 HgMn stars analyzed, we investigate about a possible variability of HD\\,71066 by comparing the spectrum observed in 2005 with an UVES spectrum observed in April 2004. We did not find any clear indication of variability. The spectra of the four stars cover the region 3030 $-$ 10000\\,\\AA. For HD\\,175640 there are two gaps at $\\lambda\\lambda$ 5759 $-$ 5835\\,\\AA\\ and 8519 $-$ 8656\\,\\AA. For the other three stars, the gaps occur at 4520 $-$ 4769\\,\\AA\\ and 7536 $-$ 7660\\,\\AA. All the spectra were reduced by the UVES pipeline Data Reduction Software (Ballester et al. 2000). We analyzed flux-calibrated spectra for the 3050-5750\\,\\AA\\ region and RED$_{-}$SCI$_{-}$POINT spectra for the 5750-9460\\,\\AA\\ interval, in that flux-calibrated reduction for the red spectra was not implemented in the pipeline reduction procedure. The measurement procedures on the spectra of HD\\,175640, HR\\,6000, and 46\\,Aql were described in Castelli \\& Hubrig (2004a) and Castelli \\& Hubrig (2007). The spectra of HD\\,71066 were normalized to the continuum using the IRAF continuum task. The equivalent widths were measured by a Gaussian fitting using the IRAF splot task. The S/N ratio is different for the different stars. In the spectra of HD\\,175640, it ranges from 200 in the near UV to 400 in the visual region. It is higher than the S/N of the spectra of the other three stars, which were observed in a different epoch. Furthermore, for each star, it is different in the different spectral intervals. For instance, for HR\\,6000, it is about 100 in the 5800$-$6800\\,\\AA\\ interval and lowers to about 25 at 7400\\,\\AA\\ (REDL spectrum). It is about 50 at 7800\\,\\AA\\ and decreases to about 25 at 9400\\,\\AA\\ (REDU spectrum). This behavior is similar for 46 Aql and HD\\,71066. At 6800\\,\\AA\\ the S/N is about 100 for HR\\,6000, 70 for 46 Aql, 100 for HD\\,71066, and 125 for HD\\,175640. \\begin{table}[] \\begin{center} \\caption{Abundances $\\log$(N$_{elem}$/N$_{tot}$) for HD\\,71066.} \\begin{tabular}{lccrccccccccccccccc} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{elem}& \\multicolumn{1}{c}{HD\\,71066}& \\multicolumn{1}{c}{Star-Sun}& \\multicolumn{1}{c}{Sun$^{a}$}& \\multicolumn{1}{c}{Thiam et al.(2010)} \\\\ & [12000K,4.1]& & & [12010,3.95]\\\\ \\hline\\noalign{\\smallskip} \\ion{He}{i} &$\\le$ $-$2.28 &$\\le$ [$-$1.23] & $-$1.05 & $-$2.30$\\pm$0.40\\\\ \\ion{Be}{ii} &$-$10.79& [$-$0.15] & $-$10.64\\\\ \\ion{C}{ii} & $-$3.90 &[$-$0.38] & $-$3.52&$-$3.89$\\pm$0.10\\\\ \\ion{N}{i} & $\\le$ $-$5.50& $\\le$$-$1.38 & $-$4.12\\\\ \\ion{O}{i} & $-$3.61$\\pm$0.05 & [$-$0.40] & $-$3.21 & $-$3.61$\\pm$0.14\\\\ \\ion{Ne}{i}&$\\le$ $-$4.70 & $\\le$[$-$0.74] & $-$3.96\\\\ \\ion{Na}{i} & $-$5.51 $\\pm$0.08 & [$+$0.20] & $-$5.71\\\\ \\ion{Mg}{i} & $-$5.32 $\\pm$0.05 & [$-$0.86]& $-$4.46\\\\ \\ion{Mg}{ii} & $-$5.40 & [$-$0.94] & $-$4.46 &$-$5.46$\\pm$0.01\\\\ \\ion{Al}{i} & $\\le$$-$7.30 &$\\le$[$-$1.73] & $-$5.57\\\\ \\ion{Al}{ii} &$\\le$$-$7.30 &$\\le$[$-$1.73] & $-$5.57\\\\ \\ion{Si}{ii} & $-$4.61$\\pm$0.19 & [$-$0.12] & $-$4.49&$-$4.58$\\pm$0.07\\\\ \\ion{P}{ii} & $-$5.06$\\pm$0.13 & [$+$1.53] & $-$6.59&$-$4.87$\\pm$0.22\\\\ \\ion{P}{iii}& $-$5.13 & [$+$1.46] &$-$6.59\\\\ \\ion{S}{ii} &$-$5.77$\\pm$0.11 & [$-$1.06] & $-$4.71&$-$5.66$\\pm$0.20\\\\ \\ion{Ca}{ii} & $-$6.50$\\pm$0.21: & [$-$0.82] & $-$5.68&$-$6.02\\\\ \\ion{Sc}{ii} & $\\le$$-$$$10.50 &$\\le$[$-$1.63] & $-$8.87\\\\ \\ion{Ti}{ii} & $-$6.45$\\pm$0.06 &[$+$0.57] & $-$7.02 &$-$6.52$\\pm$0.05\\\\ \\ion{V}{ii} &$\\le$$-$10.0 & $\\le$[$-$1.96] & $-$8.04\\\\ \\ion{Cr}{ii} & $-$6.17$\\pm$0.06 &[$+$0.20] & $-$6.37 &$-$6.28$\\pm$0.09\\\\ \\ion{Mn}{ii} & $-$5.95$\\pm$0.04 &[$+$0.70] & $-$6.65 & $-$5.81$\\pm$0.20\\\\ \\ion{Fe}{i} & $-$3.85$\\pm$0.06 & [$+$0.69] & $-$4.54&$-$3.98$\\pm$0.06\\\\ \\ion{Fe}{ii} & $-$3.85$\\pm$0.13 &[$+$0.69] & $-$4.54&$-$3.87$\\pm$0.14\\\\ \\ion{Co}{ii} &$\\le$$-$7.88 & $\\le$[$-$0.76] & $-$7.12\\\\ \\ion{Ni}{ii} &$\\le$$-$7.90 & $\\le$[$-$2.11] & $-$5.79\\\\ \\ion{Cu}{ii} &$\\le$$-$7.83 &$\\le$[0.00] & $-$7.83\\\\ \\ion{Zn}{ii} & $\\le$$-$7.94&$\\le$[$-$0.5] & $-$7.44\\\\ \\ion{As}{ii} &$-$6.3: &$+$3.37: & $-$9.67 \\\\ \\ion{Sr}{ii} & $-$8.27 & [$+$0.8] & $-$9.07 &$-$8.35\\\\ \\ion{Y}{ii} &$-$7.57$\\pm$0.08 & [$+$2.23] & $-$9.80\\\\ \\ion{Xe}{ii} &$-$5.43$\\pm$0.16 & [$+$4.44] & $-$9.87\\\\ \\ion{Nd}{iii}&$-$9.63$\\pm$0.01&[$+$0.91] &$-$10.54\\\\ \\ion{Dy}{iii}&$-$9.90&[$+$1.00] &$-$10.90\\\\ \\ion{Au}{ii} &$-$7.12$\\pm$0.03&[$+$3.91] &$-$11.03\\\\ \\ion{Hg}{i} & $-$6.40 &[$+$4.51] &$-$10.91& $-$6.38$\\pm$0.28\\\\ \\ion{Hg}{ii} & $-$6.40 & [$+$4.51] &$-$10.91&$-$6.53$\\pm$0.33\\\\ \\hline\\noalign{\\smallskip} \\end{tabular} \\end{center} $^{a}$ Solar abundances are from Grevesse \\& Sauval (1998).\\\\ \\end{table} \\begin{table*}[] \\begin{center} \\caption{The strongest emission lines in HD\\,71066, with the atomic data and configurations from the Kurucz website (see footnote\\,5) } \\begin{tabular}{lccrccccccccccccccc} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{$\\lambda$($\\AA$)}& \\multicolumn{1}{c}{elem}& \\multicolumn{1}{c}{$\\log\\,gf$}& \\multicolumn{1}{c}{$\\chi_{low}$}& \\multicolumn{1}{c}{J$_{low}$}& \\multicolumn{1}{c}{lower config.}& \\multicolumn{1}{c}{$\\chi_{up}$}& \\multicolumn{1}{c}{J$_{up}$}& \\multicolumn{1}{c}{upper config.}& \\multicolumn{1}{c}{Rc obs.}& \\multicolumn{1}{c}{Rc comp.} \\\\ \\hline\\noalign{\\smallskip} 5987.384 &\\ion{Ti}{ii}& $+$0.649 & 64979.278& 3.5 &($^{3}$F)4d e4G & 81676.439 & 4.5& ($^{3}$F)4f 2[4]& 1.012 & 0.983\\\\ 6001.400 &\\ion{Ti}{ii}& $+$0.724 & 65095.972& 4.5 &($^{3}$F)4d e4G & 81754.137 & 5.5& ($^{3}$F)4f 2[5]& 1.012 & 0.981\\\\ 6029.278 &\\ion{Ti}{ii}& $+$0.653 & 65308.434& 4.5 &($^{3}$F)4d e4H & 81889.576 & 5.5& ($^{3}$F)4f 3[6]& 1.025 & 0.984\\\\ 6125.861 &\\ion{Mn}{ii}& $+$0.788 & 82144.480& 3.0 &($^{6}$S)4d e5D & 98464.200 & 4.0& ($^{6}$S)4f $^{5}$F & 1.023 & 0.896\\\\ 6181.354 &\\ion{Cr}{ii}& $+$0.184 & 89812.420& 2.5 &($^{5}$D)4d f4D &105985.630 & 3.5& ($^{5}$D)4f 4[4]& 1.010 & 0.996\\\\ 6182.340 &\\ion{Cr}{ii}& $+$0.402 & 89336.890& 2.5 &($^{5}$D)4d e4P &105507.520 & 3.5& ($^{5}$D)4f 2[3]& 1.015 & 0.992\\\\ 6285.601 &\\ion{Cr}{ii}& $-$0.229 & 89885.080& 3.5 &($^{5}$D)4d f4D &105790.060 & 4.5& ($^{5}$D)4f $^{4}$F & 1.011 & 0.998\\\\ 6526.302 &\\ion{Cr}{ii}& $+$0.253 & 89885.080& 3.5 &($^{5}$D)4d f4D &105203.460 & 4.5& ($^{3}$F)sp r$^{4}$F & 1.010 & 0.996\\\\ 6551.373 &\\ion{Cr}{ii}& $+$0.201 & 90725.870& 3.5 &($^{5}$D)4d e4F &105985.630 & 3.5& ($^{5}$D)4f 4[4]& 1.018 & 0.997\\\\ 6585.241 &\\ion{Cr}{ii}& $+$0.815 & 90850.960& 4.5 &($^{5}$D)4d e4F &106032.240 & 5.5& ($^{5}$D)4f 4[6]& 1.028 & 0.987\\\\ 6592.341 &\\ion{Cr}{ii}& $+$0.287 & 90512.560& 1.5 &($^{5}$D)4d e4F &105677.490 & 2.5& ($^{5}$D)4f 3[3]& 1.014 & 0.996\\\\ 6636.427 &\\ion{Cr}{ii}& $+$0.573 & 90725.870& 3.5 &($^{5}$D)4d e4F &105790.060 & 4.5& ($^{5}$D)4f $^{4}$F & 1.020 & 0.992\\\\ 6961.439 &\\ion{Ti}{ii}& $+$0.663 & 67822.582& 4.5 &($^{3}$F)4d e2G & 82183.467 & 5.5& ($^{3}$F)4f 4[6]& 1.025 & 0.991\\\\ 6982.307 &\\ion{Ti}{ii}& $+$0.401 & 67606.162& 3.5 &($^{3}$F)4d e2G & 81924.126 & 4.5& ($^{3}$F)4f 3[4]& 1.015 & 0.995\\\\ 8335.148 &\\ion{C}{i} & $-$0.437 & 61981.820& 1.0 & p3s $^{1}$P & 73975.910 & 0.0& p3p $^{1}$S & 1.023 & 0.889\\\\ 9405.730 &\\ion{C}{i} & $+$0.285 & 61981.820& 1.0 & r3s $^{1}$P & 72610.720 & 2.0& p3p $^{1}$D & 1.088 & 0.730 \\\\ \\hline\\noalign{\\smallskip} \\end{tabular} \\end{center} \\end{table*} ", "conclusions": "From the high resolution stellar spectra of four HgMn stars we derived both wavelengths and $\\log\\,gf$-values for 100 \\ion{Xe}{ii} lines, which should also be observable in the spectra of numerous others chemically peculiar B-type stars. Of these lines, only 22 lines have $\\log\\,gf$-values available in the NIST database. The NIST wavelength of two of them, 4180.10\\,\\AA\\ and 4330.52\\,\\AA, differs by about 0.1\\,\\AA\\ from that observed in the spectra. There is a total of 27 lines in our sample for which the observed wavelength differs from the NIST wavelength by more than $-$0.06\\,\\AA\\ with the maximum shift of $-$0.13\\,\\AA\\ for the line at 4330.52\\,\\AA. We believe that the wavelength differences are mostly the result of uncorrect energy levels, in that they are all related to 6d or 7s levels, which have an uncertainty of about 0.5\\,cm$^{-1}$ according to Hansen \\& Persson (1987). This hypothesis seems us to be more relastic than that of some isotopic anomaly for \\ion{Xe}. For instance, using the isotopic wavelengths from Alvarez et al. (1979), Castelli \\& Hubrig (2007) excluded that the blueshift of 0.03\\,\\AA\\ observed for the \\ion{Xe}{ii} line at 6051.15\\,\\AA\\ can be due to some isotopic anomaly. Instead, because no isotopic composition was considered in our computations, owing to the lack of isotopic wavelengths for \\ion{Xe}{ii}, we could explain the larger astrophysical $\\log\\,gf$-value than the experimental one obtained for a few lines with the presence of the xenon isotopes, which should not be neglected in the computations of the strongest \\ion{Xe}{ii} line profiles. Good examples are the lines at 4844.33\\,\\AA, 5292.22\\,\\AA, and 5419.155\\,\\AA\\ (Table\\,6). On the basis of the wavelength shifts observed in the stellar spectra we redetermined the energy of three 7s, one 5d, and eighteen 6d levels. These levels, together with the old and new energy values, are listed in Table\\,7. We would like to point out that the new energy values depend, of course, on the accuracy of the energy of the lower level. The identification of the \\ion{Xe}{ii} lines and their consequent addition in the line lists, increases the accuracy of the synthetic spectra for the CP stars. In fact, it is important to be able to reproduce their high-resolution spectra well, because these stars are an excellent tool for extending laboratory spectrum analyses for several elements. An example is the determination of new high-excitation energy levels for \\ion{Fe}{ii} from the same UVES spectra of HR\\,6000 used for this paper (Castelli \\& Kurucz 2010). For instance, \\ion{As}{ii} is another element observed in some CP stars for which not even one $\\log\\,gf$-value in the optical region has been found in the literature. \\ion{As}{ii} has not only been observed in 46\\,Aql (Sadakane et al. 2001, Castelli et al. 2009), but also in HD\\,71066, as we have shown in this paper. If only one $\\log\\,gf$ value were given for it, we could derive astrophysical $\\log\\,gf$-values for the other lines, just as we did for \\ion{Xe}{ii}. The abundance analysis of HD\\,71066 has pointed out the overabundances of \\ion{Y}{ii}, \\ion{Nd}{iii}, \\ion{Dy}{iii}, and \\ion{Au}{ii} for the first time, in addition to the \\ion{Xe}{ii} and \\ion{As}{ii} overabundances. Those of other elements, in particular Hg, P, Ti, Cr, Mn, Fe, and Sr, have already been stated by Thiam et al. (2010) and confirmed by us. We found that HD\\,71066 is a typical HgMn star with Hg and Ca isotopic anomalies and emission lines for \\ion{C}{i}, \\ion{Ti}{ii}, \\ion{Cr}{ii}, and \\ion{Mn}{ii}. \\ion{He}{i} is underabundant and the shape of its profiles indicates the presence of helium vertical abundance stratification in the atmosphere. \\begin{table}[!hbp] \\caption[ ]{The final \\ion{Xe}{ii} astrophysical line list for the 3900-4525\\,\\AA\\ and 4780-8000\\,\\AA\\ intervals. The literature $\\log\\,gf$ sources are the NIST database, version 4 (NIST4) and Z\\'ieli\\'nska et al. (2002)(ZBD).} \\font\\grande=cmr7 \\grande \\begin{flushleft} \\begin{tabular}{lllllrrcllllll} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{$\\lambda$}& \\multicolumn{1}{c}{$\\log~gf$}& \\multicolumn{1}{c}{$\\log~gf$}& \\multicolumn{1}{c}{source}& \\multicolumn{1}{c}{$log\\gamma_{S}$}\\\\ \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{stellar}& \\multicolumn{1}{c}{stellar}& \\multicolumn{1}{c}{literature}& & &\\\\ \\hline\\noalign{\\smallskip} 3907.820 & $-$0.82$\\pm$0.06 & & &$-$4.684&\\\\ 4037.260 & $-$1.00$\\pm$0.00 & & \\\\ 4037.470 & $-$0.75$\\pm$0.00 & \\\\ 4057.360 & $-$0.80$\\pm$0.00 & & &$-$4.899 \\\\ 4157.980 & $-$0.60$\\pm$0.00 & & &$-$4.878 \\\\ 4162.160 & $-$1.57$\\pm$0.03 & & &$-$5.379 \\\\ 4180.007 & $-$0.35$\\pm$0.00 & $-$0.35 &NIST4& \\\\ 4193.100 & $-$0.60 \\\\ 4208.391 & $-$0.38$\\pm$0.02 \\\\ 4209.370 & $-$0.70$\\pm$0.00\\\\ 4213.620 & $-$0.22$\\pm$0.08&\\\\ 4215.620 & $-$1.05$\\pm$0.00&\\\\ 4222.900 & $+$0.64$\\pm$0.23& & &$-$4.778 \\\\ 4238.135 & $-$0.23$\\pm$0.10& & &$-$4.948 &\\\\ 4245.300 & $-$0.13$\\pm$0.07& & &$-$4.930 &\\\\ 4251.540 & $-$0.58$\\pm$0.02& & &$-$4.722 &\\\\ 4296.320 & $-$0.85$\\pm$0.00& & &$-$5.129 &\\\\ 4330.390 & $+$0.30$\\pm$0.00 & $+$0.498&NIST4&$-$4.884 &\\\\ 4369.100 & $-$0.72$\\pm$0.02& & &$-$4.890 & \\\\ 4373.700 & $-$0.70$\\pm$0.00 \\\\ 4384.910 &$\\le$ $-$1.95 & & &$-$5.358\\\\ 4393.090 & $+$0.00$\\pm$0.00& & &$-$4.927 &\\\\ 4395.770 & $+$0.00$\\pm$0.00& & &$-$4.884 &\\\\ 4414.840 & $-$0.50$\\pm$0.00& $+$0.243&NIST4&$-$5.432 \\\\ 4416.090 & $-$0.80& \\\\ 4448.025 & $+$0.10$\\pm$0.05\\\\ 4462.090 & $+$0.33$\\pm$0.00 & & &$-$4.866\\\\ $----$\\\\ 4787.77 &$-$0.82$\\pm$0.03 & & &$-$5.324\\\\ 4817.98 &$-$1.25$\\pm$0.00 & & &$-$5.351 \\\\ 4823.25 &$-$0.65$\\pm$0.00 & & &$-$4.989\\\\ 4844.33 &$+$0.61$\\pm$0.02& $+$0.491 &NIST4 & $-$5.347\\\\ & & $+$0.510$\\pm$0.027&ZBD\\\\ 4876.50 & $+$0.10$\\pm$0.00 & $+$0.255&NIST4 & $-$5.505 \\\\ 4883.53 & $-$0.25$\\pm$0.00 & & &$-$5.525 \\\\ 4884.09 & $-$0.80 \\\\ 4887.30 & $-$0.85$\\pm$0.05 & & &$-$5.423 &\\\\ 4890.085& $-$1.17$\\pm$0.04 &$-$0.754$\\pm$0.022&ZBD&$-$5.420 & \\\\ 4919.66 & $-$0.85$\\pm$0.12 & \\\\ 4921.48 & $+$0.05$\\pm$0.09& & &$-$4.442 & \\\\ 4971.68 & $-$0.75$\\pm$0.00 \\\\ 4972.70 & $-$0.55$\\pm$0.00 & & &$-$5.430 & \\\\ 4988.725& $-$0.85$\\pm$0.09 & & &$-$5.214 & \\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table} \\setcounter{table}{5} \\begin{table}[!hbp] \\caption[ ]{cont.} \\font\\grande=cmr7 \\grande \\begin{flushleft} \\begin{tabular}{lllllrrcllllll} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{$\\lambda$}& \\multicolumn{1}{c}{$\\log~gf$}& \\multicolumn{1}{c}{$\\log~gf$}& \\multicolumn{1}{c}{source}& \\multicolumn{1}{c}{$log\\gamma_{S}$}\\\\ \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{stellar}& \\multicolumn{1}{c}{stellar}& \\multicolumn{1}{c}{literature}& & &\\\\ \\hline\\noalign{\\smallskip} 5044.92 & $-$0.80$\\pm$0.00\\\\ 5080.51 & $-$0.22$\\pm$0.12\\\\ 5122.31 &$-$0.37$\\pm$0.09 & & &$-$4.951\\\\ 5188.08 &$-$1.10$\\pm$0.00 \\\\ 5260.42 &$-$0.37$\\pm$0.08 &$-$0.437 &NIST4&\\\\ 5261.95 &$+$0.25$\\pm$0.00 &$+$0.150&NIST4&$-$5.495 &\\\\ 5268.25 &$-$0.80$\\pm$0.12& & &$-$4.978 &\\\\ 5292.22 &$+$0.49$\\pm$0.06 &$+$0.351 & NIST4 &$-$5.482 &\\\\ & &$+$0.382$\\pm$0.013& ZBD\\\\ 5309.27 &$-$0.95$\\pm$0.00 &\\\\ 5313.76 &$-$0.09$\\pm$0.04 \\\\ 5339.355&$-$0.10$\\pm$0.03& $+$0.048$\\pm$0.019& ZBD \\\\ 5368.075&$-$1.05$\\pm$0.00 \\\\ 5372.405&$-$0.15$\\pm$0.06 & $-$0.211&NIST4&$-$5.551 \\\\ 5419.155&$+$0.37$\\pm$0.03 & $+$0.215&NIST4 &$-$5.481 \\\\ & & $+$0.256$\\pm$0.015& ZBD\\\\ 5438.96 &$-$0.44$\\pm$0.00 &$-$0.183& NIST4 &$-$5.544\\\\ 5450.45 &$-$0.97$\\pm$0.09 \\\\ 5460.365&$-$0.77$\\pm$0.04 &$-$0.673$\\pm$0.030&ZBD &$-$5.531&\\\\ 5472.60 &$-$0.55$\\pm$0.00 &$-$0.449 & NIST4&$-$5.482\\\\ & & $-$0.362$\\pm$0.030&ZBD\\\\ 5531.05 &$-$0.78$\\pm$0.10 & $-$0.616&NIST4 &$-$5.504\\\\ & & $-$0.632$\\pm$0.021 &ZBD\\\\ 5616.65 &$-$0.70$\\pm$0.17 &\\\\ 5659.38 &$-$0.65$\\pm$0.15 & & &$-$5.407 &\\\\ 5667.540& $-$0.53$\\pm$0.08 & & &$-$5.535\\\\ 5699.61 & $-$0.85& \\\\ 5719.587& $-$0.80$\\pm$0.00 &$-$0.746& NIST4&\\\\ & &$-$0.687$\\pm$0.023 & ZBD\\\\ 5726.88 & $-$0.28$\\pm$0.05\\\\ 5750.99 &$-$0.40$\\pm$0.05\\\\ 5758.665 &$-$0.35$\\pm$0.00 & & & $-$5.539 \\\\ 5776.39 &$-$0.70 & & & $-$5.488\\\\ 5893.29 &$-$0.90 \\\\ 5905.115&$-$0.75$\\pm$0.10 \\\\ 5945.53 &$-$0.67$\\pm$0.09 & & &$-$5.527\\\\ 5971.135 &$-$0.50 \\\\ 5976.460 &$-$0.29$\\pm$0.06 &$-$0.222 &NIST4 &$-$5.545\\\\ & & $-$0.317$\\pm$0.023 &ZBD\\\\ 6036.170&$-$0.56$\\pm$0.06 &$-$0.609 &NIST4 &$-$5.535 & \\\\ & &$-$0.562$\\pm$0.020 &ZBD\\\\ 6051.120&$-$0.28$\\pm$0.04 &$-$0.252& NIST4 &$-$5.515 & \\\\ & & $-$0.257$\\pm$0.020 & ZBD\\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table} \\setcounter{table}{5} \\begin{table}[!hbp] \\caption[ ]{cont.} \\font\\grande=cmr7 \\grande \\begin{flushleft} \\begin{tabular}{lllllrrcllllll} \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{$\\lambda$}& \\multicolumn{1}{c}{$\\log~gf$}& \\multicolumn{1}{c}{$\\log~gf$}& \\multicolumn{1}{c}{source}& \\multicolumn{1}{c}{$log\\gamma_{S}$}\\\\ \\hline\\noalign{\\smallskip} \\multicolumn{1}{c}{stellar}& \\multicolumn{1}{c}{stellar}& \\multicolumn{1}{c}{literature}& & &\\\\ \\hline\\noalign{\\smallskip} 6097.57 &$-$0.39$\\pm$0.06 &$-$0.237&NIST4 &\\\\ & & $-$0.355$\\pm$0.025 &ZBD\\\\ 6101.37 &$-$0.50$\\pm$0.28 &\\\\ 6194.07 &$+$0.05$\\pm$0.15 \\\\ 6270.81 &$-$0.18$\\pm$0.12 & $-$0.196& NIST4&$-$5.510\\\\ 6277.54 &$---$ & $-$0.894&NIST4& $-$5.543\\\\ & & $-$0.778$\\pm$0.021 & ZBD\\\\ 6300.830& $-$1.10 \\\\ 6343.95& $-$0.64$\\pm$0.10 &$-$0.786$\\pm$0.024& ZBD& \\\\ 6356.33& $-$0.25 \\\\ 6375.28 &$-$1.00 \\\\ 6512.79 &$-$1.00$\\pm$0.00 \\\\ 6528.65 &$-$0.40 \\\\ 6594.97 & 0.00$\\pm$0.00 \\\\ 6597.23 & $-$0.60$\\pm$0.00\\\\ 6620.02 &$-$0.85$\\pm$0.00& \\\\ 6694.285&$-$0.92$\\pm$0.12&$-$0.912$\\pm$0.020&ZBD \\\\ 6788.71&$-$0.50\\\\ 6790.37& $-$0.70\\\\ 6805.74&$---$&$-$0.595&NIST4\\\\ & &$-$0.547$\\pm$0.023& ZBD \\\\ 6990.835& $+$0.30$\\pm$0.05 & $+$0.200& NIST4\\\\ & & $+$0.084$\\pm$0.032 &ZBD\\\\ 7082.15 & $+$0.05 \\\\ 7164.85&$+$0.20$\\pm$0.00\\\\ 7284.34& $-$0.50 & \\\\ 7339.30 &$+$0.45? \\\\ 7787.04 &$-$0.50? \\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table} \\begin{table}[!hbp] \\caption[ ]{A few 7s, 5d, and 6d even Xe II energy levels from Hansen \\& Persson (1987) modified according to the wavelength positions observed in the UVES spectra of HR\\,6000, HD\\,71066, 46\\,Aql, and HD\\,175640} \\begin{flushleft} \\begin{tabular}{rrlllrrcllllll} \\hline\\noalign{\\smallskip} \\multicolumn{2}{c}{Term}& \\multicolumn{2}{c}{level value (cm$^{-1}$)}\\\\ & & NIST & This paper\\\\ \\hline\\noalign{\\smallskip} 5s$^{2}$5p$^{4}$($^{3}$P$_{2}$)7s & [2]$_{5/2}$ & 132518.82 & 132519.23 \\\\ & [2]$_{3/2}$ & 133189.42 & 133189.94 \\\\ 5s$^{2}$5p$^{4}$($^{3}$P$_{0}$)7s & [0]$_{1/2}$ & 140883.42 & 140883.79 \\\\ 5s$^{2}$5p$^{4}$($^{1}$D$_{2}$)5d & [0]$_{1/2}$ & 135060.97 & 135061.36 \\\\ 5s$^{2}$5p$^{4}$($^{3}$P$_{2}$)6d & [4]$_{9/2}$ & 136109.65 & 136110.13 \\\\ & [4]$_{7/2}$ & 136597.81 & 136598.48 \\\\ & [3]$_{7/2}$ & 135507.32 & 135507.72 \\\\ & [3]$_{5/2}$ & 139094.28 & 139094.83 \\\\ & [2]$_{5/2}$ & 135547.13 & 135547.53 \\\\ & [2]$_{3/2}$ & 135708.32 & 135708.72 \\\\ & [1]$_{3/2}$ & 139640.43 & 139640.61 \\\\ & [1]$_{1/2}$ & 136554.11 & 136554.47 \\\\ 5s$^{2}$5p$^{4}$($^{3}$P$_{1}$)6d & [3]$_{7/2}$ & 145587.61 & 145588.12 \\\\ & [3]$_{5/2}$ & 146927.86 & 146928.34 \\\\ & [2]$_{3/2}$ & 145940.34 & 145940.79 \\\\ & [1]$_{3/2}$ & 148085.19 & 148085.36 \\\\ & [1]$_{1/2}$ & 145222.72 & 145223.16 \\\\ 5s$^{2}$5p$^{4}$($^{3}$P$_{0}$)6d & [2]$_{5/2}$ & 144384.90 & 144385.45 \\\\ & [2]$_{3/2}$ & 144140.16 & 144140.69 \\\\ 5s$^{2}$5p$^{4}$($^{1}$D$_{2}$)6d & [4]$_{9/2}$ & 152806.73 & 152806.73 ? \\\\ & [4]$_{7/2}$ & 152708.92 & 152709.19 \\\\ & [1]$_{3/2}$ & 153584.09 & 153584.02 \\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table}" }, "1101/1101.4934_arXiv.txt": { "abstract": "We searched for binary companions to 20 young brown dwarfs in the Upper Scorpius association (145 pc, 5 Myr, nearest OB association) with the the Laser Guide Star adaptive optics system and the facility infrared camera NIRC2 on the 10 m Keck II telescope. We discovered a 0.14$\\arcsec$ companion (20.9$\\pm$0.4 AU) to the $<$0.1 M$_{\\odot}$ object SCH J16091837-20073523. From spectral deconvolution of integrated-light near-IR spectroscopy of SCH1609 using the SpeX spectrograph (Rayner et al. 2003), we estimate primary and secondary spectral types of M6$\\pm$0.5 and M7$\\pm$1.0, corresponding to masses of 79$\\pm$17 M$_{Jup}$ and 55$\\pm$25 M$_{Jup}$ at an age of 5 Myr and masses of 84$\\pm$15 M$_{Jup}$ and 60$\\pm$25 M$_{Jup}$ at an age of 10 Myr. For our survey objects with spectral types later than M8, we find an upper limit on the binary fraction of $<$9$\\%$ (1-$\\sigma$) at separations of 10 -- 500 AU. We combine the results of our survey with previous surveys of Upper Sco and similar young regions to set the strongest constraints to date on binary fraction for young substellar objects and very low mass stars. The binary fraction for low mass ($<$40 M$_{Jup}$) brown dwarfs in Upper Sco is similar to that for T dwarfs in the field; for higher mass brown dwarfs and very low mass stars, there is an excess of medium-separation (10-50 AU projected separation) young binaries with respect to the field. These medium separation binaries will likely survive to late ages. ", "introduction": "Numerous brown dwarf binaries have been discovered in the the field \\citep{clo03,bur03,bou03,bur06,liu06}. Almost all of these have projected separations of $<$15 AU, with the majority having separations of $<$7 AU. This tight binary distribution was initially viewed as evidence for the ejection scenario of brown dwarf formation \\citep{clo03}. In the ejection scenario, brown dwarfs are stellar embryos which are expelled from their natal subclusters due to interaction with other subcluster members, therefore cutting off accretion. Only tight brown dwarf binaries can survive an ejection event \\citep{rei01}. In the last decade a population of wide ($>$15 AU separation) very low mass star, brown dwarf, and ``planetary mass'' ($<$13 M$_{Jup}$) binaries have been discovered in young ($<$12 Myr) nearby clusters \\citep[][~see~Table~1~for~a~list~of~all~young~$\\leq$0.1~M$_{\\odot}$~binaries]{luh04, cha05, kra05, kra06, all06t, jay06, clo07, kon07, tod10,bej08}. These recent results suggest that the multiplicity properties of young ($\\sim$few Myr) substellar objects in star-forming regions may be substantially different from the old ($\\sim$few Gyr) field population. If common, these young binaries also provide serious constraints for current theories of brown dwarf formation, since such wide binaries cannot be formed by a non-dissipative ejection model \\citep{bat09}. However, most of these objects were either discovered serendipitously, are from surveys with unpublished statistics, or are from surveys with very few objects of comparable mass, so it is unknown how significant a population they form. Here, we conduct a systematic survey to search for such binaries in Upper Sco, the nearest OB association to the Earth. ", "conclusions": "We searched for binary companions to 20 brown dwarfs in Upper Scorpius (145 pc, 5 Myr, nearest OB association) with the laser guide star adaptive optics system and the facility infrared camera NIRC2 on the 10 m Keck II telescope. This survey is the most extensive to date for companions to very young (5 Myr), very low mass ($<$40 M$_{Jup}$) cluster brown dwarfs. We discovered a close companion (0.14$\\arcsec$, 20.9$\\pm$0.4 AU) to the very low mass object SCH J16091837-20073523. From spectral deconvolution of integrated-light near-IR spectroscopy of SCH1609-2007 using the SpeX spectrograph (Rayner et al. 2003), we estimate primary and secondary spectral types of M6$\\pm$0.5 and M7$\\pm$1.0, corresponding to masses of 79$\\pm$17 M$_{Jup}$ and 55$\\pm$25 M$_{Jup}$ at an age of 5 Myr and masses of 84$\\pm$15 M$_{Jup}$ and 60$\\pm$25 M$_{Jup}$ at an age of 10 Myr. For our survey objects with spectral types later than M8, we find an upper limit on binary fraction of $<$9$\\%$ (1-$\\sigma$) at separations greater than 10 AU. As expected from similar mass binaries in the field, we find that the binary fraction (10 -- 500 AU separations) appears to decrease monotonically with mass for young brown dwarfs. However, while proto-T-dwarfs (M$<$40 M$_{Jup}$) have a similar wide (10 -- 500 AU) binary fraction as field T dwarfs, there exists an anomalous population of wide higher mass binaries (0.07 -- 0.1 M$_{\\odot}$ primaries, separations of 10--50 AU) at young ages relative to older ages." }, "1101/1101.3513_arXiv.txt": { "abstract": "I present an analysis for fitting cosmological parameters from a Hubble Diagram of a standard candle with unknown intrinsic magnitude dispersion. The dispersion is determined from the data themselves, simultaneously with the cosmological parameters. This contrasts with the strategies used to date. The advantages of the presented analysis are that it is done in a single fit (it is not iterative), it provides a statistically founded and unbiased estimate of the intrinsic dispersion, and its cosmological-parameter uncertainties account for the intrinsic dispersion uncertainty. Applied to Type Ia supernovae, my strategy provides a statistical measure to test for sub-types and assess the significance of any magnitude corrections applied to the calibrated candle. Parameter bias and differences between likelihood distributions produced by the presented and currently-used fitters are negligibly small for existing and projected supernova data sets. ", "introduction": "The homogeneous nature of Type Ia supernovae (SNe Ia) makes them a popular tool for measuring cosmological distances. After empirical corrections based on light curve shape, color, and spectral features, the absolute magnitude (or distance modulus) of a supernova can be determined to $\\sim 0.12$ mag \\citep{2007A&A...466...11G, 2007ApJ...659..122J,2008ApJ...681..482C,2009A&A...500L..17B}. SNe Ia have been used to successfully measure the expansion rate of the universe \\citep[the Hubble Constant;][]{2001ApJ...553...47F,2009ApJ...699..539R}, discover its accelerated expansion \\citep{1998AJ....116.1009R,1999ApJ...517..565P}, and measure the properties of the dark energy responsible for that acceleration \\citep{2009ApJ...700.1097H,2010ApJ...716..712A}. The small scatter in the peak brightness of SN Ia luminosities is inferred from the small residuals in their Hubble Diagrams \\citep{1968AJ.....73.1021K}; the intrinsic supernova magnitude dispersion is measured from differences between observed magnitudes and those predicted by the cosmological model, e.g.\\ the linear Hubble law for low redshift. Although there are theoretical explanations for this dispersion including intrinsic progenitor properties, circumstellar dust, and viewing angle \\citep[see e.g.][]{2006ApJ...649..939K, 2008ApJ...686L.103G, 2003ApJ...591.1110W}, in practice the amount of dispersion is determined empirically from the data themselves. The luminosity dispersions of supernova subsets are statistics that can be used to compare and identify SN~Ia subclasses. The prevailing belief is that the intrinsic luminosity of an individual supernova, including line-of-sight effects, is encoded non-trivially within a finite set of physical and geometric parameters. The ``intrinsic'' dispersion arises from our lack of observational access to all those parameters and incomplete knowledge of how to exploit those that are available. It is possible that SN Ia subclasses with different average luminosities are responsible for some of the intrinsic dispersion seen in current data. Correlations between supernova light curves and spectral features \\citep{2005ApJ...623.1011B, 2009A&A...500L..17B, 2009ApJ...699L.139W, 2010arXiv1011.4517F} and host galaxy \\citep{2010MNRAS.406..782S,2010ApJ...722..566L} give evidence that SNe Ia need to be modeled in finer detail using an expanded suite of data. Likelihood surfaces of intrinsic dispersion for supernova subsets provide a statistical measure to test whether data are best described by a single intrinsic dispersion. This paper presents the methodology for simultaneously fitting for the intrinsic dispersion and the cosmological parameters that specify the dynamics of the cosmic expansion. Although I present straightforward textbook likelihood analysis, it has yet to be applied on supernova-cosmology data. My approach contrasts with that of \\citet{2010arXiv1006.2141S}, who suggest using Monte Carlo analysis of statistics that are insensitive to the intrinsic dispersion. This paper is organized as follows: \\S\\ref{likelihood:sec} presents the likelihood equation and contrasts it with the commonly used method. Results of simulations are given in \\S\\ref{simulation:sec} that show the quantitative differences between the results", "conclusions": "\\label{conclusions:sec} I have shown how to fit for cosmological parameters with SNe Ia when the intrinsic dispersion of the standard candle is unknown. My standard likelihood function has not been used in cosmological analysis to date. I show via simulation that, on average, our likelihood function is maximal at the values of the input parameters including the intrinsic dispersion. The presented and previously used iterative fitting methods do not give biases in the best-fit cosmological parameters and any differences in a single experiment are due to realization scatter. The fitter methods do return different intrinsic dispersions and parameter uncertainties. The procedure presented here has the advantage that it includes the covariance of the intrinsic dispersion with the other parameters in its error propagation, and the fit is done in a single iteration. The methodology can be extended to cases where multiple dispersion parameters are fit. I present an example taking the low- and high-redshift sets as being drawn from different magnitude distributions. The same approach can be used to check whether different supernova subsets (tagged for example by redshift, host-galaxy characteristics or spectral features) exhibit statistically significant differences in their population characteristics. The approach is appropriate for any analysis that uses a statistic for which the tracer has an intrinsic dispersion that must be determined from the data. For example, in weak gravitational lensing the measurement of correlated shear is obscured by the unknown intrinsic shape of individual galaxies. The intrinsic dispersion in galaxy ellipticities can be made a fit parameter determined simultaneously with those of cosmological interest. Inclusion of the likelihood-function normalization when fitting is not new to astronomy nor cosmology; \\citet{1995ApJ...438..322W} showed its importance in shot-noise-dominated photometry and it is retained in other cosmological analyses \\cite[see e.g.][]{2002MNRAS.335.1193B, 2010MNRAS.tmp.1232T}. \\citet{2010arXiv1009.5443H} do include a fit parameter in the data covariance for their supernova analysis although there it serves as a hyperparameter of the Gaussian-process prior on $w(z)$. \\citet{2010ApJ...717...40K} include the normalization term; though containing no fit parameters it is needed to directly compare the $\\chi^2$'s derived from different light-curve models. This paper gives a simplified view of how the standard candle nature of SNe Ia is used in cosmology analysis. SNe Ia are in fact calibrated candles; independent observables (light-curve shape, colors, spectral features) are correlated with peak absolute magnitude to correct and lower the dispersion in distance determinations. I advocate that intrinsic dispersion be measured as a fit parameter from the data simultaneously with the magnitude-correction and cosmological parameters. This provides a new perspective in how we search for magnitude corrections that make SNe Ia better calibrated candles. In the past we have sought parameterized magnitude corrections that minimize distance dispersion; we can now seek corrections and their inferred intrinsic dispersions that are most consistent with observations and are statistically favored over having no correction. Application of this technique to real SN data sets is the subject of ongoing work." }, "1101/1101.3349_arXiv.txt": { "abstract": "Self-annihilating or decaying dark matter in the Galactic halo might produce high energy neutrinos detectable with neutrino telescopes. We have conducted a search for such a signal using 276~days of data from the IceCube 22-string configuration detector acquired during 2007 and 2008. The effect of halo model choice in the extracted limit is reduced by performing a search that considers the outer halo region and not the Galactic Center. We constrain any large scale neutrino anisotropy and are able to set a limit on the dark matter self-annihilation cross section of $\\langle \\sigma_{A} v \\rangle \\simeq 10^{-22}{\\rm cm}^3 {\\rm s}^{-1}$ for WIMP masses above 1~TeV, assuming a monochromatic neutrino line spectrum. ", "introduction": "There is compelling observational evidence for the existence of dark matter. Although knowledge of its underlying nature remains elusive, a variety of theories provide candidate particles~\\cite{Bertone:2004pz}. Among those are Supersymmetry~\\cite{Martin:1997ns} and Universal Extra Dimensions~\\cite{Appelquist:2000nn}, both of which predict new physics at the electro-weak scale and, in most scenarios, introduce a light, and stable (or long lived) particle that exhibits the properties of a Weakly Interacting Massive Particle (WIMP)~\\cite{Steigman:1984ac}. WIMPs are an ideal dark matter candidate, predicted to have masses ranging from a few tens of GeV to several TeV. High energy neutrinos are expected to be produced as a result of the self-annihilation or decay of WIMPs. These neutrinos are detectable by high energy neutrino telescopes, making them powerful tools in the search for WIMPs and the investigation of their properties. In particular, they can be used to probe the self-annihilation cross section of dark matter candidates by looking for anomalous neutrino signals from the Galactic halo. Additionally, WIMPs could also be gravitationally captured by massive bodies like the Sun. If the annihilation rate of these captured WIMPs is regulated by the capture rate, then neutrino telescopes can be used to probe the WIMP-nucleon scattering cross section~\\cite{Abbasi:2009uz}. Recent observations of a GeV positron excess by PAMELA~\\cite{Adriani:2008zr}, an anomalous electron peak by ATIC~\\cite{ATIC:2008zzr}, and electron spectra from H.E.S.S.~\\cite{Aharonian:2009ah} and Fermi~\\cite{Abdo:2009zk}, demonstrate the importance of a multi-messenger approach to astrophysics and validate the interest in a neutrino channel. The observed lepton signals are inconsistent with each other or standard electron--positron production models~\\cite{Moskalenko:1997gh} and although they could potentially originate from nearby astrophysical sources (e.g. pulsars~\\cite{Yuksel:2008rf}), they could also be an indication of dark matter. If interpreted as the latter, it would suggest the existence of a leptophilic dark matter particle in the TeV mass range~\\cite{Meade:2009iu,Cirelli:2008pk}. Such a model would also result in significant high energy neutrino fluxes, through the decay of muons and $\\tau$-leptons. A significant fraction of neutrinos could also be produced directly as part of the annihilation~\\cite{Lindner:2010rr}, producing a line feature in the resulting neutrino spectrum. Such a mono-energetic neutrino flux is of specific interest since it can be used to set a model independent limit on the total dark matter self-annihilation cross section~\\cite{Beacom:2006tt} for the region of parameter space where gamma-ray signals would dominate. In this paper we discuss a search for neutrino signals produced by annihilating or decaying dark matter in the Galactic halo. The search is used to test the self-annihilation cross section by constraining the product of cross section and velocity averaged over the dark matter velocity distribution, $\\langle \\sigma_{A} v \\rangle$, and to probe the lifetime, $\\tau$. The search focuses on the outer Milky Way halo, where the dark matter density distributions are relatively well modelled. We do not include the Galactic Center region and thus remove any strong dependence on the choice of the halo profile. We quantify the residual weak dependence and present constraints on the dark matter self-annihilation cross section and lifetime in a model-independent way for a set of selected benchmark annihilation and decay channels, respectively. The paper is organized as follows: in the next section we describe the detector used for the data taken during 2007--2008 which is the base for our analysis. Section III discusses how we obtain an expected neutrino flux at Earth using different dark matter distributions and annihilation channels. In section IV we describe our data selection criteria and analysis strategy, which is followed by a discussion of the associated systematic uncertainties in section V. Section VI presents the result of the search, and section VII puts it in context with other experiments. Section VIII concludes by summarizing the results and giving an outlook for related searches. ", "conclusions": "} \\begin{ruledtabular} \\begin{tabular}{|l|r|} Effect & Sys. Uncertainty \\\\ \\hline Cosmic ray anisotropy & $0.2\\%$ \\\\ Exposure & $0.1\\%$ \\\\ \\hline Total Background & $0.3\\%$ \\\\ \\end{tabular} \\end{ruledtabular} \\end{table} The signal acceptance uncertainty is dominated by uncertainties in the ice properties and limitations in the detector simulation, which is uncorrelated with a number of theoretical uncertainties such as muon propagation, neutrino cross section, and bedrock uncertainty, each of which have been studied in previous analyses~\\cite{Abbasi:2009iv}. In addition, we consider the uncertainty due to Monte Carlo simulation statistics and detector exposure. The individual track pointing uncertainty (point spread function), on the order of one degree, is negligible in this analysis, which targets a large--scale anisotropy. Our dominant systematic uncertainty, the limited knowledge of ice properties as a function of depth and limitations in the detector simulation, is expected to produce an observed discrepancy between data and simulation for events near the horizon~\\cite{Abbasi:2009iv}. For nearly horizontal tracks the disagreement is maximal, with $30\\%$ more events observed in data compared to simulation predictions. Since we use the data itself to predict the number of background events in the on--source region, this discrepancy does not affect the background estimate. However, the signal acceptance can only be obtained from simulations. Hence, we must take this discrepancy into consideration for the signal acceptance uncertainty. The higher than expected observed data rate, when compared to simulation expectations, may indicate a contribution from mis-reconstructed down-going events, or a higher signal acceptance than expected. Both would cause the constraints presented later to be more conservative. The estimate for this systematic uncertainty in signal acceptance is 25-30\\%. The track reconstruction efficiency coupled with detector uptime (see Fig.~\\ref{fig:exposure}) results in a systematic uncertainty on the signal acceptance of $1\\%$. This uncertainty, combined with the theoretical uncertainties, results in a negligible contribution compared to the uncertainties in the optical properties of the ice. We therefore assume a 30\\% systematic signal acceptance uncertainty, primarily associated with that from the ice properties and limitations in the detector simulation. An additional systematic uncertainty to consider in signal acceptance is related to the photon detection efficiency of the DOMs, measured to be 8\\% in the laboratory~\\cite{Abbasi:2010vc}. The effect of this uncertainty on the passing rate of reconstructed tracks is found to range from about $1\\%$ for energetic events ($\\ge 1$~TeV), increasing to as much as 20\\% for lower energy events ($\\le 200$~GeV), as expected from annihilations assuming WIMP's of mass 200~GeV. We calculate this uncertainty for each of the considered WIMP masses and annihilation channels, then we add it in quadrature to the ice properties uncertainty discussed above. To derive the total uncertainty on the signal acceptance, we have added the systematic signal acceptance uncertainty in quadrature to the statistical uncertainty (Monte Carlo statistics). The Monte Carlo statistics uncertainty ranges from 3-6\\% (hard channels) and 4-16\\% (soft channels) in the TeV mass range dark matter, and increases to 50\\% (hard channels) and 90\\% (soft channel) at $m_{\\chi}=200$~GeV. The IceCube candidate neutrino sample, collected during 2007--2008 in the 22-string configuration, has been used to search for a neutrino anisotropy as expected from dark matter self annihilation in the Milky Way halo. Such an anisotropy was not observed and we have determined limits on the dark matter self-annihilation cross section $\\langle \\sigma_{A} v \\rangle$ at 90\\%~C.L. for WIMPs in the mass range from 200~GeV to 10~TeV. The IceCube detector sensitivity can be significantly improved by investigating the Galactic Center as a potential source. Such a search could be performed with the IceCube detector at a later construction stage and rely on selecting neutrinos interacting inside the detector volume. It would be able to significantly improve the constraints on the dark matter self-annihilation cross section given a particular choice of halo model in the case of a non observation. A large--scale anisotropy study as performed here, however, might provide a more distinct discovery signal. In the case of the Galactic Center, a dark matter signal would be more difficult to distinguish from other astrophysical neutrino sources, such as point sources (source contamination) or cosmic ray interaction with the interstellar medium." }, "1101/1101.4796_arXiv.txt": { "abstract": "Cosmic magnification is due to the weak gravitational lensing of sources in the distant Universe by foreground large-scale structure leading to coherent changes in the observed number density of the background sources. Depending on the slope of the background source number counts, cosmic magnification causes a correlation between the background and foreground galaxies, which is unexpected in the absence of lensing if the two populations are spatially disjoint. Previous attempts using submillimetre (sub-mm) sources have been hampered by small number statistics. The large number of sources detected in the {\\it Herschel} Multi-tiered Extra-galactic Survey (HerMES) Lockman-SWIRE field enables us to carry out the first robust study of the cross-correlation between sub-mm sources and sources at lower redshifts. Using ancillary data we compile two low-redshift samples from SDSS and SWIRE with $\\langle z\\rangle \\sim 0.2$ and 0.4, respectively, and cross-correlate with two sub-mm samples based on flux density and colour criteria, selecting galaxies preferentially at $z \\sim 2$. We detect cross-correlation on angular scales between $\\sim1$ and 50 arcmin and find clear evidence that this is primarily due to cosmic magnification. A small, but non-negligible signal from intrinsic clustering is likely to be present due to the tails of the redshift distribution of the sub-mm sources overlapping with those of the foreground samples. ", "introduction": "Large-scale structure at low redshifts systematically magnifies sources at higher redshifts as a result of gravitational light deflection in the weak limit. On the one hand, fewer sources will be observed, because lensing stretches the solid angle and dilutes the surface density of sources. Conversely, the effective flux limit is lowered as a result of magnification, which leads to a deeper survey. Whether there is an increase or decrease in the observed number density of sources depends on the shape of the background source number counts -- an effect known as the magnification bias (Bartelmann \\& Schneider 2001; hereafter BS01). At submillimetre (sub-mm) wavelengths the magnification bias is expected to be large and positive, resulting in an increase in the observed number density of sources compared to the case without lensing (e.g. Blain \\& Longair 1993; Blain et al. 1996; Negrello et al. 2007; Lima et al 2010). Cosmic magnification also induces an apparent angular cross-correlation between two source populations with disjoint spatial distributions. It can thus be measured by cross-correlating non-overlapping foreground and background samples. When combined with number counts, such a cross-correlation can provide constraints on cosmological parameters (e.g. $\\Omega_{{\\rm m}}, \\sigma_8$) and galaxy bias, a key ingredient in galaxy formation and evolution models (M\\'{e}nard \\& Bartelmann 2002). As the weak lensing-induced cross-correlation also probes the dark matter distribution, it provides an independent cross-check of the cosmic shear measurements, which depend on the fundamental assumption that galaxy ellipticities are intrinsically uncorrelated. Most previous investigations, using foreground galaxies selected in the optical or infrared together with background quasars, have produced controversial or inconclusive results (e.g. Seldner \\& Peebles 1979; Bartelmann \\& Schneider 1994; Bartsch et al. 1997). The best detection to date is presented in Scranton et al. (2005), where cosmic magnification is detected at an $8\\sigma$ significance level using 13 million galaxies and $\\sim$200,000 quasars from the Sloan Digital Sky Survey (SDSS). The amplitude of the weak lensing-induced cross-correlation is determined by several factors: the dark matter power spectrum and growth function, the shape of the background source number counts and the bias of the foreground sources. At sub-mm wavelengths, the power-law slope of the cumulative number count is exceptionally steep, $>$2.5 for sources in the flux range $0.02-0.5$ Jy at 250, 350 and 500 $\\mu$m (e.g. Patanchon et al. 2009; Oliver et al. 2010a; Glenn et al. 2010; Clements et al. 2010). In Scranton et al. (2005), the number count slope of the quasar sample is considerably flatter ($\\sim2$ for the brightest ones). In addition, sub-mm sources detected in deep surveys mainly reside in the high-redshift Universe with a median redshift of $z\\sim$ 2 (Chapman et al. 2003, 2005; Pope et al. 2006; Aretxaga et al. 2007; Amblard et al. 2010). The steep number counts, together with the large redshift range, make sub-mm sources an ideal background sample. So far there have been two attempts at measuring the weak lensing-induced cross-correlation between foreground optical galaxies and background sub-mm sources, but with conflicting results. Almaini et al. (2005) measured the cross-correlation between 39 SCUBA sources and optical sources at lower redshifts $\\langle z \\rangle \\sim 0.5$. They claimed evidence for a significant signal which might be caused by lensing. Conversely, Blake et al. (2006) did not find evidence for cross-correlation due to cosmic magnification using a similar number of sources. The {\\it Herschel} Multi-tiered Extra-galactic Survey (HerMES, Oliver et al. 2010b) is the largest project being undertaken by {\\it Herschel} (Pilbratt et al. 2010). In this paper, we calculate the angular cross-correlation between foreground galaxies selected from SDSS or the {\\it Spitzer} Wide-area Infrared Extragalactic (SWIRE; Lonsdale 2003, 2004) survey and background sub-mm sources detected by the Spectral and Photometric Imaging Receiver (SPIRE; Griffin et al. 2010) instrument on {\\it Herschel}. This paper is organised as follows: In Section 2, we give a brief introduction to magnification bias and the angular cross-correlation function. In Section 3, we describe the various data-sets used as foreground and background samples. Measurements of the cross-correlation between foreground and background samples are presented in Section 4. Finally, discussions and conclusions are given in Section 5. Throughout the paper, we use a spatially flat $\\Lambda$CDM cosmology with $\\Omega_{{\\rm m}} = 0.3$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Magnitudes are in the AB system. ", "conclusions": "\\label{discussions and conclusions} The unusually steep number count in the bright sub-mm regime leads to an enhanced cross-correlation signal that is due to weak gravitational lensing. In this paper, we have measured the angular cross-correlations between sub-mm sources detected by {\\it Herschel}-SPIRE in Lockman-SWIRE and foreground sources selected in the optical or near-infrared. We have also derived theoretical expectations of the weak lensing-induced cross-correlation $w_{\\rm fb}$ and the clustering-induced cross-correlation $w_{\\rm cc}$ which are in good agreement with our measurements. We find clear evidence for a lensing-induced cross-correlation between sub-mm sources at high redshifts and galaxies at low redshifts. The redshift distribution of the sub-mm sources is the biggest source of uncertainty in our analysis because most of the sources do not have spectroscopic redshifts. In principle, the clustering-induced cross-correlation $w_{\\rm cc}$ could contaminate the lensing-induced cross-correlation $w_{\\rm fb}$ if a higher than expected fraction of sub-mm sources reside in the low-redshift Universe. As the amplitude of $w_{\\rm fb}$ is mainly sensitive to the mean redshift of the background population rather than the exact shape of the $N(z)$ (M\\'{e}nard \\& Bartelmann 2002), we have carried out a simple calculation of the expected $w_{\\rm fb}$ and $w_{\\rm cc}$ amplitude by varying the mean redshift $\\langle z \\rangle$ (from 0.3 to 4.0) and the width $\\sigma_z$ (from 0.2 to 2.5), assuming the $N(z)$ of the sub-mm sources can be approximated by a Gaussian function. In all cases, to reproduce the measured cross-correlation signal, $w_{\\rm cc}$ is at most comparable to $w_{\\rm fb}$ when $\\langle z \\rangle \\sim 3.5, \\sigma_z \\sim 1.5$, $\\langle z \\rangle \\sim 2.5, \\sigma_z \\sim 1.0$ or $\\langle z \\rangle \\sim 1.5, \\sigma_z \\sim 0.5$. So the detection of the weak lensing-induced cross-correlation should be robust. It should be possible to acurately determine $N(z)$ in the future when the infrared spectral energy distributions are well understood and/or more spectroscopic redshifts are acquired for sub-mm sources. Limitations in our modelling of the cross-correlation include: using a scale- and time-independent bias factor for the galaxy-dark matter power spectrum; assuming a linearised magnification; and adopting a constant power-law number count slope independent of flux. While for this first study a simple model is adequate given the large error bars, an approach such as the halo model to describe the galaxy-dark matter power spectrum can be utilised in the future when additional data warrant an improved description (e.g. Jain et al. 2003). The expected increase in area covered by {\\it Herschel}-SPIRE will allow the detection of cosmic magnification presented in this paper to be improved and be used to constrain cosmological parameters and galaxy bias." }, "1101/1101.1450_arXiv.txt": { "abstract": "The prolific magnetar \\sgr\\ showed two outbursts in the last decade and has been closely monitored in the X-rays to track the changes in its radiative properties. We use archival \\cxo\\ and \\xmm\\ observations of \\sgr\\ to construct a history of its spectrum and persistent X-ray flux spanning a period of about seven years. We show that the decline of its X-ray flux in these two outburst episodes follows the same trend. The flux begins to decline promptly and rapidly subsequent to the flares, then decreases gradually for about 600 days, at which point it resumes a more rapid decline. Utilizing the high quality spectral data in each epoch, we also study the spectral coevolution of the source with its flux. We find that neither the magnetic field strength nor the magnetospheric properties change over the period spanned by the observations, while the surface temperature as well as the inferred emitting area both decline with time following both outbursts. We also show that the source reached the same minimum flux level in its decline from these two subsequent outbursts, suggesting that this flux level may be its steady quiescent flux. ", "introduction": "Soft Gamma Repeaters (SGRs) and Anomalous X-ray Pulsars (AXPs) belong to a class of objects called magnetars -- neutron stars whose X-ray emission is likely to be powered by the decay of their extremely strong magnetic fields (Duncan \\& Thompson 1992; Thompson \\& Duncan 1996; Thompson, Lyutikov \\& Kulkarni 2002). All seven confirmed SGRs and six out of seven confirmed AXPs\\footnotemark{}\\footnotetext{An online catalog of general properties of SGRs and AXPs can be found at http://www.physics.mcgill.ca/~pulsar/magnetar/main.html} have emitted energetic bursts of X-rays/soft gamma rays. Burst active episodes of magnetars last anywhere from few hours to months. During their bursting activity, magnetars also exhibit remarkable temporal and spectral changes in their persistent X-ray output. A detailed description of SGRs and AXPs can be found in Woods \\& Thompson (2006) and Mereghetti (2008). In the last seven years, new magnetar candidates have emerged, most prominently through transient outbursts (e.g., XTE~J1810$-$197, Halpern \\& Gotthelf 2005; CXO J164710.2-455216, Israel et al. 2007; SGR J1833$-$0832, G\\\"o\\u{g}\\\"u\\c{s} et al. 2010). These sources were too dim to be detected in X-rays during their quiescent phases (usually below our detection sensitivity), but their X-ray fluxes increased by up to few hundred times as they entered their outburst episodes. In addition, known magnetar sources also exhibit variations (triggered by bursting activity) in their persistent flux, although typically not as dramatic as those in the transient systems (e.g., Woods et al. 2004). These high X-ray luminosities were instrumental in probing the outburst mechanism of magnetars (\\\"Ozel \\& G\\\"uver 2007; G\\\"uver et al. 2007; Ng et al. 2010). In contrast, the low but persistent flux level of most magnetars requires long term (5-10 years) monitoring to understand (burst-induced) changes in their emission properties (Dib, Kaspi, \\& Gavriil 2009). \\sgr\\ has been one of the most prolific SGRs: it was discovered in 1979 (Mazets, Golenetskij \\& Guryan 1979), and was detected in a bursting mode again in 1992 (Kouveliotou et al. 1993). In May 1998 the source entered a major outburst episode that lasted about eight months and included the giant flare on 1998 August 27 (Hurley et al. 1999). ASCA and RXTE observations prior to and during the 1998 activation led to the discovery of its 5.16~s spin period (Hurley et al. 1999), its spin-down rate of $\\sim$10$^{-11}$ s/s and magnetic field of $2-8\\times10^{14}$~G (Kouveliotou et al. 1999), and thus to the confirmation of its magnetar nature. The source resumed a high level of activity in April 2001 (Guidorzi et al. 2001; Kouveliotou et al. 2001) and again in March 2006 (Vetere et al. 2006). As its burst active phases are well separated from each other, \\sgr\\ is an excellent source to investigate radiative changes both during bursting behavior as well as in burst quiescence. The first major enhancement in the persistent X-ray flux of \\sgr\\ was observed at the onset of the 1998 August 27 giant flare: the flux increased by a factor of $\\sim$700 with respect to its level before the activation (Woods et al. 2001). A detailed spectral analysis of the (much longer) flux decay period was not possible due to the lack of continuous monitoring observations with imaging instruments during the decay phase. Esposito et al. (2007) analyzed nine pointed BeppoSAX observations of \\sgr\\ spanning about five years from May 1997 to April 2002 and noted that the (2$-$10 keV) flux measured in the last pointing faded significantly with respect to earlier observations. Following the April 2001 activation, the \\cxo\\ X-ray Observatory and \\xmm\\ observed \\sgr\\ at numerous occasions, establishing a valuable dataset for understanding the long term behavior of this source in, particular, and of magnetars, in general. In this paper, we make use of all archival \\cxo\\ and \\xmm\\ observations to construct the persistent X-ray flux temporal and spectral history of \\sgr\\ spanning about seven years following the April 2001 activation. In the next section we introduce the \\cxo\\ and \\xmm\\ observations used in this study. In Section 3, we present the results of the spectral analysis and show that the magnetic field strength and the magnetospheric properties remain stable over the period spanned by the observations, while the surface temperature and the inferred emitting area both decline with time following both outbursts. We also show that the source flux shows the same trend in its decline from outburst in both episodes and ultimately reaches the same minimum flux in both cases. We discuss the implications of these results in Section 4. ", "conclusions": "We found that \\sgr\\ exhibits a monotonic but non-steady flux decline following the X-ray brightening during its 2001 and 2006 outbursts. In both outbursts, its flux dropped rapidly within a few weeks after the onset of the outburst and then at a slower rate for approximately 600~days. After this period, the flux declines again at a much faster rate. A similar decay trend was also seen in SGR~1627$-$41 (Kouveliotou et al.\\ 2003). The fact that we observe the same trend in successive outbursts from the same source suggests that the long-term effects of outbursts are not stochastic but reproducible. It is clear that magnetar bursting activity leads to long-term flux enhancements and that the additional energy powering these enhancements is stored in the crust as heat. This heat comes most likely from the energy released in the crust during the bursts, or from the energy deposited in the crust by the bombardment with magnetospheric particles during the burst. The crust, then, reradiates this additional heat over a timescale of a few years. Despite the correlation between the flux and the surface temperature, the variation in the surface temperature alone does not account for the decline in the flux, but it also requires a change in the emitting area over time. This can perhaps be understood if the crust is heated inhomogeneously, as would be expected if the initial heating episode is due to magnetic energy release. Moreover, because the thermal resistance of the crust is dominated by the uppermost layers, where the heat conductivity is strongly affected by the magnetic field, heat coming from the deeper layers of the crust could reach the surface unevenly. This leads to both uneven heating and uneven cooling, which may affect the total inferred emitting area. Furthermore, over-time, because the observations are carried out in a limited energy range, cooler parts of the crust may fall out of the observed energy band faster than the hotter regions, again reducing the inferred emitting area. Thus, the observed source flux would decline both due to a decline in temperature as $\\propto T^4$ and due to a decline in the emitting area. We find on three different occasions that the X-ray flux of \\sgr\\ was as low as $6.8 \\times 10^{-12}$~erg~cm$^{-2}$~s$^{-1}$, suggesting that this level may correspond to the source persistent X-ray flux in the absence of burst induced enhancements. The current detection thresholds of imaging instruments are $\\sim$10$^{-13}$ erg cm$^{-2}$ s$^{-1}$; it is, therefore, possible that the persistent flux levels of the so-called transient magnetars are much lower than those of the always detectable magnetars. If that is indeed the case, the magnetar engine that is responsible for a source persistent quiescent X-ray emission would seem to power a broad flux range of $\\lesssim 10^{-13}$ to $\\sim10^{-10}$~erg~cm$^{-2}$~s$^{-1}$." }, "1101/1101.6046_arXiv.txt": { "abstract": "{We consider observational constraints and fine-tuning issues in a renormalizable model of inflection point inflation, with two independent parameters. We derive constraints on the parameter space of this model arising from the WMAP 7-year power spectrum. It has previously been shown that it is possible to successfully embed this potential in the MSSM. Unfortunately, to do this requires severe fine-tuning. We address this issue by introducing a hybrid field to dynamically uplift the potential with a subsequent smooth phase transition to end inflation at the necessary point. Large parameter regions exist where this drastically reduces the fine-tuning required without ruining the viability of the model. A side effect of this mechanism is that it increases the width of the slow-roll region of the potential, thus also alleviating the problem of the fine-tuning of initial conditions. The MSSM embedding we study has been previously shown to be able to explain the smallness of the neutrino masses. The hybrid transition does not spoil this feature as there exist parameter regions where the fine-tuning parameter is as large as $10^{-1}$ and the neutrino masses remain small.} ", "introduction": "\\label{section:introduction} An outstanding goal for the theory of primordial inflation is to connect it to particle physics, and in particular to the Standard Model (SM) and its extensions \\cite{RM}. For this to happen it would appear that a low inflationary scale and sub-Planckian VEVs are necessary. These features are naturally present when inflation is generated about a {\\it point of inflection} in the potential, due to the flatness of the potential at the inflection point. This idea can be illustrated with a simple renormalizable scalar potential with two independent parameters, $A$ and $B$: \\beq \\label{genericpotential} V(\\phi)=A\\phi^2-C\\phi^3+B\\phi^4\\,, \\eeq where $C$ is determined in terms of $A$ and $B$ in order to obtain a point of inflection suitable for inflation. The VEV at which inflation occurs is closely related to the two independent parameters and can take a wide range of values below $M_{P}=2.4\\times 10^{18}$~GeV for different values of $(A,B)$. The above renormalizable potential is sufficiently simple that it can be embedded in a range of particle theories beyond the SM. Of particular interest is a model of low scale supersymmetry (SUSY), where the origin of $\\phi$ can be directly linked to SUSY partners of the SM Higgs and leptons, along with the right-handed (RH) neutrinos \\cite{Allahverdi:2006cx,Allahverdi:2007wt}.\\footnote{The combination is a {\\it gauge invariant} D-flat direction of $MSSM\\times U(1)_{B-L}$. (For a review on SUSY flat directions, see~\\cite{MSSM-REV}).} This model has been shown to produce a power spectrum of perturbations that is consistent with observation \\cite{Allahverdi:2007wt}, while explaining the small scale of the observed neutrino masses, and providing a dark matter candidate from the RH sneutrino component of the inflaton \\cite{Allahverdi:2007wt} in a simple extension of minimal supersymmetric Standard Model (MSSM).\\footnote{The first examples of an MSSM {\\it gauge invariant} inflaton are given in \\cite{Allahverdi:2006iq,Allahverdi:2006we} and the parameter space for the detection of MSSM inflatons and neutralino dark matter at the LHC was studied in \\cite{Allahverdi:2007vy,Allahverdi:2010zp}.} The advantage of such a particle physics embedding is that the model parameters are motivated by low scale SUSY within a \\emph{visible} sector. Therefore, one can track the thermal history of the universe and probe the inflaton origin at the LHC, while also constraining the potential from cosmic microwave background (CMB) observations. In Section~\\ref{section:constraints} of this paper we investigate which regions of parameter space are consistent with constraints from the WMAP 7-year results. It is known that in order to maintain sufficient flatness of the potential to reproduce the observed CMB spectrum, a fine-tuning of the parameters $A$, $B$ and $C$ is required~\\cite{Allahverdi:2006we,Bueno Sanchez:2006xk,Lalak:2007rsa,Enqvist:2007tf}. At low scales this tuning is very acute and a significant challenge to overcome. For high or intermediate scale inflation the required tuning can be reduced to some extent but remains problematic. Although the string landscape can perhaps naturally account for the fine tuning of soft SUSY breaking terms from degenerate vacua~\\cite{Allahverdi:2007wh}, a dynamical solution to the problem is desirable.\\footnote{In the context of a non-renormalizable potential for inflection point inflation, a different approach to the fine-tuning problem has been presented in ref.~\\cite{Allahverdi:2010zp}. Using the renormalization group equations, the tuning of the ratio of SUSY breaking terms can instead be viewed as an equal tuning of the non-renormalizable coupling.} The fine-tuning required to fit the spectrum constraints can also be reduced by simply raising the scale of inflation, as first pointed out in \\cite{Enqvist:2010vd}. However, on its own this violates the requirement to generate a suitable number of e-folds. We propose to include a new hybrid scalar field which provides a vacuum energy while it remains trapped in a false minimum. On being released from this false minimum, the field rolls quickly to its true minimum and brings slow-roll to a premature end, in exactly the same manner as hybrid inflation~\\cite{Linde:1993cn}. This extension can significantly reduce the amount of fine-tuning required in the model, while still matching observations and the e-fold constraint, without ruining any of the attractive features of the MSSM embedding. In Section~\\ref{section:MSSM} we briefly discuss the aspects of the SUSY embedding of \\eqref{genericpotential}, define the measure of fine-tuning of the potential, and obtain expressions for the slow-roll parameters that are used later in the paper. The bulk of Section~\\ref{section:MSSM} is however devoted to a calculation of the e-fold number that corresponds to the observed CMB scales. This calculation is not original but it provides important results that are used in Sections~\\ref{section:constraints} and~\\ref{section:hybrid}, which contain the new results of the paper. In Section~\\ref{section:constraints} we investigate the region of parameter space that allows for a period of inflation that is consistent with the e-fold constraint and constraints from the WMAP 7-year power spectrum \\cite{Komatsu:2010fb}. In Section~\\ref{section:hybrid}, we introduce the hybrid extension to the model and show how it reduces the required fine-tuning. ", "conclusions": "\\label{section:conclusion} We have discussed the very generic renormalizable inflationary potential introduced in eq.~\\eqref{genericpotential} and have analyzed the conditions required for this potential to generate a power spectrum of density perturbations compatible with the latest WMAP 7-year constraints. We have also presented a particular particle physics model for the potential \\eqref{potential}, where the origin of the inflaton is the gauge-invariant flat direction introduced in eq.~\\eqref{superpotential}, with a minimal extension of the standard model gauge group. We obtained the constraints on the parameter space for this model from cosmology, which are depicted in figure~\\ref{figure:m-h}. The allowed region of parameter space also allows an explanation for the small neutrino mass. The potential \\eqref{genericpotential} (or equivalently \\eqref{potential}) can be realised in many different theoretical models. As long as the theory allows for rapid transfer of inflaton energy into radiation at the end of inflation, the constraints on parameter space presented in figure~\\ref{figure:m-h} will apply. Therefore our results are quite general. Although in general the soft SUSY-breaking terms in this potential must be highly tuned against each other, we have shown that this is primarily a result of the requirement to produce the correct number of e-folds of inflation, rather than a constraint imposed by the observed power spectrum. With this perspective, we have presented a simple extension of the model to include a new scalar field which brings an end to inflation through the hybrid mechanism. We show this reduces the fine-tuning to very manageable levels, even achieving $\\vert\\beta\\vert\\sim10^{-1}$. This result is significant as the required fine-tuning was one of the main objections to the original model. Finally, we argue that this hybrid extension of the model also reduces any need for fine-tuning of the initial value of the inflaton field, which also makes the hybrid extension more attractive." }, "1101/1101.2874_arXiv.txt": { "abstract": "We investigate a model in which Dark Matter is stabilized by means of a $Z_2$ parity that results from the same non-abelian discrete flavor symmetry which accounts for the observed pattern of neutrino mixing. In our $A_4$ example the standard model is extended by three extra Higgs doublets and the $Z_2$ parity emerges as a remnant of the spontaneous breaking of $A_4$ after electroweak symmetry breaking. We perform an analysis of the parameter space of the model consistent with electroweak precision tests, collider searches and perturbativity. We determine the regions compatible with the observed relic dark matter density and we present prospects for detection in direct as well as indirect Dark Matter search experiments. ", "introduction": "\\label{intro} The existence of non-baryonic Dark Matter (DM) is well established by cosmological and astrophysical probes. However, despite the great experimental effort over many years, its nature still remains elusive. Elucidating the long-standing puzzle of the nature of dark matter constitutes one of the most important challenges of modern cosmology and particle physics. The various observations and experiments, however, constrain some of its properties \\cite{Bertone:2004pz,Taoso:2007qk}. Among the most important requirements a DM candidate is required to satisfy are neutrality, stability over cosmological time scales, and agreement with the observed relic density. While the neutrality of the particle is usually easy to accomodate in models, its stability in general is assumed in an {\\it ad-hoc} fashion. From a particle physics point of view, the stability suggests the existence of a symmetry that forbids the couplings that would otherwise induce the decay. Typically, the most common way to stabilize the DM particle is to invoke a $Z_2$ parity, an example of which is R parity in supersymmetry. It would certainly be more appealing to motivate such a symmetry from a top-down perspective. Different mechanisms have been suggested to achieve this \\cite{Hambye:2010zb}, for instance using $U(1)$ gauge symmetries \\cite{Frigerio:2009wf,Kadastik:2009dj,Batell:2010bp} to get for example R-parity in the MSSM from a $U(1)_{B-L}$\\cite{Martin:1992mq}) symmetery, global symmetries, accidental symmetries \\cite{Cirelli:2005uq} or custodial symmetry. A new mechanism of stabilizing the DM has been recently proposed in Ref.~\\cite{Hirsch:2010ru} in which DM stability originates from the flavor structure of the standard model. Indeed the same discrete flavor symmetry which explains the pattern of neutrino mixing~\\cite{Schwetz:2008er} can also stabilize the dark matter \\footnote{Models based on non-Abelian discrete symmetries but with a decaying dark matter candidate can be found for example in Ref.\\cite{Kajiyama:2010sb}.}. This opens an attractive link between neutrino physics and DM~\\footnote{Other mechanisms of relating DM and neutrinos include the majoron DM~\\cite{Berezinsky:1993fm,Lattanzi:2007ux,Bazzocchi:2008fh}.}; two sectors that show a clear need for physics beyond the Standard Model. The model proposed in Ref.~\\cite{Hirsch:2010ru} is based on an $A_4$ symmetry extending the Higgs sector of the SM with three scalar doublets. After electroweak symmetry breaking two of the scalars of the model acquire vacuum expectation values (vev) which spontaneously break $A_4$ leaving a residual $Z_2$. The lightest $Z_2$ neutral odd scalar is then automatically stable and will be our DM candidate. On the other hand, the fermionic sector is extended by four right handed neutrinos which are singlets of \\321. Light neutrino masses are generated via a type I see-saw mechanism~\\cite{Minkowski:1977sc,gell-mann:1980vs,yanagida:1979,mohapatra:1980ia,schechter:1980gr}, obey an inverted hierarchy with $m_{\\nu3}=0$ and vanishing reactor neutrino angle \\footnote{For a similar realisation see~\\cite{Meloni:2010sk}.}. For pioneer studies on the use of $A_4$ for neutrino physics see \\cite{Babu:2002dz}. We study the regions in parameter space of the model where the correct dark matter relic density is reproduced and the constraints from accelerators are fullfilled. We then consider the prospects for direct dark matter detection in underground experiments. We show that the model can potentially explain the DAMA annual modulation data \\cite{Bernabei:2010mq,Bernabei:2000qi} as well as the excess recently found in the COGENT experiment \\cite{Aalseth:2010vx}. We show that present upper limits on the spin independent DM scattering cross section off nucleons can already severely constrain the parameter space of the model. Indirect dark matter searches through astrophysical observations are not currently probing the model apart from some small regions of the parameter space where the dark matter annihilation cross section is enhanced via a Breit-Wigner resonance. The paper is organized as follows: in Sec.\\ref{model} we present the model, in Sec.~\\ref{constraints} the constraints from collider data are reviewed and in Section~\\ref{relic} we study the viable regions of the parameter space. In Sec.~\\ref{direct} and \\ref{indirect} we sketch the prospects for direct and indirect dark matter detection. Finally, we summarize our conclusions in Sec.~\\ref{conclusions}. ", "conclusions": "\\label{conclusions} We have studied a model where the stability of the dark matter particle arises from a flavor symmetry. The $A_4$ non-abelian discrete group accounts both for the observed pattern of neutrino mixing as well as for DM stability. We have analysed the constraints that follow from electroweak precision tests, collider searches and perturbativity. Relic dark matter density constraints exclude the region of the parameter space where simultaneously $M_{DM}\\lesssim 40$ GeV and $M_{H}\\gsim 400$ GeV because of the resulting over-abundance of dark matter. We have also analysed the prospects for direct and indirect dark matter detection and found that, although the former already excludes a large region in parameter space, we cannot constrain the mass of the DM candidate. In contrast, indirect DM detection is not yet sensitive enough to probe our predictions. However, forecasted sensitivities indicate that Fermi-LAT should start probing them in the near future. All of the above relies mainly on the properties of the scalar sector responsible for the breaking of the gauge and flavour symmetry. A basic idea of our approach is to link the origin of dark matter to the origin of neutrino mass and the understanding of the pattern of neutrino mixing, two of the most oustanding challanges in particle physics today. At this level one may ask what are the possible tests of this idea in the neutrino sector. Within the simplest scheme described in Ref.~\\cite{Hirsch:2010ru} one finds an inverted neutrino mass hierarchy, hence a neutrinoless double beta decay rate accessible to upcoming searches, while $\\theta_{13}=0$ giving no CP violation in neutrino oscillations. Note however that the connection of dark matter to neutrino properties depends strongly on how the symmetry breaking sector couples to the leptons." }, "1101/1101.4039_arXiv.txt": { "abstract": "The properties of nearby galaxies that can be observed in great detail suggest that a better theory rather than cold dark matter (CDM) would describe in a better way a mechanism by which matter is more rapidly gathered into large-scale structure such as galaxies and groups of galaxies. In this work we develop and simulate a hydrodynamical approach for the early formation of structure in the Universe, this approach is based on the fact that dark matter is on the form of some kind of scalar field (SF) with a potential that goes as $\\mu^2\\Phi^2/2+\\lambda\\Phi^4/4$, we expect that the fluctuations coming from the SF will give us some information about the matter distribution we observe these days. ", "introduction": "We begin this work remebering the framework of the standard cosmological model: a homogeneous and isotropic Universe whose evolution is best described by Friedmann's equations that come from general relativity and whose main ingridients can be described by fluids whose characteristics are very similar to those we see in our Universe. Of course, the Universe is not exactly homogeneous and isotropic but this standard model does give us a framework within which we can study the evolution of structure like the observed galaxies or clusters of galaxies from small fluctuations in the density of the early Universe. In this model 4 per cent of the mass in the Universe is in the baryons, 22 per cent is non-baryonic dark matter and the rest in some form of cosmological constant. Another idea that has been around just a bit less than hundred years and in which many of the cosmological models are based in is that of an homogeneous and isotropic Universe, although it has always been clear that this homogeneity and isotropy are only found until certain level. Now we know that the anisotropies are very important and can grow as big as the large scale structure we see today. Nowadays the most accepted model in cosmology which explains the evolution of the Universe is known as $\\Lambda$CDM, because it has achieved some observations with outstanding success, like for example the fact that the cosmic microwave background can be explained in great detail and that it provides a framework within one can understand the large-scale isotropy of the Universe and important characteristics on the origin, nature and evolution of density fluctuations which are believed to give rise to galaxies and other cosmic structure. There remain, however, certain problems at galactic scales, like the cusp profile of central densities in galactic halos, the over 500 substructures predicted by numerical simulations which are not found in observations, etc.. See for example \\cite{b16}, \\cite{b4} and \\cite{b15}. In the big bang model, gravity plays an essential role, it collects the dark matter in concentrated regions called 'Dark matter haloes'. In the large dark matter haloes, the baryons are believed to be so dense as to radiate enough energy so they will collapse into galaxies and stars. The most massive haloes, hosts for the brightest galaxies, are formed in regions with the highest local mass density. Less massive haloes, hosts for the less bright galaxies, appear in regions with low local densities, i.e, regions were the local density is not well defined, \\cite{b14}. These situations appear to be the same as in our extragalactic neighborhood, but still there are problems. Observations point out to a better understanding of the theory beginning with the less occupied space called the 'Local void', which contains just a few galaxies which are bigger than the expected. This problem would be solved, if the structure grew faster than it does in the standard theory, therefore filling the local void and giving rise to more matter in the surroundings, \\cite{b14}. Another problem arises for the so called 'Pure disk galaxies', which do not appear in numerical simulations of structure formation in the standard theory, because it is believed that their formation which is relatively slow began in the thick stellar bulges. Again this problem would be solved for the early formation of structure. The incorporation of a new kind of dark matter, different from the one proposed by the $\\Lambda$CDM model into the big bang theory holds out the possibility of resolving some of these issues. Recent works have introduced a dynamic scalar field with a certain potential $V(\\Phi)$ as a candidate to dark matter, although there is not yet an agreement for the correct form of the potential of the field. \\cite{b20}, and independently \\cite{b12} suggested bosonic dark matter (SFDM) as a model for galactic halos. Another interesting work pointing this way was done by \\cite{b13} and independently by \\cite{b21} where they used a potential of the form $\\cosh$ to explain the core density problem for disc galaxy halos in the $\\Lambda$CDM model. \\cite{b12} presented a model for the dark matter in spiral galaxies, in which they supposed that dark matter is a scalar field endowed with a scalar potential. Several recent work have also suggested that SFDM can be composed of spin-0 bosons which give rise to Bose-Einstein Condensates (BECs), which at the same time can make up the galaxies we are observing in our Universe. \\cite{b6} proposed that dark matter is composed of ultra/light scalar particles who are initially in the form of a BEC. In their work \\cite{b18} used a bosonic dark matter model to explain the structure formation via high-resolution simulations, finally \\cite{b17} reviewed the key properties that may arise from the bosonic nature of SFDM models. The main objective of this work in difference with others is to introduce SFDM and assume that dark matter its itself a scalar field that involves an auto-interacting potential of the form $V(\\Phi)=\\mu^2\\Phi^2/2+\\lambda\\Phi^4/4$, where $\\mu_{\\Phi}\\sim 10^{-22}$ eV is the mass of the scalar field, \\cite{b20}, \\cite{b10} and \\cite{b6}. With the mass $\\mu_{\\Phi} \\sim 10^{-22}$ eV and only one free parameter when $\\lambda$ is taken equals to zero, the SFDM model fits the following important features: \\begin{enumerate} \\item The cosmological evolution of the density parameters of all the components of the Universe, \\cite{b9}. \\item The rotation curves of galaxies, \\cite{b3}, and the central density profile of LSB galaxies, \\cite{b2}, \\item With this mass, the critical mass of collapse for a real scalar field is just $10^{12}\\,M_{\\odot}$, i.e., the one observed in galactic haloes, \\cite{b1}. \\item The central density profile of the dark matter is flat, \\cite{b2}. \\item The scalar field has a natural cut off, thus the substructures in clusters of galaxies is avoided naturally. With a scalar field mass of $\\mu_\\Phi\\sim10^{-22}$ eV the amount of substructures is compatible with the observed one, \\cite{b10}. \\end{enumerate} In this paper we show that the SFDM predicts galaxy formation earlier than the cold dark matter model, because they form BEC at a critical temperature $T_c >>$ TeV. So, if SFDM is right, this would imply that we have to see big galaxies at high redshifts. In order to do this, we study the density fluctuations of the scalar field from a hydrodynamical point of view, this will give us some information about the energy density of dark matter halos necessary to obtain the observational results of large-scale structure. Here we will give some tools that might be necessary for the study of the early formation of structure. In section \\ref{fondo} we analyse the analytical evolution of the SF, then in section \\ref{fluct} we treat the SF as a hydrodynamical fluid in order to study its evolution for the density contrast, in section \\ref{results} we compare our results with those obtained by the CDM model for the density contrast in the radiation dominated era just before recombination and finally we give our conclusions. \\section[]{The Background}\\label{fondo} In this section we perform a transformation in order to solve the Friedmann equations analytically with the approximation $H<<\\mu_{\\Phi}$. The scalar field (SF) we deal with depends only on time, $\\Phi=\\Phi_0(t)$, and of course the background is only time dependent as well. We use the Friedmann-Lema\\^itre-Robertson-Walker (FLRW) metric with scale factor $a(t)$. The background Universe is composed only by SFDM ($\\Phi_{0}$) endowed with a scalar potential. We begin by recalling the basic background equations. From the energy-momentum tensor $\\mathbf{T}$ for a scalar field, the scalar energy density $T_0^0$ and the scalar pressure $T_j^i$ are given by \\begin{equation} T_0^0=-\\rho_{\\Phi_0}=-\\left(\\frac{1}{2}\\dot{\\Phi}_0^2+ V\\right), \\label{rhophi0} \\end{equation} \\begin{equation} T_j^i=p_{\\Phi_0}=\\left(\\frac{1}{2}\\dot{\\Phi}_0^2-V\\right)\\delta_j^i, \\label{pphi0} \\end{equation} where the dots stand for the derivative with respect to the cosmological time and $\\delta^i_j$ is Kronecker's delta. Thus, the Equation of State (EoS) for the scalar field is $p_{\\Phi_{0}}=\\omega_{\\Phi_{0}}\\,\\rho_{\\Phi_0}$ with \\begin{equation} \\omega_{\\Phi_{0}}=\\frac{\\frac{1}{2}\\dot{\\Phi}_{0}^{2}\\,-\\,V}{\\frac{1}{2}\\dot{\\Phi}_{0}^{2}\\,+\\,V}. \\label{ec:w} \\end{equation} Notice that background scalar quantities have the subscript $0$. Now the following dimensionless variables are defined \\begin{eqnarray} x\\equiv\\frac{\\kappa}{\\sqrt{6}}\\frac{\\dot{\\Phi}_{0}}{H},\\quad u\\equiv\\frac{\\kappa}{\\sqrt{3}}\\frac{\\sqrt{V}}{H},\\nonumber \\label{eq:varb} \\end{eqnarray} being $\\kappa^{2} \\equiv 8\\pi G$ and $H \\equiv \\dot{a}/a$ the Hubble parameter. Here we take the scalar potential as $V=m^{2}\\Phi^2/2\\hbar^2+\\lambda\\Phi^4/4$, where, $\\mu=mc/\\hbar$ and $m$ is the mass given in kilograms, and from now on we will use units where $c=1$, then for the ultra-light boson particle we have that $\\mu_{\\Phi}\\sim 10^{-22}$ eV. With these variables, the density parameter $\\Omega_\\Phi$ for the background $0$ can be written as \\begin{eqnarray} \\Omega_{\\Phi_{0}}=x^2+u^2. \\label{eq:dens} \\end{eqnarray} In addition, we may write the EoS of the scalar field as \\begin{equation} \\omega_{\\Phi_{0}}=\\frac{x^{2}-u^{2}}{\\Omega_{\\Phi_{0}}}. \\label{eq:dlw} \\end{equation} Since $\\omega_{\\Phi_{0}}$ is a function of time, if its time average tends to zero, this would imply that $\\Phi^2$-dark matter can be able to mimic the EoS for CDM, see \\cite{b11} and \\cite{b9}. Now we express the SF, $\\Phi_0$, in terms of the new variables $S$ and $\\hat\\rho_0$, where $S$ is constant in the background and $\\hat{\\rho}_0$ will be the energy density of the fluid also in the background. So, our background field is proposed as \\begin{equation} \\Phi_0=(\\psi_0\\rmn e^{-\\rmn imt/\\hbar}+\\psi^*_0\\rmn e^{\\rmn imt/\\hbar}) \\end{equation} where, \\begin{equation} \\psi_0(t)=\\sqrt{\\hat{\\rho}_0(t)}\\rmn e^{\\rmn iS/\\hbar} \\end{equation} and with this our SF in the background can be finally expressed as, \\begin{equation} \\Phi_0=2\\sqrt{\\hat{\\rho}_0}\\cos(S-mt/\\hbar), \\label{tri} \\end{equation} with this we obtain \\begin{eqnarray} \\dot{\\Phi}_0^2&=&\\hat{\\rho}_0\\left[\\frac{\\dot{\\hat{\\rho}}_0}{\\hat{\\rho}_0}\\cos(S-mt/\\hbar)\\right.\\nonumber\\\\ &-&\\left.2(\\dot{S}-m/\\hbar)\\,\\sin(S-mt/\\hbar)\\right]^2 \\label{backSF} \\end{eqnarray} To simplify, observe that the uncertanty relation implies that $m\\Delta t\\sim\\hbar$, and for the background in the non-relativistic case the relation $\\dot S/m\\sim0$ is satisfied. Notice also that for the background we have that the density goes as $(\\ln\\hat{\\rho}_0)\\dot{}=-3H,$ but we also have that $H\\sim 10^{-33}$ eV $<<\\mu_{\\Phi}\\sim 10^{-22}$ eV, so with these considerations at hand for the background, in (\\ref{backSF}) we have \\begin{equation} \\dot{\\Phi}_0^2=4\\frac{m^2}{\\hbar^2}\\hat{\\rho}_0\\sin^2(S-mt/\\hbar) \\end{equation} Finally, substituting this last equation and equation (\\ref{tri}) into (\\ref{rhophi0}) when taking $\\lambda=0$, we obtain \\begin{equation} \\rho_{\\Phi_0}=2\\frac{m^2}{\\hbar^2}\\hat{\\rho}_0[\\sin^2(S-mt/\\hbar)+\\cos^2(S-mt/\\hbar)]=2\\frac{m^2}{\\hbar^2}\\hat{\\rho}_0. \\label{trigo} \\end{equation} Comparing this result with (\\ref{eq:dens}) we have that the identity $\\Omega_{\\Phi_0}=2m^2\\hat{\\rho}_0/\\hbar^2$ holds for the background, so comparing with (\\ref{trigo}), \\begin{equation} x=\\sqrt{2\\hat{\\rho}_0}\\frac{m}{\\hbar}\\sin(S-mt/\\hbar) \\label{eq:cinet} \\end{equation} \\begin{equation} u=\\sqrt{2\\hat{\\rho}_0}\\frac{m}{\\hbar}\\cos(S-mt/\\hbar). \\label{eq:potencial} \\end{equation} We plot the evolution of the energies (\\ref{eq:cinet}) and (\\ref{eq:potencial}) in Fig. \\ref{fig1}, where for the evolution we used the e-folding number $N$ defined as $N=ln(a)$ and the fact that $a\\sim t^n\\rightarrow t\\sim\\rmn{e}^{N/n}.$ In terms of the two analytic results (\\ref{eq:cinet}) and (\\ref{eq:potencial}) Fig. \\ref{fig1} shows the kinetic and the potential energies of the scalar field. \\begin{figure} \\scalebox{0.6}{\\includegraphics{fig1top.eps}} \\scalebox{0.6}{\\includegraphics{fig1bottom.eps}} \\caption{Analytical evolution of the potential (top panel) and kinetic (bottom panel) energies of the scalar field dark matter.} \\label{fig1} \\end{figure} Observe the excelent accordance with the numerical results in \\cite{b9} for the kinetic and potential energies of the background respectively. ", "conclusions": "The new observational instruments and telescopes until today have perceived objects as far as $z=8.6$, \\cite{b7}. The cosmic background radiation can bring us information from $z=1000$ to $z=2000$. But jet we can not see anything from the intermediate region, now we know of a possible galaxy that might be found at a distance of $z=10.56$ but it has jet to be confirmed. As seen earlier, as expected for the CDM model we obtained that for the matter dominated era the low-$k$ modes grow, when CDM decouples from radiation in a time just before recombination it grows in a milder way than it does in the matter dominated era (Fig. \\ref{fig2}). Although in general a scalar field is not a fluid, it can be treated as if it behaved like one.The evolution of its density can be the appropriate for the purpose of structure formation, because locations with a high density of dark matter can support the formation of galactic structure. In this work we have assumed that there is only one component to the mass density, and that this component is given by the scalar field dark matter. In this case equation (\\ref{delta}) is valid for all sub-horizon sized perturbations in our non-relativistic specie, so for sub-horizon perturbations the newtonian treatment worked with in the evolution of the perturbations suffices. The SFDM has provided to be an alternative model for the dark matter nature of the Universe. We have shown that the scalar field with an ultralight mass of $10^{-22}$ eV simulates the behavior of CDM in a Universe dominated by matter when $\\lambda=0$, because in general in a matter dominated Universe for low-$k$, $v_\\rmn{q}$ tends to be a very small quantity tending to zero, so from (\\ref{delta}) we can see that on this era we will have the CDM profile given by (\\ref{deltaCDM}), i.e., the SFDM density contrast profile is very similar to that of the $\\Lambda$CDM model, Fig. \\ref{fig2}. On the contrary for $\\lambda\\neq 0$ both models have different behavior as we can see from Fig. \\ref{fig3}, results which show that linear fluctuations on the SFDM can grow in comparison with those of CDM, even at early times when the large-scale modes (small $k$) have entered the horizon just after $a_{eq}\\sim 10^{-4}$, when it has decoupled from radiation, so the amplitudes of the density contrast start to grow faster than those for CDM around $a\\sim 10^{-2}$. Here an important point is that although CDM can grow it does so in a hierarchical way, while from Fig. \\ref{fig3} we can see that SFDM can have bigger fluctuations just before the $\\Lambda$CDM model does, i.e., it might be that no hierarchical model of structure formation is needed for SFDM and is expected that for the non-linear fluctuations the behaviour will be quite the same as soon as the scalar field condensates, in a very early epoch when the energy of the Universe was about $\\sim$ TeV. These facts can be the crucial difference between both models. As mentioned before, recent observations have taken us to very early epochs in the origin of the Universe, and have made us think that structure had already been formed, corresponding to $z\\approx 7$. It is clear from Fig. \\ref{fig3} that at recombination $z\\approx 1300$ there already existed defined perturbations in the energy density for the SFDM model, which can contribute to the early formation of structure. Then, if clusters could be formed as early as these z's, this would imply that $\\Phi^2+\\lambda\\Phi^4$ as a model for dark matter could give an explanation for the characteristic masses that are being observed, and therefore it could solve some of the problems present in the standard $\\Lambda$CDM model. Although the observational evidence seems to be in favor of some kind of cold dark matter, the last word has not been said. Astronomers hope to send satellites that will detect the finest details of the cosmic background radiation, which will help us to get information of structure at the time of recombination, from which it will be possible to deduce its evolution until now a days." }, "1101/1101.1270_arXiv.txt": { "abstract": "From the general principle that gravity originates from the coupling and thermal equilibrium between matter and vacuum background, we give two simple equations to calculate the quantum effect of gravity. From these two equations, we calculate the ratio between the dark energy density and that of the sum of the dark matter density and ordinary matter density. Without any fitting parameter, the ratio is calculated as $2.36$, which agrees quantitatively with the result $7/3$ obtained from various astronomical observations. The same quantum gravity effect explaining dark energy will also lead to abnormal gravity for a sphere full of superfluid helium. It is shown that with an atom interferometer placed in this sphere, the accuracy $\\Delta g/g$ below $10^{-8}$ could be used to test our idea, which satisfies the present experimental technique of atom interferometer. The abnormal gravity effect would have important potential applications in black hole, possible experimental test of many-world interpretation, and even in condensed matter physics and material physics. ", "introduction": " ", "conclusions": "" }, "1101/1101.1046_arXiv.txt": { "abstract": " ", "introduction": "In his contribution, Tony Moffat has described the properties of WR\\,140 (= HD 193793) and the motivation for the 2009 campaign. Fundamental to colliding-wind binaries is the power in their winds. The WC7 and O5 stars in WR\\,140 have fast ($\\sim$ 3000 km$^{-1}$) winds carrying $\\sim 2 \\times 10^{-5}$ and $\\sim 2 \\times 10^{-6} M_{\\odot}$ y$^{-1}$ mass-loss respectively. The kinetic powers of these winds are in excess of $10^4$ and $10^3 L_{\\odot}$. % Of these, $\\sim 3 \\times 10^3 L_{\\odot}$ is dissipated where the winds collide (the % wind-collision region, WCR), leading to shock-heating and compression of the plasma, X ray emission, synchrotron radio emission, changes to the profiles of some emission lines and condensation of `dust' (really something like soot) in the shock-compressed wind. Of these effects, dust formation was certainly the most unexpected. The 1977 dust formation episode from WR\\,140 was independently discovered (Williams et al. 1977, 1978) at the 1.5-m Infrared Flux Collector (now the Carlos S\\'anchez Telescope) at Iza\\~{n}a on Tenerife -- only a short distance from the Mons telescope. From 1977 to the 2001 periastron, the thermal emission from dust made by WR\\,140 was tracked with infrared photometers at a variety of wavelengths between 1 and 20 $\\mu$m and, in the years following the 2001 dust-formation episode, with infrared cameras which imaged the dust clouds. These observations have been published, but I will summarise the most recent work (Williams et al. 2009) to show how much we can learn from observing in this wavelength domain. For the 2009 campaign, we (Watson Varricatt, Andy Adamson and the author) extended the observations of the 1.083-$\\mu$m He\\,{\\sc i} line taken around the 2001 periastron passage (Varicatt, Williams \\& Ashok 2004). This line has a P Cygni profile, with a flat-topped emission component like the $\\lambda$5696 \\AA\\ C\\,{\\sc iii} line observed in the optical campaign, and develops a strong sub-peak near periastron passage. Its absorption component varies as we observe the stars through different parts of the wind. The new observations were taken in 2008 with UIST on the United Kingdom Infrared Telescope (UKIRT) at a resolution of 200 km~s$^{-1}$ and comprise five spectra extending the coverage to earlier phases ($\\phi$ = 0.93--0.95) and nine spectra near $\\phi$ = 0.99, when the WR\\,140 system was changing very rapidly. ", "conclusions": "" }, "1101/1101.3866_arXiv.txt": { "abstract": "The XENON100 experiment, located at the Laboratori Nazionali del Gran Sasso (LNGS), aims to directly detect dark matter in the form of Weakly Interacting Massive Particles (WIMPs) via their elastic scattering off xenon nuclei. We present a comprehensive study of the predicted electronic recoil background coming from radioactive decays inside the detector and shield materials, and intrinsic radioactivity in the liquid xenon. Based on GEANT4 Monte Carlo simulations using a detailed geometry together with the measured radioactivity of all detector components, we predict an electronic recoil background in the WIMP-search energy range and 30~kg fiducial mass of less than 10$^{-2}$~events$\\cdot$kg$^{-1}\\cdot$day$^{-1}\\cdot$keV$^{-1}$, consistent with the experiment's design goal. The predicted background spectrum is in very good agreement with the data taken during the commissioning of the detector in Fall~2009. ", "introduction": "\\label{intro} For all experiments dealing with very low signal rates, such as dark matter or double beta decay searches, the reduction and discrimination of the background is one of the most important and difficult tasks. As the sensitivity of these experiments keeps increasing, the fight against the background remains crucial. The XENON100 detector, which is installed in the Laboratori Nazionali del Gran Sasso (LNGS), Italy, is the second generation detector within the XENON program, dedicated to the direct detection of particle dark matter in the form of Weakly Interacting Massive Particles (WIMPs)~\\cite{wimps}. It is the successor of XENON10~\\cite{xe10-instrument}, which has set some of the best limits on WIMP-nucleon scattering cross-sections~\\cite{xe10-independent, xe10-dependent}. XENON100 aims to improve this sensitivity due to an increase of the target mass and a significant reduction of the background in the target volume. In the standard scenario, WIMPs are expected to elastically scatter off xenon nuclei resulting in low energy nuclear recoils. Neutrons passing through the detector also produce nuclear recoils of similar energy, whereas gamma rays and electrons produce electronic recoils. This opens the possibility to efficiently reject the electromagnetic background using various discrimination techniques. Experiments like XENON100 distinguish electronic interactions from nuclear recoils based on a different ratio in the yield of scintillation light (primary signal, S1) and ionization charge (secondary signal, S2). Using this discrimination technique, XENON10 and XENON100 reached an electronic recoil rejection efficiency better than 99\\% at 50\\% nuclear recoil acceptance~\\cite{xe100-independent, xe10-independent}. The main sources of electronic recoil background in XENON100 are radioactive contamination of the materials used to construct the detector and the shield, intrinsic radioactivity in the LXe target, and the decays of $^{222}$Rn and its progeny inside the detector shield. Even if the electronic recoil rejection efficiency, based on the ratio of the scintillation and ionization signals, is high, a potential statistical leakage of electronic recoil events into the nuclear recoil region can mimic a dark matter signal. One way to handle this background is a well-planned detector design avoiding the presence of radioactive materials close to the active volume. During the design phase of XENON100, all detector materials and components have been carefully selected based on measurements of their radioactive contamination in order to achieve a low background level. The background is further suppressed by improvements on the passive shield and by surrounding the target volume with an active LXe veto layer. In this paper we summarize the effort to use extensive Monte Carlo simulations to predict the electronic recoil background of XENON100 from natural radioactivity in the detector and shield components, and to study the background reduction by applying fiducial volume and veto coincidence cuts. Section~\\ref{secDetectorModel} describes the detector model which has been used in the simulations. The predicted electronic recoil background from the detector and shield materials is discussed in Section~\\ref{secDetectorMaterials}, the background from the decays of $^{222}$Rn and its progenies in the shield cavity in Section~\\ref{secRadonCavity}, and the background from radon and krypton in LXe in Section~\\ref{secIntrinsicBG}. The comparison of the background model with the measured background spectrum is presented in Section~\\ref{secComparison}, and conclusions are drawn in Section~\\ref{secConclusions}. ", "conclusions": "\\label{secConclusions} An extensive study to predict the electronic recoil background of the XENON100 experiment has been performed . The study is based on Monte Carlo simulations with GEANT4 using a detailed mass model of the detector and its shield, and the measured radioactivity values of all relevant detector components. The design goal of XENON100, to gain a factor of 100 reduction in background rate compared to XENON10 (0.6~events$\\cdot$kg$^{-1}\\cdot$day$^{-1}\\cdot$keV$^{-1}$~\\cite{xe10-independent}), has been achieved. This has been possible thanks to a selection of all detector materials, an innovative design of the cryogenic system, the use of an active LXe veto, and an improved passive shield. The predicted rate of single scatter electronic recoil events in the energy region below 100~keV, without veto coincidence cut is 15.8$\\times$10$^{-3}$ (9.6$\\times$10$^{-3}$) events$\\cdot$kg$^{-1}\\cdot$day$^{-1}\\cdot$keV$^{-1}$ for 40~kg (30~kg) fiducial mass (Table~\\ref{tab:summaryElectronRecoils}). By applying a veto cut with an average energy threshold of 100~keV, these rates are reduced to 6.1$\\times$10$^{-3}$ (4.6$\\times$10$^{-3}$)~events$\\cdot$kg$^{-1}\\cdot$day$^{-1}\\cdot$keV$^{-1}$ for 40~kg (30~kg) fiducial mass. The discrimination between electron and nuclear recoils based on the ratio of proportional to primary scintillation light (S2/S1) is not considered in this paper, and provides a further background reduction of~$>$99\\%. From the good agreement between Monte Carlo simulation and measured data, as shown in Figures~\\ref{figDataMC} and \\ref{figDataMCzoom}, and the predicted background rates from Table~\\ref{tab:summaryElectronRecoils}, it can be concluded that the electronic recoil background in the XENON100 experiment during the commissioning run in Fall 2009~\\cite{xe100-independent} is dominated by the natural radioactivity in the detector materials. With an optimized fiducial volume cut and an active veto cut, the background rate in the energy region of interest can be reduced down to a level where radioactive $^{85}$Kr in LXe starts to dominate. The results of the present work are not only important for understanding the electromagnetic background in the XENON100 experiment and the validation of the background model, but can be also useful for the design of next-generation detectors for dark matter searches, such as XENON1t or DARWIN~\\cite{darwin}." }, "1101/1101.3782_arXiv.txt": { "abstract": "{We integrate the 2D MHD ideal equations of a straight slab to simulate observational results associated with fundamental sausage trapped modes.} {Starting from a non--equilibrium state with a dense chromospheric layer, we analyse the evolution of the internal plasma dynamics of magnetic loops, subject to line--tying boundary conditions, and with the coronal parameters described in Asai et al. (2001) and Melnikov et al. (2002) to investigate the onset and damping of sausage modes.} {To integrate the equations we used a high resolution shock--capturing (HRSC) method specially designed to deal appropriately with flow discontinuities. }{Due to non--linearities and inhomogeneities, pure modes are difficult to sustain and always occur coupled among them so as to satisfy, e.g., the line--tying constraint. We found that, in one case, the resonant coupling of the sausage fundamental mode with a slow one results in a non--dissipative damping of the former. } {In scenarios of thick and dense loops, where the analytical theory predicts the existence of fundamental trapped sausage modes, the coupling of fast and slow quasi--periodic modes -with a node at the center of the longitudinal speed- occur contributing to the damping of the fast mode. If a discontinuity in the total pressure between the loop and the corona is assumed, a fundamental fast sausage transitory leaky regime is spontaneously produced and an external compressional Alfv\\'en wave takes away the magnetic energy.} ", "introduction": "Coronal heating remains today an open field of research. One of the most important contribution to it is attributed to active region overdense loops which are bright in soft X--ray and EUV (Aschwanden et al. \\citealp{ac10}). Over the past decade the interest in determining the presence of waves and oscillations that can contribute to the heating mechanisms has increased on the basis of the development of theoretical and observational studies. The coronal seismological remote diagnostics revealed an effective media to identify a wide--spectrum of fast and slow intensity oscillations. Measurements of characteristic periods, dispersion relations, speeds and damping times as well as different theoretical models that intend to describe non-dissipative damping mechanisms and leakage processes of systems with high Reynolds number provide new remote sensing diagnostic tools to reveal unknown or more accurate solar physics parameters, e.g., magnetic field strength and transport coefficients (see, e.g. Aschwanden \\citealp{asc4} and Nakariakov and Verwichte \\citealp{nak}, for recent reviews). The pure sausage mode, modelled as a magnetic cylinder or a magnetic slab, is known to be a compressible MHD fast mode that perturb the plasma in the radial direction causing a symmetric contraction and widening of the tube without distortion of its axis (Edwin and Roberts \\citealp{rob}). This mode was observationally detected in the corona by Nakariakov et al. \\cite{nak2} using the Nobeyama Radioheliograph. EUV and soft X--ray band periodic variations of the thermal emission intensity are associated with density perturbations and Doppler broadening of the emission lines. Zaitsev and Stepanov \\cite{zai} proposed that a modulation of hard X--ray and white--light emission from the loop footpoints can occur due to the change in the loop radius and the consequent change of the mirror ratio, causing the periodic precipitation of non--thermal electrons of flaring loops. These quasi--periodic pulsations, triggered by sausage modes, were also associated with flaring events of other stars (Mitra--Kraev et al. \\citealp{mit}). The standard theoretical models (Edwin and Roberts \\citealp{rob}; \\citealp{rob2}; Zaitsev and Stepanov \\citealp{zai}) predict that the fast global sausage mode has two regimes separated by a cutoff at the long wavenumbers: the leaky one, where the energy is lost through the boundaries of the loop and carried away by an external solution, occurring in thin and long loops when the internal Alfv\\'en speed is larger than the external one; and the trapped regime, occurring in thick and short loops, where the energy of the oscillation is confined to the interior of the loop structure. Aschwanden et al. \\cite{asc} showed that the sausage cutoff imposes such a high electron density that they can only occur in flare loops. They found that previously reported fast global mode observations had oscillations confined to a small segment of the loop with nodes limiting the segment, indicating that the oscillation must be associated with higher harmonic modes rather than fundamental ones. However, they found that microwave and soft X--ray observations by Asai et al. \\cite{asa} and Melnikov et al. \\cite{mel} are consistent with fast sausage MHD oscillations at the fundamental harmonic. The leaky regime has recently been studied by Pascoe et al. \\cite{pas}, who found by numerical modeling that long loops with sufficiently small density contrast can support global sausage leaky modes of detectable quality, and thus they argue that this mode can also be responsible of the quasi--periodic pulsations of flaring emission. An important issue is that the global sausage leaky mode period is determined by the Alfv\\'en speed outside the loop. Also, Verwichte et al. \\cite{ver} noted that wave energy of loops with internal Alfv\\'en speed locally below the external one can tunnel through an evanescent barrier into the surrounding corona due to the inhomogeneity of the media that produces relative variations of the Alfv\\'en speeds. Recently, Srivastava et al. \\cite{sri} based on the observation of multiple sausage oscillations suggested that fundamental post flare cool loops are not in equilibrium causing plasma motion along the loop. Pascoe et al. \\cite{pas2} investigate the effects of non--uniform cross--sections and found that the global sausage mode is very effectively excited in flaring loops. The sound $(v_{s})$ and Alfv\\'en $(v_{A})$ speeds are associated with two main properties of the low corona, its compressibility and elasticity, respectively. The low coronal values of the parameter $\\beta \\simeq v_{s}^{2}/v_{A}^{2}<<1,$ imply that the sound speed is much smaller than the Alfv\\'en one. Due to the high Reynolds number in the corona and to the rapid fast mode decays, this must be firstly accomplished in an almost non--dissipative way. When more realistic inhomogeneous distributions of the physical quantities are taken into account, both in the longitudinal and radial direction (e.g., dense chromospheric layer), the pure character of the MHD modes is lost and the nonlinear coupling of modes takes place. In mediums with usual strong sudden energy release sources, the drain of the larger Alfv\\'enic elastic energy to the lower magnetoacoustic compressive perturbations would not be unusual, leading to the eventual non--linear steepening of modes and to the formation of coronal shocks (Costa et al. \\citealp{cos}). The quantitative evaluation of the statistical distribution and energy content in shock waves is important to estimate its contribution as a source for the coronal heating. In this paper our aim is to investigate the onset and damping of coronal sausage modes considering nonlinear and inhomogeneous initial states. We simulate two studied observational flaring loops supposed capable to develop global sausage trapped modes, described in Asai et al. \\cite{asa} (Case A) and Melnikov et al. \\cite{mel} (Case M), respectively. \\begin{figure}% \\begin{center} \\includegraphics[width=7.4cm]{f0.ps} \\end{center} \\caption{Divergence cleaning test. Dimensional time obtained multiplying by $t_{0}=1.39$sec for Case A and $t_{0}=3.25$sec for Case M.} \\label{fig:cero} \\end{figure} \\begin{figure*}% \\begin{center} \\includegraphics[width=7.4cm]{f1A.ps} \\hspace*{19pt} \\includegraphics[width=7.4cm]{f1M.ps} \\end{center} \\caption{Density as a function of time, across the loop width ($x$), and at the center of the loop length, ($y=0$). Left) Case A, dimensional values obtained multiplying by $l_{0}=10$Mm and $t_{0}=1.39$sec; Right) Case M, dimensional values obtained multiplying by $l_{0}=16.3$Mm and $t_{0}=3.25$sec. } \\label{fig:uno} \\end{figure*} \\begin{table*} \\begin{center} \\begin{tabular}{|l|c|c|c|c|c|c|} \\hline \\multicolumn{1}{|c|}{\\bf Regions} & \\multicolumn{1}{c|}{$n_{A}$[$10^{10}$ \\ $cm^{-3}$]\\rule{0pt}{9pt}} & \\multicolumn{1}{c|}{$T_{A}$[$MK$]}& \\multicolumn{1}{c|}{$B_{A}$[$G$]}& \\multicolumn{1}{c|}{$n_{M}$[$10^{10} \\ cm^{-3}$]} & \\multicolumn{1}{c|}{$T_{M}$[$MK$]} & \\multicolumn{1}{c|}{$B_{M}$[$G$]} \\\\ \\hline Inner ($0$\\rule{0pt}{9pt}) & $4,3$ & $2.4 $ & $152$ & $9.8$ & $9$ & $90$ \\\\ \\hline External ($e$\\rule{0pt}{9pt}) & $0.21$ & $2.5 $ & $153$ & $0.16$ & $8.8$ & $104$ \\\\ \\hline Chromosphere ($c$\\rule{0pt}{9pt}) & $19000$ & $0.02 $ & $106$ & $1900$ & $0.1$ & $72$ \\\\ \\hline \\end{tabular} \\end{center} \\caption{\\label{tab:table1} Initial parameter values: particle number density $n,$ temperature $T,$ magnetic field $B$. Left) Case A values taken from Asai et al. \\cite{asa}; Right) Case M values taken from Melnikov et al. \\cite{mel}. } \\end{table*} ", "conclusions": "We integrated the ideal MHD equations to simulate fundamental trapped sausage modes observationally described (Asai et al. \\citealp{asa} (Case A) and Melnikov et al. \\citealp{mel} (Case M)) incorporating a dense chromospheric region. We could reproduce the observational parameters described in the literature, i.e, periods, densities, lengths and magnetic fields. As in Melnikov et al. \\cite{mel}, we found that, for both cases, there are two peak fast frequencies with different contributions depending on the location of the signal. The modes were interpreted as the fundamental (more intense at the top of the loop) and the second sausage harmonic (more intense at the footpoints) (see Fig.~\\ref{fig:dos}a for Case A and Fig.~\\ref{fig:dos}b and Fig.~\\ref{fig:cuatro}a for Case M). We analyzed the coupling of the fundamental sausage modes with longitudinal components. In both cases we obtained slow frequency components, of $P\\sim 17$sec and $P\\sim 54$sec, respectively. Moreover, we found that Case M has another slow contribution, of the same frequency as the fundamental sausage mode, apparently driven by this frequency. The slow component of Case A is rapidly damped ($\\sim 5$ periods); whereas the -more broad slow spectra- of Case M lasts not damped for various periods. We suggest that this is due to the internal transferring of the fast mode energy into the slow energy which is accomplished more efficiently in the resonant case. Thus, we call this mechanism an internal damping one. \\\\ \\indent We also showed that certain initial conditions associated with typical flaring loop energy releases (impulsive depositions of energy) that could be achieved by considering jump conditions across the radio, results in the drain of energy into the exterior in the form of a leaky compressible Alfv\\'en mode, allowing the final damping of the initially trapped fundamental sausage mode. This external damping mechanism which is efficient to damp modes in an initially trapped configuration is in accordance with typical observational damping times of this modes ($\\leqslant 10$periods)." }, "1101/1101.1787_arXiv.txt": { "abstract": "{From an optical spectroscopic survey of 3CR radio galaxies with $z<0.3$, we discovered a new spectroscopic class of powerful radio-loud AGN. The defining characteristics of these galaxies are that compared with radio galaxies of similar radio luminosity they have: a \\oiiihb\\ ratio of $\\sim$ 0.5, indicative of an extremely low level of gas excitation; a large deficit of \\oiii\\ emission and radio core power. We interpret these objects as relic AGN, i.e. sources that experienced a large drop in their level of nuclear activity, causing a decrease in their nuclear and line luminosity. This class opens a novel approach to investigating lifetimes and duty cycles of AGN.} ", "introduction": "\\label{introduction} Optical spectroscopy has played a major role in enhancing our understanding of active galactic nuclei (AGN). \\citet{heckman80} and \\citet{baldwin81} demonstrated that optical lines can be used as tools to classify in general emission-line objects, particularly AGN. Diagnostic diagrams comparing emission line ratios can distinguish H~II regions from gas clouds ionized by nuclear activity \\citep{veilleux87}. Moreover, AGN can be separated into Seyferts and Low Ionization Nuclear Emission-line Regions \\citep[LINERs,][]{heckman80} since they form separate branches in the diagnostic diagrams \\citep{kewley06}. We performed an optical spectroscopic survey of the 113 radio galaxies (RG) belonging to the 3CR sample and with $z<$ 0.3 \\citep{spinrad85}, using the Telescopio Nazionale Galileo \\citep{buttiglione09,buttiglione11}. Most RGs belong to two main spectroscopic classes, those of high and low excitation galaxies (HEG and LEG respectively, originally introduced by \\citealt{laing94}, the analogous to Seyfert and LINER for radio-loud AGN). In \\citet{buttiglione10}, we also reported the discovery of a new class, characterized by an extremely low level of gas excitation whose properties are discussed in this Letter. We adopt the following cosmological parameters: $H_o = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.73$, and $\\Omega_m = 0.27$. ", "conclusions": "\\label{summary} The energy carried by the jets of radio-loud AGN has a profound impact on the evolution of their hosts and the energy balance of the intracluster medium (e.g. \\citealt{fabian03,croton06}). To ascertain the effects of AGN feedback, it is necessary to measure not only the jet power, but also the timescales and frequency of the energy release, i.e. measure AGN lifetimes and duty cycles. Several approaches, based purely on radio properties, have been used previously (e.g. \\citealt{alexander87,shabala08,bird08}) to investigate these issues. These analyses have led to evidence \\citep{giovannini88,parma07} of relic (fossil) RG based on the low core dominance or very steep spectral indices (a manifestation of the energy losses of the radio-emitting relativistic electrons) of some systems. Using radio observations it is possible to derive model-dependent estimates of the RG lifetimes or, in case of fossil RG, the time elapsed since the `death' of the active nucleus, $\\tau_{\\rm D}$. Relic RGs, distinguished by their optical spectroscopic properties, open a new and complementary path to exploring the lifetime and duty cycles of radio galaxies. The fraction of relics in the RG population is given by the ratio of the duration of the relic phase to the RG lifetime. In principle, both parameters can be derived once all RG in a complete sample are spectroscopically classified, a situation that is (almost) met by our 3CR spectroscopic survey. Considering the range of extended luminosity where relic RG are observed, Log L$_{178} > 33$ \\ergsHz, there are 83 objects in the 3CR with $z<0.3$, three of which are relics, one has a star-forming-like spectrum, 46 are HEG, and 20 are LEG. This leads to the fraction of relics being 3/83. This fraction can increase substantially depending on the nature of the 13 objects in the sample that could not be classified spectroscopically. In particular, as explained in more detail below, some of them can be considered as `candidate relics'. A second, less important, uncertainty in the relic fraction would also exist if HEG only can evolve into relic systems. The duration of the relic phase could in principle be estimated using the spectroscopic evolution, by comparing the temporal changes in line ratios with those predicted by photoionization models. Furthermore, the intensity of the various lines can be compared with those of RG of similar extended luminosity, i.e. with the relic progenitors. This approach requires a direct measurement of the NLR density, the parameter that drives its temporal evolution. Unfortunately, the density in the relic galaxies is essentially unconstrained by our observations as we are unable to measure the [S~II] doublet ratio in two of the ELEG, and it is poorly constrained for 3C~314.1. However, future measurements should be easily attainable with dedicated observations as the intensity of the [S~II] doublet is similar to that of the \\Ha\\ line, whose flux was accurately measured in our spectra. Although the extended radio emission is, in this context, the structure responding on the longer timescale to the changes in the nuclear properties, the radio morphology of the three relics can provide useful insights into their evolution. All of them can be classified as FR~II following the original definition of \\citet{fanaroff74}, but they differ when examined in more detail: -- 3C~028 has well-defined twin jets linking the host to a double-lobed structure, with well-defined hot spots, but lacks a radio core \\citep{feretti84}. There is no apparent brightness discontinuity in the jets on either side of the host, an indication that the drop in nuclear activity occurred recently. -- 3C~314.1 has a structure known as a ``fat double'' as it lacks jets and hot spots \\citep{leahy91}. This is an indication that the switching-off of the jets of this galaxy occurred at least $\\sim 7 \\times 10^5$ years ago, based on the light travel time to the edge of the radio source $\\sim$ 250 kpc from the nucleus. -- 3C~348 is considered to be a ``born again\" RG based on its peculiar radio morphology \\citep{gizani03}: two jets emerge from the core and propagate within a relaxed double-lobed, steep-spectrum structure. On the basis of the size of the inner radio structure embedded in the FR~II fossil ($\\sim$ 200 kpc), the relic phase started at least $3\\times 10^5$ years ago. Despite the presence of a restarted AGN, the relic line emission appears to be still the dominant component observed in this galaxy; this implies that the current phase must be associated with a significantly lower level of nuclear luminosity than the earlier phase of activity. Evidence of the different evolutionary stages in the three sources is supported by comparing their \\oiii\\ luminosity with respect to HEG of similar radio luminosity (see Fig. \\ref{zoom}): for 3C~028 this is lower by only a factor of $\\lesssim$ 10, while 3C~314.1 and 3C~348 exhibit much larger deficits (a factor of 10$^2$ - 10$^3$) . An independent estimate of the relic phase duration, $\\tau_{\\rm D}$, can be obtained by {\\sl directly} spatially mapping the change in state of the AGN using deeper optical spectroscopy of the off-nuclear emission line regions (e.g. \\citealt{robinson87}). Owing to light travel effects, information about the drop in the nuclear emission might not have reached these outer regions. In a sort of reversed light-echo effect, the regions located at radii $r > c \\, \\tau_{\\rm D}$ would still show the original high excitation state, characterized by a ratio \\oiiihb\\ $\\sim 20$ times higher than in the nuclear regions. In practice, geometrical effects affect the estimate of $\\tau_{\\rm D}$, since the isochrones are paraboloids centered on the nucleus and the derived timescale depends on the orientation of the gas cloud with respect to the line of sight. This degeneracy can be broken when emission lines are detected on both sides of the nucleus. Finally, we consider in more detail the 13 spectroscopically unclassified sources (usually because the \\oiii\\ and/or the H$\\beta$ lines are too faint to be measured). Not all of them should be considered as potential additional relics since the uncertain classification is caused by a variety of factors\\footnote{ For example, 3C~111 and 3C~445 are broad line RG, not classified because their broad Balmer lines hides completely the narrow components; 3C~132 is seen through a region of very high galactic absorption ($A_V \\sim 4$); in 3C~346 the \\Ha\\ line coincides with a telluric band.}. However, relative to RG of similar total radio luminosity nine unclassified sources also have a relatively low \\oiii\\ luminosity (see Fig. \\ref{zoom}) and we therefore consider them to be plausible `candidate relics'. The importance of these objects is two-fold: first of all, only after a proper spectroscopic validation will it be possible to derive the ratio of the number of relic to active galaxies, and consequently the relative duration of the two phases; secondly, it is possible that these sources might be in a different evolutionary stage than the three relic galaxies discussed here. A clearer characterization of the class of relic RG would require a study of these objects in greater detail from both the radio and spectroscopic point of view. Relics AGN are also likely to exist among radio-quiet objects; for example, NGC~5252, with its high-excitation extended NLR surrounding a nuclear region with a LINER spectrum, is a likely relic QSO \\citep{goncalves98,capetti5252}. In the region of the spectroscopic diagrams typical of the 3CR relics there is indeed a substantial number of SDSS emission-line galaxies (see Fig. 1). However, the lack of radio diagnostics makes it identify genuine relics among the broad distribution of radio-quiet AGN in the diagnostic planes. Furthermore, relics are likely to be rare in a flux-limited sample of emission line galaxies, particularly when their definition requires the detection of the short-lived \\oiii\\ line. As an alternative, one could rely on non-classical diagnostic diagrams, including lines originating from high excitation gas unaffected by the charge exchange reaction (e.g. HeII$\\lambda$4686) that causes the extremely short-decay time for the [O~III] line. The temporal evolution within these same diagnostic diagrams can provide us with an estimate of the relic age. However, the likelihood of observing these effects will in general depend on both the gas density and (through light travel times) the size of the emitting region. In principle, extrapolating backward in time it will be possible to associate these sources with their active counterparts at the correct level of luminosity and thus estimate their lifetimes." }, "1101/1101.1114_arXiv.txt": { "abstract": "We investigate the X-ray absorption structure of oxygen in the interstellar medium by analyzing {\\it XMM}-Newton observations of the low-mass X-ray binary Sco X-1. We use simple models based on the O~{\\sc i} atomic photoabsorption cross section from different sources to fit the data and evaluate the impact of the atomic data on the interpretation of the observations. We show that relatively small differences in the atomic calculations can yield spurious results, and that the most complete and accurate set of atomic cross sections successfully reproduce the observed data in the $21.0{-}24.5$~{\\AA} wavelength region of the spectrum. Our fits indicate that the absorption is mainly due to neutral gas with an ionization parameter of $\\xi=10^{-4}$~erg~cm~s$^{-1}$ and an oxygen column density of $N_{\\mathrm{O}}\\approx (8{-}10)\\times 10^{17}$~cm$^{-2}$. Our models are able to reproduce both the K edge and the K$\\alpha$ absorption line from O~{\\sc i} which are the two main features in this region. We find no conclusive evidence for absorption by other than atomic oxygen. ", "introduction": "X-ray spectroscopy provides a powerful tool for understanding the physical and chemical properties of the diffuse interstellar medium (ISM). The X-ray band covers the emission and absorption spectra produced by inner-shell transitions of the most abundant ions from carbon to iron. The interaction of X-ray photons from bright point sources (e.g. galactic X-ray binaries) with the ISM gives rise to absorption lines and edges in the spectrum. The energy position and the shape of these features depend on whether the absorption is due to free atoms or molecules, and on whether these atoms or molecules are in the gas or in the solid phase. Neutral oxygen is a major constituent of the ISM which makes it one of the most important elements in astronomical observations. Precise knowledge of the neutral oxygen atomic parameters is needed for the correct modeling of the observed spectra. Calculations of the photoabsorption cross section of the ground state of O~{\\sc i} were carried out by \\cite{mcl98} (hereafter {\\sc mck98}) using the {\\it R}-matrix method, giving a detailed comparison with the experimental results of \\cite{sto97}. Although they claimed overall agreement, there are significant discrepancies in the positions of the inner-shell excited resonances and in the near-threshold resonance profiles. This problem was overcome in the $LS$-coupling calculation of \\cite{gor00} (hereafter {\\sc gmc00}) by taking into account core relaxation effects and the smearing of the K edge due to Auger damping. A more complete {\\it R}-matrix calculation was carried out in intermediate coupling by \\cite{gar05} (hereafter {\\sc gar05}) for all the ions in the oxygen isonuclear sequence. There is now very good agreement between the {\\sc gar05} calculations and both the experimental cross section \\citep{sto97} and the {\\sc gmc00} results. The oxygen inner-shell features in the X-ray spectrum of galactic sources have been used to provide abundance determinations in the ISM as well as estimates of the oxygen ionization fractions \\citep{sch02,tak02,jue04,tur04,ued05,jue06}. However, studies in the IR and UV have shown that oxygen can also be found in solid particles \\citep{dra03,whi03}. It has been argued that oscillatory modulations near the K edge, usually referred to as the X-ray absorption fine structure (XAFS), could be detected. These are condensed matter modulations of the atomic cross section due to the presence of solid particles \\citep{lee05,lee09}. Studies of the soft X-rays from galactic sources have reported possible detections of molecules that could be linked to XAFS signatures in the edges of several elements such as Ne, Si, and Mg \\citep{pae01,lee02,ued05}; in the L edge of Fe \\citep{lee01,kas09}; and in particular, in the oxygen K edge \\citep{dev03,cos05,dev09,pin10}. Although these signatures could be important in the K edges of molecules involving higher $Z$ elements, namely Mg, Si, and Fe, oxygen is $\\sim 10{-}20$ times more abundant than any of these, thus potentially providing enough signal-to-noise to detect the XAFS signatures. Theoretical models of XAFS in the astrophysical context have been developed by \\cite{woo95,woo97}; and \\cite{for98}. See also \\cite{lee11} and references therein for details on the theory of XAFS. {\\it XMM}-Newton observations of the X-ray source Scorpius~X-1 (Sco~X-1) reveal strong absorption in the wavelength region corresponding to neutral oxygen. Located at a $\\sim 2.8$~kpc distance \\citep{bra99} and with a flux of $F\\sim 3.4\\times 10^8$~erg~cm$^{-2}$~s$^{-1}$ (in the $2{-}10$~keV energy band), it is the brightest X-ray source in the sky other than the Sun and the diffuse X-ray background radiation. Its high X-ray flux provides very good statistics in relatively short exposure times, giving the opportunity to study signatures of oxygen absorption in the ISM with great detail. For several galactic sources, including Sco~X-1, \\cite{devr03} analyzed high-resolution X-ray spectra taken with the reflection grating spectrometer (RGS) in the {\\it XMM}-Newton satellite. By comparing low and high extinction sources, they were able to separate the ISM and the instrumental components of the O~{\\sc i} K edge; moreover, \\cite{dev09} searched for XAFS signatures in the spectrum of Sco~X-1. The XAFS signature is derived from the differences between the observed flux and that predicted theoretically. However, the model used by these authors is based on the atomic oxygen absorption cross section calculated by {\\sc mck98}. In this Letter we show the importance of the accuracy of the atomic data used in the modeling of the detailed features of the oxygen absorption in the ISM. We demonstrate that small variations in the K-edge structure derived from different atomic calculations yield spurious results when applied to astronomical observations. In Section~\\ref{secobs}, we describe the observational data used in our study while the theoretical models are delineated in Section~\\ref{secmod}. Results derived from model fits of the observed data are presented in Section~\\ref{secres}, and finally, the main conclusions are summarized in Section~\\ref{seccon}. ", "conclusions": "\\label{seccon} In the present report we have shown the relevance of accurate atomic data in the modeling of the X-ray spectra from cosmic sources. In particular, we have studied the Sco X-1 spectrum produced by the RGS1 instrument on board of the {\\it XMM} Newton satellite covering the $21.0{-}24.5$~{\\AA} wavelength region. Absorption occurs when the X-rays interact with the cold gas of the ISM, the main spectral features in this region being the absorption K edge and K$\\alpha$ line from neutral oxygen. We found a good fit using a self-consistent photoionization model which includes the most recent atomic data for the oxygen isonuclear sequence by {\\sc gar05}. Our fits indicate that the absorbing gas has an ionization parameter of $\\xi=10^{-4}$~erg~cm~s$^{-1}$ and a hydrogen column density of $N_{\\mathrm{O}}\\approx (8{-}10)\\times 10^{17}$~cm$^{-2}$. Simple models based on the raw atomic photoabsorption cross sections of O~{\\sc i} from three different calculations were used to evaluate data sensitivity. We show that models based on the {\\sc mck98} atomic cross sections are unable to reproduce the K-edge structure in detail, while those based on {\\sc gmc00} and {\\sc gar05} yield more accurate fits of the main spectral features. The fits using the most up to date models do not show evidence for absorption by anything other than atomic oxygen. The analysis presented here indicates that the atomic data uncertainties in combination with the limited resolution of the grating spectrum make detection of molecular or solid material challenging. Although oxygen is expected to be found in molecular form or locked into solids in the ISM, the use of accurate atomic calculations to correctly account for the atomic oxygen contribution is crucial when searching for XAFS or similar features in the X-ray spectra of astronomical sources." }, "1101/1101.1322_arXiv.txt": { "abstract": "{ We present cosmological simulations of galaxy clusters, with focus on the cluster outskirts. We show that large-scale cosmic accretion and mergers produce significant internal gas motions and inhomogeneous gas distribution (\"clumpiness\") in the intracluster medium (ICM) and introduce biases in measurements of the ICM profiles and the cluster mass. We also show that non-thermal pressure provided by the gas motions is one of the dominant sources of theoretical uncertainties in cosmic microwave background secondary anisotropies. We briefly discuss implications for cluster cosmology and future prospects for understanding the physics of cluster outskirts using computer simulations and multi-wavelength cluster surveys. } ", "introduction": "In recent years, galaxy clusters have emerged as one of the most unique and powerful laboratories for cosmology and astrophysics. Being the largest and most magnificent structures in the Universe, clusters of galaxies serve as excellent tracers of the growth of cosmic structures. Current generation of X-ray cluster surveys have provided independent confirmation of cosmic acceleration and significantly tighten constraints on the nature of dark energy \\citep{allen_etal08,vikhlinin_etal09} and alternative theories of gravity \\citep[e.g.,][]{schmidt_etal09}. Several ongoing and new X-ray (e.g., {\\it eROSITA}) and Sunyaev-Zel'dovich effect (SZE) cluster surveys (e.g., {\\it SPT, ACT, Planck}) are underway to improve current cosmological constraints. Outskirts of galaxy clusters have special importance for cluster cosmology, because they are believed to be much less susceptible to complicated cluster astrophysics, such as radiative gas cooling, star formation, and energy injection from active galactic nuclei. Dominant physical processes in the outskirts are limited to the gravity-driven collisionless dynamics of dark matter and hydrodynamics of the intracluster medium (ICM). In the hierarchical structure formation model, galaxy clusters grow by accreting clumps and diffuse gas from the surrounding large-scale structure in their outer envelope. Numerical simulations predict that the large-scale cosmic accretion and mergers give rise to internal gas motions and inhomogeneous gas distribution in the ICM. However, until very recently, observational studies of the ICM have been limited to radii considerably smaller than the virial radius of clusters. Recently, {\\it Suzaku} X-ray observations have extended X-ray measurements of the ICM profile out to and beyond the virial radius for several clusters \\citep{bautz_etal09,george_etal09,reiprich_etal09,hoshino_etal10,kawaharada_etal10}. While these measurements are still quite uncertain, initial results suggested that the observed ICM profiles may deviate significantly from the prediction of hydrodynamical cluster simulations \\citep[e.g.,][]{george_etal09}. In addition to testing models of structure formation, these new measurements will be important for controlling systematic uncertainties in cluster-based cosmological measurements. In this work, I will present theoretical modeling of the outskirts of galaxy clusters based on cosmological simulations, with highlights on implications for the interpretation of forthcoming multi-wavelength observations of galaxy clusters. The simulations we present here are described in \\citet{nagai_etal07a} and \\citet{nagai_etal07b}, and we refer the readers to these papers for more details. \\begin{figure}[t] \\begin{center} \\vspace{-4mm} \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{f1.eps}} \\vspace{-8mm} \\caption{\\footnotesize {\\it Top panel:} Ratio of pressure from random gas motions to total pressure as a function of radius. Relaxed clusters are represented by solid lines while unrelaxed clusters are represented by dashed lines. {\\it Bottom panel:} Averaged mass profiles $M(< r)$ of the relaxed clusters, normalized by $M_{500}$. The solid line shows the actual mass profile from simulation, the long dashed line shows the mass profile from hydrostatic equilibrium including random gas and thermal pressure, and the short dashed line shows the mass profile from hydrostatic equilibrium including thermal pressure only. Hashed region shows the 1-$\\sigma$ error of the mean. From \\citet{lau_etal09}.} \\vspace{-5mm} \\label{fig:clump_phase} \\end{center} \\end{figure} ", "conclusions": "" }, "1101/1101.3327_arXiv.txt": { "abstract": "Certain thermal non-equilibrium situations, outside of the astrophysical realm, suggest that entropy production extrema, instead of entropy extrema, are related to stationary states. In an effort to better understand the evolution of collisionless self-gravitating systems, we investigate the role of entropy production and develop expressions for the entropy production rate in two particular statistical families that describe self-gravitating systems. From these entropy production descriptions, we derive the requirements for extremizing the entropy production rate in terms of specific forms for the relaxation function in the Boltzmann equation. We discuss some implications of these relaxation functions and point to future work that will further test this novel thermodynamic viewpoint of collisionless relaxation. ", "introduction": "Galaxy-hosting dark matter systems and the galactic stellar systems themselves act collisionlessly over Hubble timescales. Understanding the evolution of self-gravitating, collisionless systems is a foundational element to the larger picture of galaxy evolution. Much of the advancement in understanding collisionless evolution has come from $N$-body simulations that focus on the motions of large numbers of individual mass elements. Such simulations allow astrophysicists to determine the density and velocity distributions that result from a variety of initial conditions \\citep[\\eg][]{va82,ma85,nfw96,m98}. While $N$-body simulations have given the astrophysics community a powerful method for making predictions that can be tested against observations, the modeled evolutions are sufficiently complex that a full physical picture of the evolution remains lacking. Effects such as dynamical friction, collisionless relaxation processes, and even numerical artifacts can all play roles of varying importance in these simulations. Astrophysicists now have substantial empirical evidence that collisionless systems in cosmological contexts have ``universal'' equilibrium distributions of mass and velocities, \\eg\\ the radial power-law behavior of density divided by velocity dispersion cubed \\citep{tn01}. We are interested in gaining a better understanding of the physical origin of these seemingly special distributions. Before beginning to build a physical picture of collisionless evolution, let us quickly review some of the ideas that are involved. Thermal equilibrium is the state that systems relax towards when energy can be exchanged between the components of the system. A system in thermal equilibrium is characterized by a minimum internal energy or, equivalently, a maximum entropy. In a thermodynamic sense, entropy is linked to the amount of energy transferred as heat, bringing with it a connotation of randomness. In a statistical sense, entropy is simply a redefinition of the accounting of energy states in a system. Relaxation is a generic term for any process that erases the memory of initial conditions in a system. For example, gases reach the completely relaxed state of thermal equilibrium through collisions. \\citet{lb67} demonstrated that non-degenerate collisionless self-gravitating systems with finite mass and energy do not have a state of maximum entropy when one works with the standard distribution function and assumes that the occupation numbers of phase-space volumes is large enough that Stirling's approximation applies \\citep[For an alternative view see][]{hw10}. This lack of entropy extremum implies that there is no thermal equilibrium state for such systems. And yet, mechanical equilibrium, a state with no net force at any location, is possible. In stellar interiors, for example, the lack of thermal equilibrium---evidenced by gradients in temperature and pressure---is required for mechanical equilibrium. Note that the thermal non-equilibrium description refers to the entire star. It is common to treat the gas in a star as being in local thermodynamic equilibrium, where small regions of the star are modeled as having uniform temperatures. In this work, we consciously make a distinction between full thermal equilibrium and stationary (time-independent) situations. An example that clarifies this difference is familiar to those of us nearer the poles. On a cold winter day, a house is heated to a temperature that is well above the outside air temperature. Heat flows out of the house, and entropy is generated in a wall during the energy transfer. As long as the home's heating system runs and the outdoor temperature remains constant, this process will be steady-state, yet it does not represent a thermal equilibrium as the entropy of the home-outside air system continually increases. Dark matter halos and galactic stellar systems are similar to the above, in the sense that they have temperature gradients and are thus thermally non-equilibrium systems, even though they are steady-state and are in mechanical equilibrium. Thermal non-equilibrium systems are common in nature and have been the subject of numerous studies. For certain non-astrophysical, non-gravitating systems it was found that stationary states, like mechanical equilibrium, coincide with states of extreme entropy production \\citep[][and references therein]{dgm84,g08}. In some instances, it appears that steady-states are reached by minimizing entropy production, while in other cases maximizing entropy production is the requirement. The most common modern formulation of the principle of maximum entropy production is due to \\citet{z61}. The principle of miminum entropy production \\citep{p78} appears to have a smaller range of applicability, but does find uses in physics, chemistry and biology. The theoretical interpretation, and the relation between the two seemingly contradictory principles is still being worked out \\citep[\\eg,][]{ms06,b10}. A further, and more general reason to explore entropy production in self-gravitating systems is because irreversability and entropy production are closely related concepts. The ``microscopic'' equation of motion of gravity (and of other microscopic physics) is time-reversible, while the global evolution of the system is not, in the sense that it has a definite arrow of time. A natural agent that can make the transition from reversability to irreversability in macroscopic systems is the thermodynamic entropy production. This ability of thermodynamics to pick the arrow of time through non-zero (positive) entropy production has been utilized in a cosmological context by \\citet{pg86}. These authors propose a scenario where in the early Universe some of the energy of the space-time (gravity) was transferred to matter. Because matter production is an irreversible, entropy creating process, the evolution of the Universe proceeds with an arrow of time. In this paper we investigate the consequences of applying the entropy production extremization principle to self-gravitating systems. We do not prove that such a principle should be applicable here; we merely explore what this principle implies. In particular, we want to find what conditions are necessary to extremize entropy production and how those conditions differ from maximizing the total entropy. We begin with a review of relevant ideas regarding entropy from both a thermodynamic and statistical point of view in Section~\\ref{back}. In \\S~\\ref{kinetic}, entropy production equations for the two statistical families we are interested in are then derived. For these cases, we also develop the specific forms for collisionless relaxation processes required to extremize the entropy production. We present a discussion of our results and conclude in Section~\\ref{discuss}. ", "conclusions": "\\label{discuss} The seminal paper by \\cite{lb67} has shown that, under certain assumptions, there is no maximum entropy state for self-gravitating systems. Alternative descriptions of the same systems which yield self-consistent results under entropy extremization may be found in \\citet{m96} and \\citet{hw10}. Lynden-Bell's conclusion is exemplified by the order-of-magnitude calculation in \\citet[][S.\\ 4.7]{t86,bt87}, using Maxwell-Boltzmann statistics. The result of this calculation is that when a collisionless system rearranges its mass to become more centrally concentrated in the core, the entropy of the contracting core mass decreases. The outer envelope, expanding in response to the contracting core to satisfy the virial theorem, should have an increase in entropy. As a result, the entropy for the entire system increases. Since a system can always increase its entropy by contracting the core and expanding the envelope, there is no maximum entropy state, and hence maximizing entropy will not lead to a satisfactory description of a steady state. Yet, we know from numerous high-resolution $N$-body simulations that long-lived steady states do exist. How does one find these theoretically? Apparently, one has to resort to means other other than entropy extremization. In this paper we try one alternative approach. We apply a principle of extremizing entropy production rate to self-gravitating systems. This principle has been used widely to describe thermal non-equilibrium, but not in systems that are self-gravitating. Our basic hypothesis is that a steady state is obtained by extremizing the entropy production. We present expressions for the entropy production rates for two types of statistics, Maxwell-Boltzmann (MB) and Lynden-Bell (LB), as Equations~\\ref{mbsprod} and \\ref{lbsprod}, respectively. We then find expressions (Equations~\\ref{mbexts} and \\ref{lbexts}) for the relaxation term that forms the right hand side of the coarse-grained Boltzmann equation. The meaning of these expressions and the interpretation of our results, under the assumption that this idea is applicable to self-gravitating systems, are discussed below. \\subsection{Entropy production} The development of the expressions for entropy production rates $\\sigma$ is a central result of this work. The descriptions of $\\sigma$ for the MB and LB statistical families are given in Equations~\\ref{mbsprod} and \\ref{lbsprod}, respectively. Since we are dealing with collisionless systems exclusively, one might expect these to be zero. In fact, if we were considering the fine-grained distribution function, there would be no entropy change, no thermodynamic evolution, as the collisionless $\\dif f/\\dif t=0$ Boltzmann equation would apply. However, we are considering the coarse-grained function. As a system evolves, the fine-grained distribution function becomes stretched and twisted in phase-space. Because the coarse-grained function averages the fine-grained function with nearby empty regions of phase-space, the coarse-grained function changes as the system evolves. Now, recall that entropy represents the number of accessible states. On the level of the fine-grained function, the number of accessible states stays the same. However, going from a fine-grained to coarse-grained description implies that there are now regions of phase-space not occupied by the fine-grained function that are accessible to the coarse-grained function. This implies that there are more microscopic ways of realizing a given macroscopic state, leading to more possible states and larger entropy. Thus, coarse-graining an evolving system results in entropy production even in a collisionless system. In terms of physical processes, the evolution is due to the larger-scale phase-space evolution of the system driven by collisionless relaxation processes, like violent relaxation and phase mixing. The above argues that entropy production takes places during evolution of collisionless systems. But our analysis shows that entropy production takes place even during the steady state. Let us start by writing down the expression for the entropy production during the steady-state by combining Equations~\\ref{mbsprod} and \\ref{mbexts} for the MB case, \\begin{eqnarray}\\label{siglast} \\sigma_{\\rm MB} & = & -\\kb \\int Q \\Gamma_{\\rm MB}(\\eta) \\: \\dif\\vv \\nonumber \\\\ & = & -\\kb Q\\Gamma_{\\rm MB}(f=\\eta)V_{\\rm velocity}, \\end{eqnarray} where $V_{\\rm velocity}$ is the volume of occupied velocity space. This is consistent with the second law of thermodynamics since $Q<0$ and all other terms are positive. A similar expression can be found for the LB case, using Equations~\\ref{lbsprod} and \\ref{lbexts}. Since a steady state is described by an unchanging value of $\\sigma$, any non-zero value of $\\sigma$ persists even when a system has reached mechanical equilibrium. It is interesting to think of the source of this continued entropy production. After a system stops evolving on the macroscopic scale, it still continues to evolve on ever decreasing microscopic scales as the fine-grained function continues to stretch and twist almost indefinitely. The corresponding continual coarse-graining of the ever evolving fine-grained function on smaller and smaller scales, results in constant, non-zero entropy production. \\subsection{Interpreting the Boltzmann Equation} Our expression for the Boltzmann equation states that the relaxation function, $\\Gamma$, determines the rate of change of the coarse-grained distribution function, \\begin{equation} \\frac{\\dif f}{\\dif t} = \\Gamma(f). \\label{boltzfull} \\end{equation} Despite the right hand side being non-zero (given by Equations~\\ref{mbexts} and \\ref{lbexts} for the MB and LB statistics, respectively), the above equation does not contradict the assumption of a stationary state. A stationary, or steady, state is the Eulerian viewpoint, \\ie\\ $\\partial f/\\partial t = 0$, while the Boltzmann equation above is a Lagrangian viewpoint. $\\partial f/\\partial t = 0$ does not imply $\\dif f/\\dif t = 0$. In other words, the relaxation function $\\Gamma$ does not determine the explicit time-dependence of $f$, which must be zero for stationary states, but rather describes a flux of occupied cells through phase-space, as do the velocity-driven ($\\vv \\bcdot \\vnab f$) and acceleration-driven ($\\va \\bcdot \\vnab_v f$) flux terms (c.f. Equation~\\ref{boltz}). In the context of the Lagrangian derivative, we can think of $t$ simply as a parameter that indicates the location along a particle's or cell's path through phase-space. In collisional systems, the relaxation function is called the collision term and is usually dealt with in a Fokker-Planck approximation scheme. In these systems the particles, through two-body encounters, gradually disperse over the whole available phase-space, and so $\\dif f/\\dif t$ following any given particle in an evolving system does not stay constant, but generally decreases with time. In a general collisionless system, this term represents the processes like violent relaxation and phase mixing, on the coarse-grained scale. In a steady-state collisionless system, the large scale processes like violent relaxation no longer operate and the only changes happen on microscopic scales. In this context, the left hand side of Equation~\\ref{boltzfull} describes how a particle, or a cell moves through the system (and $t$ is the parameter). Therefore it is not unexpected that over some portions of its motion the coarse-grained density around it will be increasing and over others, it will be decreasing. For the MB case, $\\Gamma_{\\rm MB} > 0$ for $f>0$, implying that $f$ should continually grow. However, this is impossible as the coarse-grained distribution function is limited to a maximum value $\\eta$. This contradiction arises as MB statistics are valid only when $f \\ll \\eta$ so that macro-cells are not multiply occupied. On the other hand, the LB case does not present any contradictions with the Boltzmann equation. $\\Gamma_{\\rm LB}$ is zero when $f=0$ and $f=\\eta$, and is positive over the vast majority of the intervening range. This behavior guarantees that when the coarse-grained density reaches its maximum value, the relaxation term disappears and the system behaves collisionlessly even at the coarse-grained level. To sum up, if extremizing entropy production in self-gravitating systems does lead to steady-state configurations, then Equation~\\ref{boltzfull} with the appropriate expressions for $\\Gamma$ (as given by Equations~\\ref{mbexts} and \\ref{lbexts}) describes the steady state of self-gravitating collisionless systems. The relaxation term $\\Gamma$ describes the continual evolution of the coarse-grained distribution function, which is due to the combimation of the dynamical evolution of the fine-grained DF on microscopic scales, and coarse-graining." }, "1101/1101.5671_arXiv.txt": { "abstract": "Using a combination of self-consistent and test-particle techniques, Identikit~1 provided a way to vary the initial geometry of a galactic collision and instantly visualize the outcome. Identikit~2 uses the same techniques to define a mapping from the current morphology and kinematics of a tidal encounter back to the initial conditions. By requiring that various regions along a tidal feature all originate from a single disc with a unique orientation, this mapping can be used to derive the initial collision geometry. In addition, Identikit~2 offers a robust way to measure how well a particular model reproduces the morphology and kinematics of a pair of interacting galaxies. A set of eight self-consistent simulations is used to demonstrate the algorithm's ability to search a ten-dimensional parameter space and find near-optimal matches; all eight systems are successfully reconstructed. ", "introduction": "Dynamical modeling of specific pairs of interacting galaxies is a subject with considerable history. \\citet[][hereafter TT72]{TT72} bolstered their interpretation of bridges and tails as tidal features by presenting test-particle models of Arp~295, M~51, NGC~4676, and NGC~4038/9; the power of such models was demonstrated when \\citet{S74} confirmed TT72's prediction for the relative velocities of the two galaxies making up NGC~4676. As observational and numerical techniques have improved, modeling of interacting systems has generated an extensive literature (see \\citealt{BH09}, hereafter BH09, for a partial list). The motivation for dynamical modeling has evolved over time. While early studies focused on testing the tidal theory of galactic encounters, more recent work has used dynamical modeling to help interpret observations, probe the structure and dynamics of unseen matter, and reconstruct the dynamical histories of merging galaxies. Despite the advances of the past few decades, it remains \\textok{a challenge} to create models matching the detailed morphology and kinematics of a pair of colliding galaxies. One fundamental \\textok{complication} is the inherent uncertainty in inferring the distribution \\textit{and} dynamics of dark matter strictly by its effects on luminous material. Another is the \\textok{violent reprocessing} of interstellar material -- including rapid star formation -- in galaxy collisions. However, there are three rather technical issues which also limit progress: \\begin{enumerate} \\renewcommand{\\theenumi}{(\\arabic{enumi})} \\item A galactic collision is described by a large number of parameters which interact in highly non-linear ways. \\item Simulating galactic collisions is computationally intensive. \\item The criteria for a successful match are not easily translated into quantitative terms. \\end{enumerate} The parameters necessary to simulate an encounter of two disc galaxies fall into three groups, as illustrated in Fig.~\\ref{parameters}. The first group specifies the initial orbits of the galaxies; assuming these orbits are asymptotically Keplerian at early times, the required parameters are the periapsis separation $p$, the orbital eccentricity $e$, and the mass ratio $\\mu$. The second group describes the spin vector of disc $d$ (where $d = {1}, {2}$) with respect to the angular momentum of the relative orbit and the separation vector between the galaxies at periapsis; this vector is parametrized by the inclination $i_{d}$ and argument to periapsis $\\omega_{d}$ of each disc. Together with any parameters needed to describe the internal structures of the two galaxies, the first and second groups specify the \\textit{initial conditions} for a galactic encounter. The third group consists of the time $t$ since first periapsis, and parameters which map the simulation onto the observational plane: three Euler angles $\\theta_\\alpha$ specifying the viewing direction, scaling factors $\\mathcal{L}$ and $\\mathcal{V}$ \\textok{which transform dimensionless simulation positions and velocities, respectively, into real physical quantities}, and the centre-of-mass position on the plane of the sky $\\vect{R}_\\mathrm{c}$ and radial velocity $V_\\mathrm{c}$. These parameters may be chosen \\textit{after} a simulation has been run. \\begin{figure} \\begin{center} \\includegraphics[clip=true,width=0.25\\columnwidth]{parameters.ps} \\caption{An abstract representation of the sixteen-dimensional parameter space of galaxy interactions. The radial coordinate represents the initial orbit, the azimuthal coordinate represents the disc orientations, and the vertical coordinate represents the parameters chosen after a simulation is run. A conventional N-body simulation explores the parameter subspace represented by the dotted line, while a single Identikit simulation explores the entire cylindrical surface. \\label{parameters}} \\end{center} \\end{figure} Of these sixteen parameters, only a few have a~priori constraints. TT72 argued that the orbital eccentricity should be $e \\simeq 1$; this is generally supported by cosmological simulations \\citep{KB06}, although the $e$ distribution extracted from these simulations includes a tail to $e < 1$. The mass ratio $\\mu$ may be estimated from the relative luminosities of the two galaxies -- provided that the galaxies are still distinct and that interaction-induced star formation has not significantly altered their luminosities. Finally, the scale factors $\\mathcal{L}$ and $\\mathcal{V}$ are are not completely arbitrary since the pre-encounter galaxies should have radii and circular velocities comparable to those of other disc galaxies. \\subsection{Identikit 1} Identikit simulations combine test-particle and self-consistent techniques (BH09). Each galaxy is modeled by an initially spherical configuration of massive particles with cumulative mass profile $m(r)$, in which is embedded a spherical swarm of massless test particles on initially circular orbits \\textok{with angular momenta uniformly distributed over all directions}. Two such models are launched towards each other with orbital parameters $(p, \\mu, e)$. During the ensuing encounter, the massive components interact self-consistently, approximating the time-dependent potential and orbit decay of a fully self-consistent galactic collision. The test particles mimic the tidal response of embedded discs with all possible spin vectors; once such a simulation has been run, selecting the appropriate subset of test particles yields a good approximation to the tidal response of any particular disc. In the simplest Identikit implementation, the test particles initially populating each galaxy model have the same radial distribution as the discs they are intended to mimic. Each test particle $i$ is associated with a normalized vector $\\vect{s}_i \\in \\mathsf{S}^2$ which records the direction of the particle's \\textit{initial} angular momentum with respect to the centre of its parent galaxy. An Identikit simulation yields a Monte Carlo representation of an `extended' distribution function, \\begin{equation} g_{d}(\\vect{r}, \\vect{v}, \\vect{s}; t) \\doteq \\sum\\nolimits_{\\textstyle i} \\delta^3(\\vect{r} - \\vect{r}_i(t)) \\, \\delta^3(\\vect{v} - \\vect{v}_i(t)) \\, \\delta^2(\\vect{s} - \\vect{s}_i) \\, . \\end{equation} This function gives the phase-space number density at time $t$ of test particles from disc ${d}$ with position $\\vect{r}$, velocity $\\vect{v}$ which initially had angular momentum direction $\\vect{s}$. Here and throughout, `$\\doteq$' is used throughout to indicate an explicit Monte Carlo expression. On the right-hand side, $\\vect{r}_i(t)$ and $\\vect{v}_i(t)$ are the position and velocity of test particle $i$ at time $t$; these are understood to \\textok{also} depend on the orbital parameters $(p, \\mu, e)$. Once this extended function has been constructed, the distribution function $f_{d}$ for a disc with a specific initial spin $\\vect{s}_0$ is estimated by \\begin{equation} f_{d} (\\vect{r}, \\vect{v}; t) \\propto \\int\\nolimits_{\\textstyle \\tilde{\\vect{s}}_0} d^2\\vect{s} \\, g_{d}(\\vect{r}, \\vect{v}, \\vect{s}; t) \\doteq \\sum\\nolimits_{\\textstyle \\vect{s}_i \\in \\tilde{\\vect{s}}_0} \\, \\delta^3(\\vect{r} - \\vect{r}_i(t)) \\, \\delta^3(\\vect{v} - \\vect{v}_i(t)) \\, , \\end{equation} where $\\tilde{\\vect{s}}_0 = \\{ \\vect{s} \\in \\mathsf{S}^2 \\,|\\, \\vect{s} \\cdot \\vect{s}_0 \\ge 1-\\sigma \\}$ and $\\sigma \\ll 1$ is a tolerance parameter which determines the solid angle contributing to the estimate of $f_{d}$. The reason why $f_{d}$ is estimated by integrating over a finite solid angle $\\tilde{\\vect{s}}_0$ is not obvious; it may seem enough to simply evaluate $g_{d}(\\vect{r}, \\vect{v}, \\vect{s}_0; t)$. However, $g_{d}$ and $f_{d}$ are both represented in a Monte Carlo fashion. To sample $f_{d}$ well enough for a visual comparison with observational data requires a few thousand particles; if $g_{d}$ is represented by $N_\\mathrm{test} \\sim 10^5$ to~$10^6$ test particles per galaxy, this requires $\\sigma \\simeq 10^{-2}$. This simple version of Identikit has some drawbacks. Only a small percentage of the test particles are initially placed at large radii where they are responsive to tidal forces; \\textok{this wastes computer time}. \\textok{Moreover}, simulated discs \\textok{defined by $\\vect{s}_i \\cdot \\vect{s}_0 \\ge 1 - \\sigma$} have scale heights which increase linearly with radius \\textok{and look unrealistic when viewed edge-on}. Both of these flaws can be addressed by radially biasing the distribution of test particles. Following BH09, the test particle density is multiplied by a factor of $r^2$, and each test particle $i$ is given a weight $\\xi_i = \\mathrm{max}(r^\\mathrm{init}_i/r_\\mathrm{min}, 1)^{-2}$, where $r^\\mathrm{init}_i$ is the initial orbital radius of particle $i$, and $r_\\mathrm{min}$ is a small cut-off radius. The extended distribution function is then \\begin{equation} g_{d}(\\vect{r}, \\vect{v}, \\vect{s}, \\xi; t) \\doteq \\sum\\nolimits_{\\textstyle i} \\delta^3(\\vect{r} - \\vect{r}_i(t)) \\, \\delta^3(\\vect{v} - \\vect{v}_i(t)) \\, \\delta^2(\\vect{s} - \\vect{s}_i) \\delta(\\xi - \\xi_i) \\, , \\end{equation} and the expression for $f_{d}$ becomes \\begin{equation} f_{d} (\\vect{r}, \\vect{v}; t) \\propto \\int d\\xi \\int\\nolimits_{\\textstyle \\tilde{\\vect{s}}_0(\\xi)} d^2\\vect{s} \\, g_{d}(\\vect{r}, \\vect{v}, \\vect{s}, \\xi; t) \\doteq \\sum\\nolimits_{\\textstyle \\vect{s}_i \\in \\tilde{\\vect{s}}_0(\\xi_i)} \\, \\delta^3(\\vect{r} - \\vect{r}_i(t)) \\, \\delta^3(\\vect{v} - \\vect{v}_i(t)) \\, , \\end{equation} where $\\tilde{\\vect{s}}_0(\\xi) = \\{ \\vect{s} \\in \\mathsf{S}^2 \\,|\\, \\vect{s} \\cdot \\vect{s}_0 \\ge 1 - \\sigma\\xi \\}$. ", "conclusions": "The problem of finding a model matching the morphology and kinematics of a pair of colliding galaxies has generally been solved \\textok{by} a process of trial-and-error, informed by physical insight into the dynamics of tidal interactions. In this approach, observational data can't be used to derive initial conditions directly; instead, the outcome of a model calculation is compared to the observations, and the initial conditions are modified on the basis of this comparison. Identikit~2 offers a shortcut: an important component of the initial conditions -- the initial spin vectors of the interacting discs -- can be derived directly from the observed morphology and kinematics of the tidal features. Moreover, by providing a way to assess the quality of a solution, the algorithm can be used to search parameter space automatically. The Identikit~2 algorithm derives the initial spin vector of a tidally interacting disc by simultaneously populating a set of phase-space regions which trace that disc's tidal features. It may seem odd not to make any use of the \\textit{amount} of material found in each region, but this `omission' is deliberate. The tracers commonly used to study the kinematics of interacting systems don't necessarily obey a continuity equation; for example, neutral hydrogen may be ionized or converted to molecular form, so the amount of H{\\footnotesize{I}} found in a given region can't be predicted by purely dynamical models. On the other hand, even if much of the H{\\footnotesize{I}} in a tidal feature has been converted to \\textok{another phase}, the remaining \\textok{H{\\footnotesize{I}} may} still provide a useful constraint on disc spin as long as it has followed a free-fall trajectory. Since the algorithm does not use a $\\chi^2$ statistic to quantify goodness-of-fit, it's not straightforward to obtain precise confidence limits on solutions. Nonetheless, inspection of the function $\\Lambda$ and its constituent factors (Fig.~\\ref{view_ball}) offers some insight into a solution's uniqueness and accuracy. It may be worth exploring the landscape of this function in more detail. For example, instead of simply maximizing $\\Lambda$, the algorithm could examine solutions around the peak, and look for secondary peaks which may represent alternate solutions. A similar examination of the functions $\\Omega^*_{d}$ could likewise reveal uncertainties and alternate solutions for spin direction. \\subsection{Other parameters} \\label{other_parameters} The version of the algorithm tested here performs a blind search of viewing direction, parameterized by $(\\theta_X, \\theta_Y)$. It constrains four other viewing parameters -- the line-of-sight rotation $\\theta_Z$, the length scale $\\mathcal{L}$, and two components of the offset $\\vect{R}_\\mathrm{c}$ -- from the positions of the galaxy centres, and computes initial disc spins $(i_{d}, \\omega_{d})$ using the regions tracking each disc. This accounts for ten parameters out of the sixteen described in Fig.~\\ref{parameters}. Preliminary experiments indicate that some of the other parameters can also be determined by maximizing $\\Lambda(t, \\dots)$. For example, the test in \\S~\\ref{test_param_search} was repeated varying the time since periapsis $t$ between $0.5$ and $1.5$; the times maximizing $\\Lambda$ clustered around the actual value $t_\\mathrm{true} = 1$, with an r.m.s.~of $0.18$. Further tests with multiple parameters in play are necessary; it will be interesting to see if the algorithm can \\textit{simultaneously} fit for the time $t$, periapsis separation $p$, and velocity scale $\\mathcal{V}$ as \\textok{well as} BH09 did. The number of unknowns remaining depends on the type of system to be modeled as well as the nature of the data available. Suppose that accurate systemic velocities for both members of a pair of well-separated galaxies are available. By forcing the model nuclei \\textok{to} coincide with their real counterparts in velocity as well as position, locking can determine the velocity scale $\\mathcal{V}$ and offset $V_\\mathrm{c}$ in addition to $\\theta_Z$, $\\mathcal{L}$, and $\\vect{R}_\\mathrm{c}$. Adopting $e = 1$ and using photometry to estimate the mass ratio $\\mu$ leaves just four parameters -- the viewing direction $(\\theta_X, \\theta_Y)$, the periapsis separation $p$, and time since periapsis $t$ -- as unknowns. In this case a blind search of these parameters seems reasonable. However, sufficiently accurate systemic velocities may be hard to determine; galactic nuclei have large velocity dispersions, and different tracers (stars, H${}_\\alpha$, H{\\footnotesize{I}}, CO) often give results differing by several tens of km/s. Only in cases where the systemic velocities of the nuclei differ by more than the uncertainties is velocity locking likely to yield useful constraints. At the other extreme, fully merged systems such as NGC~7252 \\citep{S77, HGvGS94, HM95} represent the most difficult class of objects to model. The position and systemic velocity of a merger remnant constrain the offsets $\\vect{R}_\\mathrm{c}$ and $V_\\mathrm{c}$, but even assuming $e = 1$, a total of \\textit{eight} orbit and viewing parameters (Fig.~\\ref{parameters}) remain indeterminate. Since the real galaxies have already merged, only models run past merger need be considered. Nonetheless, the available parameter space is still very large, and blindly searching for possible solutions may not be very rewarding. Several groups have used genetic algorithms to automate the search for models of interacting galaxies \\citep{W98, TK01}. In these algorithms, a population of candidate solutions compete to match the observational data; the less successful candidates are eliminated, and the most successful reproduce to replenish the population. After enough generations have passed, the population converges toward a solution matching the observations. Unlike simple `hill-climbing' strategies, genetic algorithms are unlikely to be trapped by local maxima (a familiar \\textok{problem with} naive trial-and-error modeling). To date, most genetic algorithms have determined the reproductive fitness of candidate solutions by comparing low-resolution images of the actual system and candidate pixel by pixel, with only limited use of velocity data \\citep{WD01}. For such a comparison to be truly meaningful, the observed material -- typically stars or H{\\footnotesize{I}} -- must obey a continuity equation; as noted above, this may be violated in real systems. Moreover, as Figs.~\\ref{example1} and~\\ref{example2} illustrate, morphology alone is not enough to strongly constrain disc spins; low-resolution versions of the $(X,Y)$ images in the upper left of these figures would be almost indistinguishable. If further tests confirm that maximizing $\\Lambda(t, \\dots)$ is an effective way of constraining viewing and orbit parameters, it may be possible to combine Identikit~2 with a genetic algorithm. In this hybrid approach, each member of the candidate population would define a specific choice of viewing direction, periapsis time and separation, velocity scale, and possibly orbital eccentricity. The remaining viewing parameters would be determined by locking the centres, and initial spin directions could then be derived directly. The resulting $\\Lambda$ value could be used to determine the fitness of the candidate solution. This approach combines the strengths of both algorithms. A genetic algorithm should be able to out-perform a blind search without getting stuck on local maxima. Meanwhile, Identikit~2 could efficiently determine disc spins and provide a robust way to define reproductive fitness which fully includes velocity information and does not assume continuity. \\subsection{Mass models} \\label{mass_models} The mass models adopted in Identikit simulations will almost certainly influence the algorithm's accuracy. In the tests presented here, the \\textit{same} mass model has been used to construct self-consistent simulations and their Identikit reproductions. However, rotation curves of real disc galaxies exhibit a variety of shapes \\citep*{CvG91, CGH06}; small galaxies often have rising curves, while massive galaxies may have flat or even falling curves. This diversity presumably arises because real galaxies have a range of initial angular momenta, bulge/disc/halo mass ratios, and assembly histories \\citep*{MMW98}. Several studies have shown that rotation curve shape (equivalently, potential well or halo structure) strongly influences the development of tidal features; in particular, long tidal tails can be inhibited by sufficiently deep galactic potential wells \\citetext{\\citealp*{DMH96}, \\citeyear{DMH99}; \\citealp{SW99}}. What happens if the mass model used in an Identikit simulation does not match the structure of the galaxies being modeled? Suppose for a moment that the Identikit mass model and the real galaxy have the similar rotation curves but apportion mass differently between various components. In this situation the Identikit algorithm will probably yield good encounter parameters despite the mismatch. One test is shown in Figs.~\\ref{cdf_view} and~\\ref{cdf_spin}, where the results obtained by giving all particles equal weights (red lines) accurately reproduce the viewing and spin directions of all eight test systems. This can be interpreted as an experiment in which real galaxies with exponential discs (surface density $\\Sigma \\propto e^{- \\alpha R}$) are matched to Identikit models with radially biased discs ($\\Sigma \\propto R^2 e^{- \\alpha R}$) but identical rotation curves. At the opposite extreme, suppose an attempt is made to match a pair of galaxies with long, well-developed tidal tails -- which imply relatively shallow galactic potential wells -- using an Identikit model with a very deep potential well. Since such a model would be unable to produce long tidal tails, it's unlikely that \\textit{any} set of encounter parameters could populate phase-space regions near the ends of the tails. As a result, the algorithm would fail to find a solution. This `failure' points to a flaw in the adopted mass model; the obvious recourse is to try a model with a shallower potential well. Between these extremes is a grey area in which Identikit~2 may yield a good match to the morphology and kinematics of an interacting system without providing accurate values for all encounter parameters. Consider the problem of modeling a system observed just after first encounter, in which tidal features have not yet had time to develop and probe the full extent of the galactic potential wells. Some constraint on potential well depth may still be afforded by the relative velocities of the two galaxies, but depth is likely to be degenerate with orbital eccentricity; for example, a large line-of-sight velocity difference may arise because the galaxies have deep potential wells \\textit{or} because their initial orbit was hyperbolic ($e > 1$). What can be learned about galactic structure by matching the morphology and kinematics of a pair of interacting galaxies with a dynamical model? This larger question is independent of the specific technique -- Identikit~2, genetic algorithm, or trial-and-error -- used to produce the model. Clearly there's no simple answer; the outcome will vary from system to system. Identikit~2 offers a practical way to explore this question without laborious trial-and-error modeling. It's straightforward to construct self-consistent simulations with various mass models; these could then be tested against Identikit models with different rotation curves. Of particular interest will be tests varying the ratio of circular to escape velocity, which appears to predict the extent of the tidal tails produced in an encounter \\citetext{\\citealp{SW99}; \\citealp{DMH99}}. \\subsection{Coda} Determining additional encounter parameters (\\S~\\ref{other_parameters}) and exploring the results of different mass models (\\S~\\ref{mass_models}) both present intriguing theoretical problems. Beyond these, the effects of `pre-existing conditions' such as bars and warps suggest additional lines of investigation. However, the algorithm already appears quite effective at reconstructing galactic collisions. It will be very interesting to see if it works with real galaxy data. Source code for this algorithm is available at \\texttt{http://www.ifa.hawaii.edu/faculty/barnes/research/identikit/}." }, "1101/1101.2051_arXiv.txt": { "abstract": "We describe a new method which achieves high precision Very Long Baseline Interferometry (VLBI) astrometry in observations at millimeter wavelengths. It combines fast frequency-switching observations, to correct for the dominant non-dispersive tropospheric fluctuations, with slow source-switching observations, for the remaining ionospheric dispersive terms. We call this method Source-Frequency Phase Referencing. Provided that the switching cycles match the properties of the propagation media, one can recover the source astrometry. We present an analytic description of the two-step calibration strategy, along with an error analysis to characterize its performance. Also, we provide observational demonstrations of a successful application with observations using the Very Long Baseline Array at 86 GHz of the pairs of sources 3C274 \\& 3C273 and 1308+326 \\& 1308+328, under various conditions. We conclude that this method is widely applicable to millimeter VLBI observations of many target sources, and unique in providing {\\it bona-fide} astrometrically registered images and high precision relative astrometric measurements in mm-VLBI using existing and newly built instruments. ", "introduction": "\\label{sec:intro} The comparative study of the radiation emitted at multiple radio bands has proved to be a useful tool in astronomy for the investigation of the nature of the emission mechanisms and to probe the physical conditions of the emitting regions. Multi-frequency observations with the high spatial resolution obtained with Very Long Baseline Interferometry (VLBI) are suitable for the study of extragalactic radio sources, such as AGNs, providing detailed images of the radiation from the relativistic jets, which are launched from the central engine that powers the sources. Observations at increasingly high frequencies offer the prospect of an increasingly deep exploration of the inner jet region, closer to the central engine. By comparing well aligned high resolution images at multiple frequencies it is possible to map the spectral index across the jet structure. The spectral index map carries direct information about the physical conditions in the jet regions, and, potentially, with observations at the highest frequencies, on the structure of the central engine \\citep{bh_shadow_1}. Also, the standard model \\citep{blandford_79} predicts changes in the apparent position of the observed ``core'' component, at the base of the jet, in observations at different frequencies as a result of opacity effects in the jet. These position changes are called {\\it core-shifts} and hold a direct relationship with the conditions in the nuclear region at the base of the jet where the ``core'' is located. For both studies the precise alignment of the source images is mandatory to assess true intrinsic source properties using multi-frequency comparison techniques, otherwise alignment errors will result in misleading conclusions. Standard VLBI images, which are created using self-calibration techniques, provide exquisite detail on the source structure but lack astrometric information. The astrometry is lost in the process of removing the residual contributions arising from imprecise modeling of the propagation effects through the atmosphere and the use of independent frequency standards at each telescope, among others. The special analysis technique of Phase Referencing ({\\em hereafter} PR) is required to preserve the astrometric information. PR relies in the use of interleaving observations of an external calibrator source to correct the errors present in the target dataset, rather than using the target data themselves as in standard VLBI analysis \\citep{alef88}. By doing this it is possible to achieve high-precision (relative) {\\it bona-fide} astrometric measurements of the angular separation between the two sources. The switching time and switching angle are critical parameters for the success of PR techniques. Typical switching values are estimated using the temporal and spatial structure-function of the atmospheric fluctuations, and are dependent on the observing frequency \\citep{memo_20}. Conventional phase referencing has been successfully used at cm-wavelengths (from 1.4 to 43-GHz) for which atmospheric effects are moderate and calibrator sources are easy to find. VLBI observations at mm-wavelengths are challenging because of the lower sensitivity of the instruments, intrinsically lower source fluxes and shorter coherence times imposed by the rapid variations of the water vapor content in the troposphere. For the same reasons, VLBI astrometry with conventional PR at high frequencies, beyond 43 GHz, is practically impossible due to the extremely short telescope switching times involved. The only successful demonstration was with the VLBA at 86-GHz for a pair of sources only 14$^\\prime$ apart \\citep{porcas_02_pr86,porcas_03_pr86}. An alternative approach to overcome the tropospheric limitations in observations at high frequencies consists in using fast frequency-switching, instead of fast source-switching as in PR, on the grounds that the tropospheric excess path delay, being independent of the observing frequency (i.e. it is a non-dispersive medium), can be corrected for using dual frequency observations. This technique has been attempted in VLBI resulting in longer effective coherence times at mm-wavelengths, but it failed to recover astrometry due to remaining dispersive ionospheric and instrumental errors \\citep{middelberg_05_fs}. We propose that astrometry at high frequencies can be achieved by combining alternating observations at two frequencies, to correct for the non-dispersive propagation media effects, and of two sources, to correct for the remaining dispersive effects, providing suitable switching times and switching angles are used. We term this new technique {\\it Source-Frequency Phase Referencing} ({\\em hereafter} SFPR). The direct outcomes of this technique are: high precision {\\it bona-fide} astrometric measurements of the angular separation between emitting regions in the two frequency bands, and increased coherence time in VLBI observations at the highest frequencies. Hence it allows {\\it bona-fide} astrometric registration of VLBI maps in the high frequency regime, beyond the threshold for conventional PR techniques. For example, applied to AGN-jets it would allow spectral index and core-shift measurements; applied to spectral line VLBI observations this would allow the alignment of the spatial distribution of emission arising from multiple maser transitions of a given molecule. Such information is of great interest in astrophysics (e.g. \\cite{lob_98_cj,m87_dodson,soria_07,rioja_pasj}). Moreover, the combination of SFPR and conventional PR techniques holds the prospect of providing high precision relative astrometric measurements of positions with respect to an external reference (i.e. a calibrator source). This would enable VLBI multi-epoch proper motion and parallax studies at the highest frequencies. This paper presents an analytical description of this new technique that enables high precision VLBI astrometric measurements in the highest frequency regime, {along with an experimental demonstration using VLBA observations at 43 and 86\\,GHz}. Also, we present a comparative error analysis and discuss the feasibility of the new technique in the context of existing and newly built instruments. A comprehensive computer-simulation study of the SFPR performance will be presented elsewhere. ", "conclusions": "\\label{sec:disc} \\subsection{\\it Validation of SFPR method:} We have developed a new two-step calibration technique called SFPR that, by precisely compensating for the effect of the propagation medium in VLBI observations, enables high precision astrometry even at the highest frequencies where conventional PR techniques fail. Previous attempts using fast frequency-switching observations achieved an increased coherence time as a result of the dual-frequency tropospheric calibration, which enabled the detection of weak sources but failed to provide astrometry due to remaining dispersive errors \\citep{middelberg_05_fs}. Our method addresses this issue with a second calibration step, that corrects for those and hence enables astrometry. We have provided experimental demonstration of the ability of the SFPR method to disentangle the astrometric signature from the other contributions using VLBA observations at 86 GHz, the highest VLBA frequency, and 43 GHz. This ``chromatic'' astrometric signature corresponds to the angular separation between the emitting regions at the two observing frequency bands in the target source, assuming an achromatic calibrator. Therefore, our SFPR method is a valuable tool for studies that require comparison and {\\it bona-fide} astrometric registration of images at two or more frequencies, even at the highest frequencies possible with VLBI. Additionally, when combined with conventional PR observations at $\\nu^{{\\rm low}}$, SFPR can provide `PR-like' astrometry at $\\nu^{{\\rm high}}$. That is, measurements of the positions with respect to an external calibrator. Such astrometric measurements can be used for position stability, proper motion and parallax studies at frequencies beyond the traditional limit of phase referencing ($\\sim$43 GHz). In summary this method offers the means to expand the benefits that conventional PR techniques offer in the moderate frequency regime, into the highest frequencies used in VLBI. In previous works dual frequency observations have been used for the detection of weak sources at mm-VLBI; now, with the SFPR technique, {\\it bona-fide} high precision astrometric mm-VLBI (and sub-mm) can also be performed. We have carried out an analytical error analysis to characterize the performance of the SFPR method. This analysis shows that when using frequency-switching observations the dominant source of errors are the random fluctuations in the {\\it dynamic} component of the troposphere. Also that the magnitude of this error depends on the frequency switching cycle at the observations and the weather conditions, irrespective of the pair angular separation. These findings are in good qualitative and quantitative agreement with the experimental results from our SFPR observations of two pairs of sources with very different angular separations. This validates the new astrometric method and the error analysis, confirming its potential in mm-VLBI. Furthermore it gives confidence in our extrapolations into domains yet untested, that is to the sub-mm regime, and to simultaneous dual frequency observations. \\subsection{\\it Broad scope of application:} The constraints for successful SFPR observations are relatively easy to fulfill. The SFPR method has been successfully demonstrated with VLBA observations of a pair of sources with a large (10$^o$) angular separation, at 43/86\\,GHz. These results are encouraging as they suggest that the SFPR method would work with any other combination of integer-ratio frequencies provided suitable frequency switching cycles are used, and that the angular separation between the calibrator and target sources can be large and telescope switching cycles long; certainly much more than those required for conventional PR. The key elements for success of the new method in the high frequency regime are, firstly, that the frequency switching operation in SFPR can be carried out much faster than the source switching counterpart in PR. Secondly, that finding a suitable SFPR calibrator source is relatively easy, unlike for PR, because wider angular separations and longer switching cycles are allowable, along with the extended coherence at the high frequencies. The SFPR method can be implemented as a regular observing mode with existing instruments, such as the VLBA, which support fast frequency switching operations. The best astrometric performance is achieved when multiple frequency bands can be observed simultaneously, as this provides an exact tropospheric correction in all weather conditions, eliminates the need of phase connection and increases the on-source time. The Korean VLBI Network (KVN), equipped with multi-channel receivers at 22/43/86/129 GHz, and telescopes like Yebes and Haystack are among the instruments that can carry out simultaneous observations of multiple high frequency bands to achieve maximum benefit from the SFPR technique. Based on our analysis it is clear that simultaneous frequency observations can be very useful in high frequency VLBI. We strongly suggest that this capability is included in all next-generation instruments. We conclude that this method is broadly applicable to mm-VLBI observations of many target sources, and unique in providing {\\it bona-fide} astrometrically registered images and high precision relative astrometric measurements using existing and newly built instruments. \\subsection{\\it Applications to space-VLBI:} Additionally to the errors in the atmospheric propagation models, the geometric errors also introduce inaccuracies in the PR analysis, which can ultimately prevent its application. While for VLBI ground arrays the telescope coordinates can be accurately measured with dedicated geodesy campaigns this is of particular concern for space-VLBI observations, since the precise orbit determination for a satellite antenna is much more complicated. For example, for VSOP-2, the accuracy in the orbit reconstruction required for successful PR observations at 43\\,GHz must be better than 10-cm (A07), which is challenging. For comparison, the typical orbit determination accuracy for its predecessor, the HALCA satellite, was 2--5 meters using Doppler measurements from the Ku-band link \\citep{porcas_vsop_00,rioja_vsop_09}. Strategies to achieve the 10-cm level of accuracy using global satellite navigation systems and satellite laser ranging techniques are presented in \\cite{asaki_orbit}. Alternatively, the SFPR method automatically corrects for any geometric errors, including any orbit determination errors, irrespective of their magnitude. There are no specific requirements on the orbit accuracy for SFPR analysis other than those imposed by the correlator fringe field of view, which is typically many meters. In addition to the astrometric applications, the increased sensitivity resulting from longer coherence time is very useful because of the limited size of an orbiting antenna, particularly at the higher frequencies. Also, the fast frequency-switching operation for SFPR is less demanding than fast source-switching for PR, reducing the requirements on the satellite attitude control system. Further discussions can be found in \\citet{memo_32}. Therefore, we believe this method will be very useful for space VLBI missions. In particular, applied to VSOP-2 observations, it would enable increased sensitivity allowing the detection of weaker sources, and permit astrometric measurements and long term monitoring projects at 43 GHz, by using the calibration derived from interleaving observations at 22 GHz even with a coarse orbit determination. For other future space VLBI missions at high frequencies, e.g. ``Millimetron'' \\citep{millimetron}, having simultaneous observations at multiple frequencies combined with all or some aspects of the SFPR calibration techniques would enable enhanced sensitivity and astrometric capabilities. \\subsection{\\it Astrometric Precision \\& Applications} High precision astrometric and sensitive measurements are valuable tools to provide insight into astrophysical phenomena, as demonstrated by application of phase referencing techniques at a moderate frequency regime (i.e. up to 43 GHz). The SFPR method enables such measurements in the high frequency regime by means of an improved atmospheric calibration. Our analytical error analysis shows that the tropospheric {\\it static} component is readily compensated using SFPR techniques and that, in general the {\\it dynamic} component would be the dominant source of errors with fast frequency switching observations, leading to $\\sim\\,$a few tens of micro-arcseconds astrometric precision. Based on the repeatability of results from observations with the VLBA we estimate an astrometric precision of $\\sim \\, 20\\,\\mu$as. With simultaneous dual frequency observations both the {\\it static} and {\\it dynamic} components of the troposphere would be precisely compensated, and the much smaller ionospheric residuals become the dominant source of errors. Increasing the reference frequency can result in negligible ionospheric residuals, therefore SFPR techniques at high frequencies offer the prospect of achieving the theoretical astrometric precision set by the interferometer beam size and the signal-to-noise ratio \\citep{thompson_isra_2}. SFPR techniques applied to VLBI spectral line maser observations would allow a precise {\\it bona-fide} astrometric spatial registration of, for example, the SiO emission structures at different frequency bands (43, 86, 129 GHz) in the same source for studies of the circumstellar environment in AGB stars. When applied to AGN studies, mm and sub-mm VLBI observations probe the inner-jet regions. SFPR can add the measurement of core-shifts with micro-arcsecond precision, plus enable deeper observations for the detection of weak sources. For astrometric measurements relative to an external reference, in combination with conventional PR at $\\nu^{{\\rm low}}$, the final precision would be the quadratic sum of errors from both techniques. This would allow the application to proper motion studies of maser emission at the high frequencies, and precise astrometric monitoring programs of the `jet foot-prints' to unveil the cause of the observed jet-wobbling phenomena." }, "1101/1101.0054_arXiv.txt": { "abstract": "The zeroth-order component of the cosine expansion of the projected three-point correlation function is proposed for clustering analysis of cosmic large scale structure. These functions are third order statistics but can be measured similarly to the projected two-point correlations. Numerical experiments with N-body simulations indicate that the advocated statistics are redshift distortion free within $10\\%$ in the non-linear regime on scales $\\sim 0.2-10 h^{-1}$Mpc. Halo model prediction of the zeroth-order component of the projected three-point correlation function agrees with simulations within $\\sim 10\\%$. This lays the ground work for using these functions to perform joint analyses with the projected two-point correlation functions, exploring galaxy clustering properties in the framework of the halo model and relevant extensions. ", "introduction": "Observed large scale structure in the Universe is generally conjectured to arise from Gaussian initial condition or nearly so; the rather high level non-Gaussianity at present is due to the action of gravitational force and gas physics. The three-point correlation function (3PCF) is of the lowest order among correlation functions capable of probing such non-Gaussianity. With the recent increase of interest and the corresponding attempts to extract more information about structure formation processes and primordial non-Gaussianity from fine clustering patterns of galaxies, the 3PCF (or its counterpart in Fourier space, bispectrum) has attracted much attention in recent years \\citep[e.g][]{KayoEtal2004, NicholEtal2006, SmithEtal2008, JeongKomatsu2009, Sefusatti2009}. However, 3PCF is well known for its low return of investment compared with the two-point correlation function (2PCF). One major obstacle hindering the interpretation and consequently the application of 3PCF is the redshift distortion induced by the peculiar velocities of galaxies. Although effects of redshift distortion on 2PCF (or power spectrum) are not yet well understood analytically \\citep[e.g.][]{Scoccimarro2004}, approximations by incorporating pairwise velocity distribution have been proposed, validated and applied successfully to statistical analyses \\citep{Peebles1980, DavisPeebles1983, White2001, Seljak2001, KangEtal2002, Tinker2007, SmithEtal2008}. In the case of 3PCF (or bispectrum) analogous approach would involve higher order statistics of peculiar velocities. The complicated entanglement of redshift distortions with nonlinear gravitational dynamics and nonlinear biasing renders theoretical prediction extremely difficult in configuration space. In Fourier space and with the distant observer approximation, prediction of the bispectrum in redshift space in various perturbative and empirical schemes has been moderately successful, although none have been able to show satisfactory agreement with simulations \\citep{MatsubaraSuto1994, HivonEtal1995, VerdeEtal1998, ScoccimarroEtal1999}. The mostly accurate model to date appears to be the work of \\citet{SmithEtal2008}, a halo model extension implemented with higher order perturbation theory. One can eliminate the complexity of redshift distortion with projection of the correlation functions upon the plane perpendicular to the line-of-sight (LOS). Projected correlation functions are obtained by integrating over the anisotropic correlation functions along LOS, which effectively removes redshift distortions if the conservation of total number of galaxy pairs and triplets along LOS can be satisfied. Since thickness of a realistic sample is finite, galaxies near radial edges could enter or leave the sample space by their apparent movement due to peculiar velocity, such conservation is only approximately achieved if the sample is shallow, or redshifts are photometric. Violation of the conservation condition may bring non-negligible systematical bias on large scales \\citep{NockEtal2010}. Nevertheless, this is not a problem for most modern spectroscopic galaxy samples, and the bias actually can be minimized by careful design of estimation methodology. In comparison with the projected 2PCF that has been widely used to investigate clustering dependence on galaxy intrinsic properties, evolution history and environment and to distinguish cosmological models \\citep[e.g.][]{HawkinsEtal2003, ZhengEtal2007, BaldaufEtal2010, ZehaviEtal2010}, exploration and application of the projected 3PCF has been limited in the literature \\citep{JingBoerner1998, JingBorner2004a, Zheng2004, McBrideEtal2010}. Lack of accurate theoretical models of 3PCF prevents proper interpretation of measurements. \\citet{ScoccimarroCouchman2001} offered a phenomenological model based on hyper-extended perturbation theory for the bispectrum in the nonlinear regime. Their fitting formula is accurate on smaller scales but in the weakly and mildly nonlinear regimes it is improved upon by the empirical model of \\citet{PanColesSzapudi2007}. Both fail on very small scales, and neither can capture the signal of baryonic oscillation in bispectrum appropriately \\citep{SefusattiEtal2010}. The approach of halo model appears more promising, as it can reproduce most measurements in simulations for the bispectrum \\citep[e.g.][]{MaFry2000a, MaFry2000b, ScoccimarroEtal2001, SmithWattsSheth2006, SmithEtal2008} and the 3PCF in configuration space \\citep[e.g.][]{TakadaJain2003, WangEtal2004, FosalbaPanSzapudi2005}. In spite of disagreement with simulations for some configurations of 3PCF, the halo model is still more attractive than the phenomenological models for its clean and physically motivated parametrization to galaxy biasing through e.g. the machinery of the halo occupation distribution \\citep[HOD,][]{BerlindWeinberg2002}. Another reason for the scarce exploration of projected 3PCF is the complexity of estimation. Computational requirement of 3PCF is demanding for currently available computers when millions of points are typical. The additional task of decomposing the separations among three points for projected 3PCF adds to the CPU load. Furthermore, the 3PCF is already more prone to Poisson noise than the 2PCF, and typical bin width of scales for projected 3PCF is even smaller than for the normal 3PCF. In order to suppress discreteness effects for a reliable estimation, a high number density of points in the sample is crucial, but often unrealistic for real surveys. By analogy to the monople of 3PCF advocated by \\citet{PanSzapudi2005a, PanSzapudi2005b}, we show that a third-order statistical function similar to the angular average of the projected 3PCF is redshift distortion free and relatively easy to estimate and model theoretically. In the next section, the definitions, and relation with 3PCF together with estimation algorithm is described. Section 3 presents numerical properties of the new statistical measure while in section 4 we demonstrate the consistency of halo models to simulations of the new function. Summary and discussion are in the last section. ", "conclusions": "In this paper we propose a third-order correlation function for characterising galaxy clustering properties. The statistics $Z_0$ we advocate is the zeroth-order component of the projected 3PCF. Although $Z_0$ is a 3PCF, its estimation takes roughly the same amount of computing operation as the projected 2PCF. The algorithm can be easily implemented after moderate modification of a code for the projected 2PCF. Various numerical experiments confirm that $Z_0$ can be deemed to be redshift distortion free within approximately $10\\%$ for the regime where the scale perpendicular to LOS is $0.2<\\sigma<10h^{-1}$Mpc. In addition, the maximal integration scale $\\pi_{max}$ parallel to LOS during estimation ought to be greater than $\\sim 120h^{-1}$Mpc. A serious concern is that shot noise could ruin the estimation in the strongly nonlinear regime if the number density of points in a sample is too low. This requirement for a robust $Z_0$ measurement is tighter than for the projected 2PCF, but still weaker than the normal projected 3PCF, since $Z_0$ is an integral of the former. The criterion we suggest is $DDD>\\sim 100$. As we expected, the halo model provides satisfactory prediction to dark matter $Z_0$ of simulations within $\\sim10\\%$, if the classical Sheth-Tormen mass functions are used. Our computation indicates that extending the halo boundary is enough to yield good fit to simulations, while a hard cut-off to mass function is not as effective as previous works claimed. Substituting new functions of the halo mass distribution and halo biasing in high precision does not lead to significantly better agreement with simulations. Since the angular dependence in the projected 3PCF and the normal 3PCF is smeared out in $Z_0$, we conjecture that using an anisotropic halo profile probably will not significantly improve accuracy. A significant bias of halo model predicted $Z_0$ compared to simulations emerges in the weakly nonlinear regime, where halo models boil down to second-order perturbation theory; the latter is already known to be poor in predicting dark matter 3PCF. A more precise bispectrum from higher order perturbation theories may offer a way to increase precision \\citep[e.g.][]{Valageas2008, Sefusatti2009, BartoloEtal2010}. The principal reason for proposing $Z_0$ is to provide an efficient redshift distortion free 3PCF, complementary to the standard projected 2PCF, for galaxy clustering analyses. It is well known that the projected 2PCF itself is a Gaussian statistic only and thus has its limitations. Third order correlation functions, mainly carrying information about non-Gaussianity, are more sensitive to details of the galaxy distribution. Non-Gaussianity of galaxy distribution is generated by the nonlinear action of gravitational force and gas physics if the primordial density fluctuation of the universe after inflation is Gaussian. The degeneracy shown in projected 2PCF \\citep[e.g.][]{ZuEtal2008} may be broken if third order correlation functions are employed. The redshift distortion free feature of $Z_0$ on scales less than $10h^{-1}$Mpc defines its potential in investigating the relation of galaxies with their host halos, and the formation histories of galaxies and halos. Furthermore, the success of halo model prediction on dark matter $Z_0$ encourages us to apply $Z_0$ for analysing galaxies. In principle, with measurements from galaxy samples, $Z_0$ enables us to generalize and diagnose schemes of HOD, conditional luminosity function \\citep[CLF, ][]{YangEtal2003} and semi-analytical models \\citep[e.g.][]{Baugh2006} to third order statistics at cost of one additional free parameter, the halo boundary. Our present work is restricted to dark matter only, the behavior of $Z_0$ for biased objects remains unclear. Testing with mock galaxy samples before applying to real data will be necessary." }, "1101/1101.3488_arXiv.txt": { "abstract": "Long duration gamma-ray bursts are commonly associated with the deaths of massive stars. Spectroscopic studies using the afterglow as a light source provide a unique opportunity to unveil the medium surrounding it, probing the densest region of their galaxies. This material is usually in a low ionisation state and at large distances from the burst site, hence representing the normal interstellar medium in the galaxy. Here we present the case of GRB 090426 at $z=2.609$, whose optical spectrum indicates an almost fully ionised medium together with a low column density of neutral hydrogen. For the first time, we also observe variations in the Ly$\\alpha$ absorption line. Photoionisation modeling shows that we are probing material from the vicinity of the burst ($\\sim 80$ pc). The host galaxy is a complex of two luminous interacting galaxies, which might suggest that this burst could have occurred in an isolated star-forming region outside its host galaxy created in the interaction of the two galaxies. ", "introduction": "Optical afterglow studies of gamma-ray bursts (GRBs) provide a powerful tool for unveiling the interstellar medium (ISM) properties of their host galaxies. Long GRBs are connected to the death of a massive star, and their host galaxies are subsequently expected to be sites of heavy star-formation extending back to the early Universe. Metallicities determined from absorption lines of the ISM show higher values than those of, for example, QSO absorbers (Fynbo et al. 2006a; Savaglio et al. 2006; Savaglio 2009), though this might be only a sightline effect (Fynbo et al. 2008). Kinematics of absorption lines give some indications for galactic outflows (e.g. Th\\\"one et al. 2007), which are an expected phenomenon in star-forming galaxies. Usually, absorption lines in GRB afterglow spectra are constant in time. In a few cases metal absorption lines together with their corresponding fine structure transitions, caused by UV pumping from the intense afterglow radiation, vary in intensity following the change in the ionizing flux of the burst. Modeling of the line variability allows us to better constrain the distance from the GRB to the absorbing material, which typically ranges from $>150$\\,pc to a few kpc (e.g. Vreeswijk et al. 2007, D'Elia et al. 2009). The presence of neutral material in the spectra implies distances on the order of a few hundred pc to several kpc (e.g. Prochaska et al. 2007, Th\\\"one et al. 2007). For strong NV absorption lines, it has been suggested that those lines must arise from material close to the GRB site, at $<$100\\,pc (Prochaska et al. 2007). GRB 080310 showed variation in all resonant Fe II transitions which can only be explained with photoionization by the burst at a distance of $\\sim$ 100\\,pc (A. De Cia et al. in prep.). Excluding those few special cases, optical GRB afterglow spectroscopy seems to predominantly probe the ISM of the host galaxy and not the material around the GRB progenitor. Studying the circumburst environment, however, would allow us to verify the predictions of current progenitor models. Prompt emission X--ray data showed some evidence for variations in the absorbing column density (Amati et al. 2000). In the optical, absorption features from circumburst material are difficult to detect due to the flash ionization of the surrounding material by the GRB out to several tens of parsecs (Robinson et al. 2010), even though the GRB is predicted to reside in a dense environment. In addition, any remaining absorbing material would be hidden under the stronger absorption lines from the line-of-sight ISM in the host galaxy. There are several potential signatures which could demonstrate that we are observing material close to the GRB instead of the general ISM of the galaxy: time variability of the absorption line strength (Perna \\& Loeb 1998), a high ionization state of the medium (Prochaska et al. 2008), or absorption lines at large velocities compared to the redshift of the GRB. The last has been observed in a few cases (Fox et al. 2008, M\\o ller et al. 2002), but it cannot be excluded that those systems are intervening absorbers and hence not connected to the GRB. It was also suggested that a large difference between the optical and X-ray column densities might be a hint of observing highly ionized material, eventually close to the GRB (Watson et al. 2007, Campana et al. 2010). An ideal laboratory would be a GRB progenitor that exploded in a relatively isolated environment with little ISM along the line of sight. So far, no conclusive evidence for such conditions has been found, probably because such physical conditions are rare and only a few afterglows have time resolved afterglow spectroscopy. In this paper we present the first example for such a scenario, GRB 090426. GRB 090426 was discovered by the BAT (Burst Alert Telescope) $\\gamma$-ray telescope onboard the {\\it Swift} satellite (Gehrels et al. 2004) on April 26, 2009, 12:48 UT (Cummings et al. 2009) and had a duration of $T_{90} = 1.2 \\pm 0.3$\\,s (Sato et al. 2009). Soon after its discovery an optical counterpart was found by {\\it Swift} (Cummings et al. 2009) and the redshift was determined from an optical spectrum to be $z = 2.609$ (Levesque et al. 2010). Both its restframe duration (0.3\\,s) as well as its observed duration put GRB 090426 into the short burst category according to the commonly used classification based on the duration of the prompt emission. In \\S 2 we describe the optical spectra and afterglow observations. In \\S 3 we discuss the characteristics of GRB 090426 itself and some issues regarding its classification. \\S 4 describes the properties of the ISM in the line of sight as derived from optical and X--ray observations. In \\S 5 we deal in detail with the Ly$\\alpha$ line variation and model its behaviour with a photoionisation code. \\S 6 describes the properties of the host galaxy of GRB 090426. Discussion and conclusions are reported in \\S 7. \\noindent ", "conclusions": "The spectrum of GRB 090426 is the first compelling case of observed photoionised material from the star-forming region of the progenitor itself. Its low hydrogen column density is varying, the absorption lines show a highly ionised medium of high column density, and the large discrepancy between X--ray and optical absorption gives another indication for photoionised material being present. GRB spectra have shown a large range of hydrogen column densities, usually explained by different sightlines of the GRB through its host galaxy (Fynbo et al. 2009; Jakobsson et al. 2006), with some column densities even lower than the one measured for GRB 090426. However, if GRB 090426 exploded in, for example, the halo of the galaxy or a very low density ISM like a globular cluster, we would also expect low column densities of the metal absorption lines. A location in the halo had been suggested for GRB 070125 (Cenko et al. 2008) and GRB 071003 (Perley et al. 2008), both of which show a bright but featureless continuum with extremely weak absorption lines. GRBs with log $N_\\mathrm{HI}$/cm$^{-2}$ $<$ 20.0 tend to show a higher fraction of ionised versus neutral material, possibly explained by a lack of UV shielding by the lower density of hydrogen; however, some of them still contain a relatively large column density of neutral material. In addition, the afterglow showed that the GRB must have exploded in a non-low-density medium (Levesque et al. 2010). The strange properties of the spectrum of GRB 090426 can therefore not simply be explained by a sightline effect. Considering the fact that this GRB exploded in what seems to be a galaxy merger, another scenario might explain the strange spectrum we observe. The lack of absorption from the normal ISM of the host suggests that the progenitor did not reside in a star-forming region inside its host galaxy. A possible scenario is therefore that the progenitor was located in an isolated star-forming region outside the host itself, e.g. in something like a tidal tail created by the merging of the two galaxies. Thus, the sightline to the GRB would only intersect the material of the star-forming region, ionised by the GRB flux and eventually piled up by the wind from massive Wolf-Rayet stars, and not the dense ISM of the galaxy. We note, however, that the afterglow position is, in projection, close to the center of the galaxy. Another strange feature about this burst was its relatively short duration paired with a soft high energy emission spectrum. Naturally, the question arises if the unusual high-energy properties and those of the afterglow spectrum could be connected. From Fig. \\ref{090426:HR} we see that all other low column density GRBs occupy the normal parameter space for long bursts concerning HR and duration, so there seems to be no obvious correlation between HR or duration and low column density. Since the material observed in the afterglow spectrum is relatively far away from the burst compared to the prompt emission, which is expected to be produced by internal shocks within the jet of the GRB, a connection is not to be expected. Concerning the progenitor of the burst, Antonelli et al. (2009) and Levesque et al. (2010) had already concluded that, despite its short duration, the burst was likely due to the collapse of a massive star. Arguments include the star-forming host galaxy, the high afterglow luminosity, consistency with the Amati relation, and the strength of the absorption lines. In past years, the classification of bursts purely on its duration or another single observational fact has also been questioned. Evidence for a collapsar progenitor has been observed even in the absence of a supernova signature (GRB 060505 and GRB 060614, see e.g. Fynbo et al. 2006b), which we cannot verify for GRB 090426 since the SN would be too faint to detect with current instruments, and a range of high redshift bursts have been detected with short intrinsic duration (see e.g. L\\\"u et al. 2010). Suggestions for how to produce a collapsar with such a short intrinsic duration include viewing the GRB from off-axis, such that we see only a small part of the jet (Lazzati et al. 2009), or the possibility that the central engine had been turned off before the jet reached the surface of the str, with the destruction wave catching up with the jet and leading to a shorter prompt duration (Mizuta \\& Aloy 2009). % No other GRB observed so far has shown such extreme properties of the absorption lines, despite a number of spectra with relatively low column densities. However, none of them had time resolved spectral observations that would allow a line variability study (with the exception of GRB 080310, a potentially similar case to GRB 090426 concerning the properties of the afterglow spectrum; see A. De Cia et al. in prep.). % A larger number of rapid response observations and time series of spectra using a high sensitivity and medium- to high-resolution spectrographs (in order to fit reliable column densities) might allow us to study a few more of those rare events where we directly see the interaction of the GRB with its environment. Extremely rapid response, on the timescale of a few minutes, might even allow us to observe the material ejected by the progenitor itself." }, "1101/1101.1502_arXiv.txt": { "abstract": "The Blazhko effect is a long term, generally irregular modulation of the light curves that occurs in a sizeable number of RR~Lyrae stars. The physical origin of the effect has been a puzzle ever since its discovery over a hundred years ago. We build here upon the recent observational and theoretical work of \\citet{setal} on RRab stars who found with hydrodynamical simulations that the fundamental pulsation mode can get destabilized by a $9:2$ resonant interaction with the 9th overtone. Alternating pulsation cycles arise, although these remain periodic, \\ie not modulated as in the observations. \\\\ Here we use the amplitude equation formalism to study this nonlinear, resonant interaction between the two modes. We show that not only does the fundamental pulsation mode break up into a period two cycle through the nonlinear, resonant interaction with the overtone, but that the amplitudes are modulated, and that in a broad range of parameters the modulations are irregular as in the observations. This irregular behavior is in fact chaotic and arises from a strange attractor in the dynamics. ", "introduction": "A large subclass of RR~Lyrae stars undergo light curve modulations, most of them on the time scale of some 60 periods, although the range extends from some tens to some hundreds of periods. The fundamental pulsation cycle itself lasts ~$\\sim$~0.5\\thi day. The effect was discovered by \\citet{blazhko} over a hundred years ago, and a number of explanations have been proposed, such as closely spaced pulsation modes, a modal $1:2$ resonance, an oblique rotator model, a nonradial modal interaction, and convective cycles \\citep{stothers,molnark}. However, none of these mechanisms is without fault. The giant step toward the explanation of the Blazhko effect has come from the unprecedently precise and continuous {\\sl Kepler} space telescope observations and their analysis \\citep{setal}. Because we are going to build on this work we first present a summary of these findings. Fourteen RR Lyr stars in their sample undergo Blazhko modulations. Unexpectedly, three of these stars also display period doublings, \\ie the shapes of the light curves shows cycle to cycle alternations. The depths of these alternations change during the Blazhko cycle. Another recent observational finding is that the Blazhko cycle does not repeat regularly, see \\eg \\cite{chadid,kolenberg,sodor}. This behaviour poses another important constraint for the physical explanation of the Blazhko effect, as do the observed variations of the mean physical parameters of the stars during the Blazhko cycle \\citep{jurcsik1,jurcsik2}. On theoretical side \\citet{setal} performed a systematic numerical hydrodynamical modeling survey of RR~Lyr models. They found that over a relatively broad region of astrophysical model parameters the fundamental pulsation is unstable and develops into a pulsation with alternating cycles. Guided by earlier work \\citep{mb} that had shown that half-integer resonances can cause a bifurcation to alternating cycles, \\citet{setal} did some sleuthing work. By computing the Floquet stability coefficients \\citep{hartman, bmk} of the fundamental pulsation and searching for resonances of the fundamental mode with successive overtones showed that it is the 9th overtone that is in a $9:2$ resonance and destabilizes the fundamental pulsation cycle. (It is of course a coincidence that it should be the 9th overtone that is in a $9:2$ resonance.) This 9th overtone in these RR~Lyr models turns out to be egregious in that it is a surface mode (dubbed 'strange mode' when it was first encountered by \\citet{strange1} and \\citet{strange2}). The $9:2$ resonance appears in a relatively narrow, winding band in an $\\Log L- \\Log T$ diagram \\citep{KMS}. \\begin{figure*}[ht!] \\begin{center} \\includegraphics[scale=0.9]{plotsig_apjl.ps} \\caption{Temporal modulation of the two amplitudes of the irregularly modulated period two pulsation.} \\end{center} \\label{fig1} \\end{figure*} The numerical hydrodynamical simulations \\citep{KMS} were able to produce alternating cycles, and the ancillary analyses established that the $9:2$ resonance is the cause of the symmetry breaking bifurcation. However, these hydrodynamical simulations were unable to produce either regular or irregular Blazhko like modulations. When a surface mode plays a dynamical role, the numerical hydrodynamical simulations become particularly sensitive to the mixing length parameters, to the zoning in the outer regions as well as to the surface boundary condition. It is therefore extremely difficult to make completely trustworthy and robust simulations. In fact it is an open question whether hydrodynamics with a time dependent mixing length treatment of convection is able to produce Blazhko modulations. For these reasons it is of importance to resort to an alternative, complementary approach as we do in this paper. ", "conclusions": "\\cite{KMS} and \\cite{setal} discovered with the help of numerical hydrodynamical simulations that the Blazhko effect is most likely associated with the half-integer ($9:2$) resonance between the fundamental pulsation mode and an overtone that destabilizes the fundamental RR Lyr full amplitude pulsation. Because of the half integer nature of the resonance ($9P_9 = 2P_0$) it takes therefore two fundamental periods for the pulsation to repeat, hence the occurrence of alternating cycles. The hydrodynamical modeling indeed found this symmetry breaking to give rise to cycle to cycle alternations, \\ie short term variations. The Blazhko effect, however, is a long term irregular amplitude modulation of the pulsations which hydrodyamical simulations have not produced so far. In this paper, using an entirely different approach, namely the amplitude equation formalism, we demonstrate that {\\sl irregular} amplitude modulations can occur quite naturally as a result of the nonlinear, resonant mode coupling between the 9th overtone and the fundamental mode. The phenomenon occurs over a broad range of physical parameters and is therefore quite robust. Furthermore we find that the range of 'periods' of the Blazhko like amplitude modulations are in concordance with the observed ones. It is important to emphasize that the same half integer resonance, responsible for the destabilization of fundamental RR Lyr pulsations, is also capable of producing period doubling or amplitude modulations, depending on the coefficients in the equations. However, observations may have other surprises in store and the Blazhko effect may turn out to be more complicated, so that this simple 2 mode coupling may not account for all its complexity. This is the case for CZ Lacertae (\\cite{sodor}), where the Blazhko-modulation is multi-periodic. There exist, in effect, additional modal resonances or near resonances that may need to be added in the amplitude equation description, and that may then lead to more complicated light curves. Finally, now that the fact that period alternations can occur in classical variable stars such as RR~Lyr stars has been accepted by the observational community, time might be ripe to also target BL~Her stars in which {\\sl theory had actually predicted such alternations} more than 15 years ago \\citep{blher1,blher2,blher3}, but which at the time had been received very skeptically at best. It is true that a recent unpublished Fourier analysis by Buchler and Moskalik of the OGLE data of BL~Her stars does not reveal any conclusive evidence for alternations, but then that was certainly also the case for the pre-{\\sl Kepler} of RR~Lyr." }, "1101/1101.4920_arXiv.txt": { "abstract": "I am mainly interested in the formation and destruction of young star clusters in nearby star forming galaxies such as the Antennae, M83, and M51. One of the first analysis steps is to throw out all those pesky stars that keep contaminating my young cluster samples. Recently, spurred on by our new WFC3 Early Release Science data of galaxies including M83, NGC~4214, M82, NGC~2841, and Cen~A, we began taking a closer look at the stellar component. Questions we are addressing are: 1)~what are the most luminous stars, 2)~how can we use them to help study the destruction of star clusters and the population of the field, 3)~what fraction of stars, at least the bright stars, are formed in the field, in associations, and in compact clusters. In this contribution we describe some of the beginning steps in this process. More specifically, we describe how we separate stars from clusters in our galaxies, and describe how candidate Luminous Blue Variables (LBVs) and ``Single Star'' HII (SSHII) regions have been identified. ", "introduction": "The ability of the Hubble Space Telescope to resolve star clusters in nearby galaxies has lead to rapid growth in this field in the past two decades. For example, the discovery of candidate young globular clusters in NGC~1275 by Holtzman et~al.\\ (1995) and the subsequent studies of the Antennae by Whitmore et~al.\\ (1995, 1999, 2007, 2010a) have shown that it is possible to study the formation of these systems in the local universe rather than having to try to figure out how they formed some 13 Gyr ago. However, most of the work on individual stars has remained focussed on the Milky Way and nearby Local Group galaxies. In this contribution we point out a few examples of what is possible working on individual stars in galaxies more distant than 3~Mpc. One of the first analysis steps for studies of star clusters in galaxies is to identify and then ``throw out'' all the stars that would otherwise contaminate the cluster sample. Recently, spurred on by our new WFC3 Early Release Science (see Chandar et~al.\\ 2010) data of galaxies including M83, NGC~4214, M82, NGC~2841, and Cen~A, we began taking a closer look at the stellar component. This was largely motivated by the improved quality of the $U\\!-\\!B$ vs.\\ $V\\!-\\!I$ diagrams enabled by the increase in the discovery efficiency of WFC3 by a factor of $\\approx$50 over ACS (due to larger field of view) and WFPC2 (due to higher quantum efficiency). An example is shown in Figure~1, which is taken from Chandar et~al.\\ (2010). \\begin{figure} \\plotone{whitmoref1.eps} \\caption{($U\\!-\\!B$) vs.\\ ($V\\!-\\!I$) two-color diagrams for cluster candidates that are outside the nuclear region in M83 (first column of panels), within the nuclear region (middle column of panels), and for stellar candidates throughout the galaxy (based on resolution; along the right column of panels). Each row shows objects in the indicated magnitude range, starting with bright sources at the top and moving to fainter objects at the bottom. The solid line shows predictions in the appropriate WFC3 filters from the cluster models of Bruzual \\& Charlot (2010 -- private communication, but see also Bruzual \\& Charlot 2003). The line of small dots in the panels on the right show the predicted Padova isochrones for individual stars of the appropriate luminosity. See Chandar et al. (2010) for details. \\label{fig:12_plot} } \\end{figure} The quality of the agreement between the Bruzual-Charlot cluster models and the observations is excellent (e.g., upper left panels in Figure~1). This bodes well for the ability to age date clusters using the new generation of WFC3 observations. However, in the context of this meeting, the excellent agreement of the photometric observations of candidate stars (e.g., the bottom right panel) with the Padova stellar models, suggests that the study of individual stars with WFC3 is likely to be equally enhanced. ", "conclusions": "WFC3 data represent a quantum jump in our ability to study cluster populations in nearby galaxies. A byproduct of this is the ability to study individual stars in detail for stars beyond 3~Mpc. In this contributions we: \\begin{enumerate} \\item Demonstrate how both size and color information can be used to separate stars and clusters. \\item Find a population of stars in M83 that may be LBVs. \\item Identify a population of ``Single Star'' HII regions (SSHII) in M83. Many of these are in the field, showing that not all stars form in clusters. \\end{enumerate} One might ask the question how much beyond 3~Mpc is it possible to perform similar studies. Without trying to set a limit, we simply note that it is straightforward to observe individual young stars in galaxies such as the Antennae galaxies ($\\approx$20 Mpc), as shown in Figure~5 (from Whitmore et~al.\\ 2010a)." }, "1101/1101.2671_arXiv.txt": { "abstract": "{The measurement of Doppler velocities in spectroscopic solar observations requires a reference for the local frame of rest. The rotational and radial velocities of the Earth and the rotation of the Sun introduce velocity offsets in the observations. Normally, good references for velocities are missing (e.g. telluric lines), especially in filter-based spectropolarimetric observations. }% {We determine an absolute reference for line-of-sight velocities measured from solar observations for any heliocentric angle, calibrating the convective line shift of spatially-averaged profiles on quiet sun from a 3D hydrodynamical simulation. This method works whenever there is quiet sun in the field-of-view, and it has the advantage of being relatively insensitive to uncertainties in the atomic data. }% {We carry out radiative transfer computations in LTE for selected \\ion{C}{i} and \\ion{Fe}{i} lines, whereas the \\ion{Ca}{ii} infrared lines are synthesized in non-LTE. Radiative transfer calculations are done with a modified version of \\textsc{Multi}, using the snapshots of a non-magnetic 3D hydrodynamical simulation of the photosphere.} % {The resulting synthetic profiles show the expected C-shaped bisector at disk center. The degree of asymmetry and the line shifts, however, show a clear dependence on the heliocentric angle and the properties of the lines. The profiles at $\\mu=1$ are compared with observed profiles to prove their reliability, and they are tested against errors induced by the LTE calculations, inaccuracies in the atomic data and the 3D simulation.}% {Theoretical quiet-sun profiles of lines commonly used by solar observers are provided to the community. Those can be used as absolute references for line-of-sight velocities. The limb effect is produced by the projection of the 3D atmosphere along the line of sight. Non-LTE effects on \\ion{Fe}{i} lines are found to have a small impact on the convective shifts of the lines, reinforcing the usability of the LTE approximation in this case. We estimate the precision of the disk-center line shifts to be approximately 50~m~s$^{-1}$, but the off-center profiles remain to be tested against observations.}% ", "introduction": "The need for high spatial and spectral resolution in solar observations stimulates new advances in telescope and instrumentation technology, adaptive optics and image reconstruction techniques. However, the accuracy of the measurements is often limited by calibration issues. This is evident when using Doppler shifts of spectral lines to acquire line-of-sight velocities in the solar plasma, where the task of finding a local standard of rest is often problematic. The current work is partly motivated by the installation of the CRisp Imaging Spectropolarimeter \\citep[CRISP,][]{2008ApJ...689L..69S} at the Swedish 1-m Solar Telescope \\citep[SST,][]{2003SPIE.4853..341S}. Instruments like this usually lack laboratory wavelength references, while the use of telluric lines has limitations as discussed below. The obvious (and commonly used) solution of letting the spectral line of interest -- averaged over a large enough area and long enough time -- define a local frame of rest is confounded by the fact that convective motions leave strong fingerprints on any spectral line formed in the solar photosphere. The statistical average of bright blueshifted and dark redshifted profiles formed in granules and intergranular lanes respectively, produces the blueshifted C-shaped bisectors of photospheric lines as reviewed by \\citet{1982ARA&A..20...61D}. Below, we discuss some methods that have been used to define a velocity reference for solar observations. \\begin{itemize} \\item \\emph{Sunspot umbrae} are common standards \\citep[e.g.,][]{1977ApJ...213..900B,2008ApJ...689L..69S,2010ApJ...713.1282O}, but this only works when a suitable sunspot is in the field-of-view and relies on the assumption that the umbra is at rest because of the convective motions being suppressed by the strong magnetic fields. Furthermore, the line must be measurable in the umbra and wavelength shifts which are induced by blends from lines uniquely formed in the umbra (e.g. molecules), must be accounted for. Finally, observations of umbrae are always plagued by stray light from the much brighter quiet photosphere. \\item \\emph{Telluric lines} can be used to define a laboratory-frame reference, which can then be converted to the Sun. \\citet{1997ApJ...474..810M} and \\citet{2008ApJ...676..698B} used telluric lines to derive the wavelength scale for their observations, then converted to an absolute wavelength scale on the Sun given the ephemeris constants, time of the observations, solar rotation and the laboratory wavelengths for all observed spectral lines. \\item \\emph{A spectral atlas}, most commonly the atlas acquired with the Fourier Transform Spectrometer at the McMath-Pierce Telescope (hereafter called the FTS atlas) of \\citet{fts-atlas}, can be used to calibrate observations. This was done by \\citet{2007ApJ...655..615L}, who used the laboratory wavelength to define the reference for velocities, assuming an absence of systematic errors in the FTS atlas itself and correcting for the gravitational redshift (633~$\\mbox{m}\\ \\mbox{s}^{-1}$). This method only works at disk center where the atlas was recorded. \\item \\emph{Numerical models} can be used to predict convective line shifts and thus provide a velocity reference. A two-component model of the solar photosphere has been derived by \\citet{2002A&A...385.1056B} from the inversion of \\ion{Fe}{i} spectral lines. Since it does not contain any horizontal velocities, it then is in principle limited to calibrating disk-center observations, for which it has been used by \\citet{2004A&A...415..717T} and \\citet{2009A&A...508.1453F}. However, \\citet{2004A&A...427..319B} used it with the empirical results of \\citet{1988A&AS...72..473B} to also estimate lineshifts off solar center. \\citet{2007ApJ...655..615L} employ a 3D numerical granulation model to compute the convective blueshift of the \\ion{C}{i}~5380~\\AA\\ line and calibrate their observations. This approach was necessary because the laboratory wavelength of the \\ion{C}{i} line is not known with enough precision to use the atlas calibration that is mentioned above. \\end{itemize} We note that telluric lines cannot always be observed with the solar diagnostic lines and that when this is the case, observing them can be expensive. Observations with tunable filters require high cadence because of seeing variability and solar evolution. Any time spent recording telluric lines thus decreases the spatial resolution and the signal-to-noise of the science data. This would also apply to any laboratory spectral-line source used for calibration. Furthermore, approaches using telluric lines (or other calibration lines) and spectral atlases all require very accurate atomic data to define the zero velocity point and thus convert the wavelength scale into velocity scale. Such data are unfortunately not available for many of the lines that are extensively used in solar physics as is apparent in the work of \\citet{2007ApJ...655..615L} described above. Their approach is proposed in this study: using realistic 3D numerical simulations of the solar photosphere to compute spatially-averaged line profiles. As the radiative transfer computation is carried out in the local frame of rest of the Sun and the adopted atomic data is known, the velocity shifts of the spatially-averaged profile can be computed accurately for different heliocentric angles, relative to the assumed center of the line. This method relies on the degree of realism of the 3D simulation and requires a large statistical sample of spectra to compute the spatially-averaged profile. \\\\ The aim of this work is to provide the community with a set of theoretical spatially-averaged quiet-Sun profiles of lines commonly used in solar physics, that are computed for a range of heliocentric angles. In Sect.~\\ref{llist} the atomic data and radiative transfer details are discussed. The 3D HD simulation is described in Sect.~\\ref{hydromod}. In Sect.~\\ref{res} the results are presented together with an error analysis and a real filter-based observation is calibrated using the new data. Section~\\ref{edata} describes the electronic data made available and Sect.~\\ref{conclusions} summarizes the main conclusions of this work. ", "conclusions": "We made use of the approach described by \\citep{2007ApJ...655..615L} to obtain an absolute reference for velocities in solar observations. For this purpose, we calibrated the shift in synthetic spectra produced by convective motions at different heliocentric angles for \\ion{C}{i}, \\ion{Fe}{i}, and \\ion{Ca}{ii} lines as commonly used for solar physics diagnostics. We propose to use the profiles as absolute local references for line-of-sight velocities to allow accurate velocity measurements. A realistic 3D hydrodynamical simulation of the non-magnetic solar photosphere was used to synthesize the spectra. The computations were done in LTE for \\ion{C}{i} and \\ion{Fe}{i} lines whereas non-LTE computations were carried out in the \\ion{Ca}{ii} infrared lines. As the conversion from wavelength to velocity scale is done using the same laboratory wavelength as used for computing the profiles, the measured line shifts are relatively unaffected by errors in the central wavelength. Uncertainties in the rest of atomic parameters are partially compensated when the profiles are fitted to reproduce the strength of the FTS atlas. Our results show a systematic center-to-limb variation of the bisectors, which is determined to a large extent by the 3D structure of the atmosphere, which is projected along the line-of-sight. However, there is a small amount of redshift (relative to $\\mu=1.0$) produced by the change in the formation height with the heliocentric angle (non-3D effects), which is not related to projection effects. A careful analysis of the models suggests that a large number of 3D snapshots must be used in studies of line shifts so the oscillations are properly sampled and the statistics of granules and intergranules are complete, whereas non-LTE effects seem to have a minor influence on the shift of our \\ion{Fe}{i} lines. The largest uncertainty is related to the non-magnetic nature of the simulation. Nevertheless, we estimate our calibrations to be accurate up to $\\sim50 \\ \\mbox{m}\\ \\mbox{s}^{-1}$ at $\\mu=1$, remaining untested towards the solar limb. We would like to emphasize the somewhat limited usability of the \\ion{Ca}{ii} calibration, as observations are polluted by a significant number of blends, which should be masked in order to compare with the calibration data. Our synthetic profiles reproduce spatially-averaged solar observations accurately, but it would be interesting to test for the impact of magnetic fields on the convective shift of the lines, and confirm the reliability of our calculations near the limb." }, "1101/1101.2447_arXiv.txt": { "abstract": "{The most accepted scenario for the origin of fossil groups is that they are galaxy associations in which the merging rate was fast and efficient. These systems have assembled half of their mass at early epoch of the Universe, subsequently growing by minor mergers, and therefore could contain a fossil record of the galaxy structure formation.}{We have started an observational project in order to characterize a large sample of fossil groups. In this paper we present the analysis of the fossil system RX J105453.3+552102.}{Optical deep images were used for studying the properties of the brightest group galaxy and for computing the photometric luminosity function of the group. We have also performed a detail dynamical analysis of the system based on redshift data for 116 galaxies. Combining galaxy velocities and positions we selected 78 group members.}{RX J105453.3+552102 is located at $\\left=0.47$, and shows a quite large line--of--sight velocity dispersion $\\sigma_{v}\\sim 1000$ \\kss. Assuming the dynamical equilibrium, we estimated a virial mass of $M(2$ in the $R$-filter) between the two brightest members of the system within half of its virial radius. These galaxy associations also show an extended bright X-ray emission ($L_{\\rm X}>10^{42}$ $h_{50}^{-2}$ erg s$^{-1}$) surrounding the brightest group galaxy. According to this definition, these systems are as common as poor and rich galaxy clusters together ($n\\sim (1-4) \\times 10^{-6}$ $h_{50}^{-3}$ Mpc$^{-3}$; Vikhlinin et al. \\cite{vikhlinin98}; Jones et al. \\cite{jones03}; Santos et al. \\cite{santos07}; La Barbera et al. \\cite{labarbera09}; Voevodkin et al. \\cite{voevodkin10}). Numerical simulations show that FGs could be particular cases of structure formation. Thus, according to these simulations, FGs have been formed inside highly concentrated DM halos at an early epoch of the Universe, assembling half of their dark matter mass at $z>1$, and subsequently growing by minor mergers. In contrast, non-fossil groups show, on average, a later formation (D'Onghia et al. \\cite{donghia05}; von Benda-Beckmann et al. \\cite{vonbenda08}). This early formation leaves enough time for $L^{*}$ galaxies to merge into a massive elliptical-type galaxy located at the center of the group, producing a lack of intermediate-luminosity galaxies and a large magnitude gap between the brightest and the second brightest galaxy of the group. FGs also have special dynamical properties which speed up the merging efficiency. In particular, the in-fall of massive satellites in FGs took place on orbits with low angular momentum, which might be the main responsible of the anisotropy of the group galaxies, in such a way that groups with highly radially anisotropic velocity distributions tend to become fossil (Sommer-Larsen \\cite{somer06}). Simulations also indicate that FGs have only been able to accrete on average one galaxy since $z\\sim 1$, compared to $\\sim 3$ galaxies for normal groups (see von Benda-Beckmann et al. \\cite{vonbenda08}). This means that FGs provide unique clues on the history of cosmic mass assembly and the relationship between baryons and their host halos. They also could have a fossil record of the structure formation of galaxies at early epochs of the Universe. Observations are broadly in agreement with the formation framework of FGs proposed by numerical simulations. Thus, Khosroshahi et al. (\\cite{khosro07}) compared the scaling relations of a sample of FGs and non-fossil systems and found that FGs follow the X-ray luminosity-temperature relation ($L_{\\rm X}$-$T_{\\rm X}$) as clusters and groups. However, there are significant differences in the optical vs. X-ray luminosities ($L_{\\rm opt}-$$L_{\\rm X}$), X-ray luminosity vs. cluster velocity dispersion ($L_{\\rm X}-\\sigma$) and X-ray temperature vs. cluster velocity dispersion ($T_{\\rm X}-\\sigma$) relations. In particular, for a given $\\sigma$, FGs are located in more luminous and hotter X-ray halos than normal groups and clusters. They also have larger X-ray luminosities than normal groups for a given $L_{\\rm opt}$ (but see also Voevodkin et al. \\cite{voevodkin10}). These differences could be due to an early formation epoch of FGs as suggested by simulations (Khosroshahi et al. \\cite{khosro07}). Detailed X-ray observations of some FGs also indicate that these systems were assembled at early epochs in high centrally concentrated DM halos with large mass-to-light-ratio ($M/L$) relations. Nevertheless, they do not show cooling cores as those detected in galaxy clusters, which points toward the presence of other heating mechanisms, like AGN feedback (Sun et al. \\cite{sun04}; Khosroshahi et al. \\cite{khosro04},\\cite{khosro06}; Mendes de Oliveira et al. \\cite{mendes09}). The absence of recent galaxy or cluster major mergers together with the lack of cool cores make FGs the ideal objects to study the effects of AGN feedback and the link between galaxy evolution and intra-group medium (IGM). Optical and near-infrared observations indicate that the faint-end slope ($\\alpha$) of the luminosity function (LF) of FGs spans a wide range of values. Thus, the Schechter function fitted to the LF of these systems shows values in the range -1.6$<\\alpha<$-0.6 (Cypriano et al. \\cite{cypriano06}; Khosroshahi et al. \\cite{khosro06}; Mendes de Oliveira et al. \\cite{mendes06}, \\cite{mendes09}). This suggests that some FGs are dwarf rich systems like similar size/mass galaxy clusters, while others show a lack of dwarf galaxies. It has been pointed out that these differences between fossil and non-fossil systems can reflect different substructure distribution (Jones et al. \\cite{jones00}). Thus, FGs could have one order of magnitude less substructure with respect to the standard cosmological model predictions (D'Onghia \\& Lake \\cite{donghia04}). Nevertheless, the number of LFs measured for FGs is scarce and the system where this different substructure was measured has only 40$\\%$ of the Virgo mass. It should be pointed out the fact that many systems classified in the past as FGs turned to be fossil clusters. On mass scale of groups it is not completely clear when the transition from galaxy formation to galaxy cluster formation happens. The low mass FGs are intermediate systems in this respect and can give hints of how and at which extent the substructures are accreted. Thus, studies on low mass FGs might give a hint on the abundance of dwarf galaxies in systems with mass scale intermediate between a galaxy and a galaxy cluster as compared to the standard cosmological predictions. The brightest group galaxies (BGGs) located at the center of FGs are among the most massive galaxies known in the Universe. They contain the key for understanding the formation and evolution of FGs. Observations show that BGGs have also different observational properties than other bright elliptical (E) galaxies. In particular, they present discy isophotes in the center and their luminosity correlates with the velocity dispersion of the group (Khosroshahi et al. \\cite{khosro06}). These different properties suggest a different formation scenario for bright Es in fossil and non-fossil systems. While bright Es in FGs would grow by gas-rich mergers, giant Es in non-fossil systems would suffer more dry mergers. However, recent samples of BGGs do not find these differences (La Barbera et al. \\cite{labarbera09}). All previous results have the drawback that they were obtained using small samples of FGs. This could be the reason of some contradictory results found by different studies. The lack of a large and homogeneous statistical study of this kind of systems make the previous results not conclusive. A systematic study of a large sample of FGs remains to be done. \\subsection{Fossil Groups Origins (FOGO) project.} We have started a large observing program on FGs. The aim of this project is to carry out a systematic, multiwavelength study of a sample of 34 FGs selected from the Sloan Digital Sky Survey (SDSS; Santos et al. \\cite{santos07}). This sample is ideal for providing strong constraints on the observational properties of the galaxy populations in FGs due to its unique characteristics. The sample spans the last 5 Gyr of galaxy evolution ($02$. The fossil group classification is clearly strongly dependent on the magnitude estimation of the bright galaxy group. Unfortunately, this is not an easy task due to the BGG is often located in high density galaxy environments. Thus, there is a large difference between SDSS model ($m_{r,model}=17.69$) and Petrosian magnitudes ($m_{r,petro}=18.10$) for the BGG of RX J105453.3+552102. From our photometry, the magnitude of the BGG calculated by SExtractor is $m_{r,SEx}=18.08$, and the magnitude obtained from its surface brightness fit is $m_{r,fit}=17.49$. Notice the agreement between $m_{r,petro}$ and $m_{r,SEx}$ and between $m_{r,model}$ and $m_{r,fit}$. Nevertheless, the model magnitudes are always brighter than aperture ones because they are computed integrating until infitive radius the fitted surface brightness profiles of the galaxies. The 0.2 mag difference between $m_{r,model}$ and $m_{r,fit}$ could be due to the best S\\'ersic fitted model by SDSS has $n=1$, while our best fitted model has $n\\sim 2$. The differences between model and aperture magnitudes are crucial for comparing magnitudes of the same class. Thus, when considering our $m_{r,SEx}$ magnitudes of the cluster galaxies we obtained a magnitude gap between the BGG and the second rank galaxy within 500 kpc radius of $\\Delta m_{12}=1.92\\pm0.09$. This magnitude gap can be seen in Fig.~\\ref{colormag2}. The value of 1.92 is very close to the classical FG definition given by Jones et al. (2003), but does not allow to classify RX J105453.3+552102 as a fossil group. Taking into account the errors there is a probability of $\\sim 20\\%$ to have $\\Delta m_{12}>2$. We also computed from our photometry the model magnitudes of the second brightest galaxies of the group within 500 kpc radius. In this case $\\Delta m_{12}=1.87\\pm0.15$. Recently, Dariush et al. (\\cite{dariush10}) have proposed another photometrical definition of FGs based on the magnitude gap between the BGG and the fourth ranked galaxy ($\\Delta m_{14}$). Analysing groups and clusters of galaxies using the Millennium Simulation, they found that early-formed galaxy association are better identified as those showing $\\Delta m_{14}>2.5$ mag. In our case, the RX J105453.3+552102 group has $\\Delta m_{14}=2.47\\pm0.09$ mag (using our $m_{r,SEx}$) and, again, this group cannot be classified as fossil (see Fig.~\\ref{colormag2}). In this case, the probability that the system has $\\Delta m_{14}>2.5$ is $\\sim 35\\%$. The same value of $\\Delta m_{14}$ was obtained when $m_{r,fit}$ of the 4th brightest galaxy of the group within 500 kpc was considered. As shown above, the classification of a system as fossil or not can be quite sensible to the BGG magnitude estimation or to the presence of bright interlopers in the cluster field. More in general, the classification scheme of a fossil group might be improved in several ways, e.g., taking into account a radius scaling with $R_{\\rm 200}$ and a magnitude band changing with redshift. However, the discussion of this scheme is out of the aims of FOGO project which are rather to check of how many groups in the catalog of Santos et al. (\\cite{santos07}) actually have ''fossil'' nature and to study their properties. As for RX J105453.3+552102, its real nature, the likely past dynamical history, and properties are discussed in the next sections. \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{15364fg14.ps}} \\caption{ Distance to the group center vs SExtractor r-band magnitude for all galaxies (grey circles). We have also overplotted cluster members (red triangles) and non-cluster members (black circles). The horizontal line shows 500 kpc distance from the galaxy group center (the distance used in the ``fossil group'' definition). The vertical dotted and dashed lines show $\\Delta m_{12}=2.0$ and $\\Delta m_{14}=2.5$ respectively, taking $m_{r,SEx}$ as the BGG magnitude.} \\label{colormag2} \\end{figure} \\subsection{The dynamical state of the cluster} The presence of a large magnitude gap has been always taken as indication of relaxed and early-formed galaxy systems. Nevertheless, in the Millennium Simulation can be seen that most of the early-formed systems do not show large magnitude gaps (see Dariush et al. \\cite{dariush10}). Thus, overcoming the empirical definitions of ``fossil group'', one should consider whether RX J105453.3+552102 is or is not an old and undisturbed system that has underegone little infall of $L^*$ galaxies since its initial collapse. Recent major mergers with other galaxy systems would yield some observable smoking-guns. The first would be the presence of substructure in RX J105453.3+552102. We have used a battery of different tests in 1D, 2D and 3D to take into account the geometry of a possible cluster merger (Pinkney et al. \\cite{pin96}). We found no evidence of substructure. The only possible hint is the peculiar velocity of the BGG galaxy (significant at the $>95\\%$ c.l.), which is often connected to evidence of substructure (e.g. Bird \\cite{bir94}). However, the velocity of the BGG in the cluster rest frame is only $\\sim 300$ \\ks and the relative peculiar velocity with respect to the cluster velocity dispersion is 0.3, which is not a particularly large value among clusters (see fig.~2 of Coziol et al. \\cite{coz09}). Moreover, RX J105453.3+552102 seems well isolated in the phase-space as show by Fig.~\\ref{figvd} (see den Hartog \\& Katgert \\cite{den96} and Aguerri et al. \\cite{aguerri07} for other clusters). This supports the idea that RX J105453.3+552102 is far from an important accretion episode. Nevertheless, we have shown that, albeit the overall velocity distribution of the spectroscopical galaxy sample belonging to RX J105453.3+552102 is Gaussian, there is a significant departure from Gaussianity in the outer regions ($R>2'$) which we have interpreted as a possible signature of radial anisotropy of the galaxies in the group outskirts. If confirmed in more detailed dynamical analysis, this will represent an interesting piece of information to be added into the fossil group information scenarios (D'Onghia et al. 2005; Sommer-Larsen 2006). The second signature of an old and undisturbed cluster comes from the BGG itself. In fact, RX J105453.3+552102 BGG shows a small $n$ S\\'ersic index value ($n\\approx2$) and clear discy isophotes in the external regions. According to the findings of numerical simulations, surface brightness profiles with small $n$ values result from gas rich mergers (e.g., Khochfar \\& Burkert \\cite{kho05}). If the RX J105453.3+552102 BGG has indeed been formed from the merger of all major galaxies within the inner regions of the system in very early times, then some of these mergers would have been gas--rich. In contrast, the merger between two clusters and the following equal--mass dry mergers of the corresponding dominant ellipticals would produce surface brightness profiles with $n\\approx4$. Another piece of evidence in favour of the above scenario is the absence of multiple nuclei both from our photometric data and from spectroscopic data. I fact, we took three spectra with different slit position angles crossing the BGG nucleus and giving equal $z$ values. The third, possible piece of evidence that RX J105453.3+552102 is a relaxed old cluster comes from the dip observed in the photometric LF at $M_{r}\\sim-19.5$. This dip is more prominent in the LF$_{\\rm phot}$ computed with the galaxies located within 0.5 Mpc radius. Similar bimodality in the LF has also been reported in other relaxed clusters (Yagi et al. \\cite{yagi02}) and groups (Hunsberger et al. \\cite{huns98}, Miles et al. \\cite{mil04}, Mendes de Oliveira et al. \\cite{mendes06}). We can conclude that although RX J105453.3+552102 do not follow the prescription of a fossil system, is a relaxed and old galaxy cluster with no indication of recent infall of $L^{*}$ galaxies and therefore, it is a genuine FG. \\subsection{RX J105453.3+552102 a relaxed and massive system at z$\\approx 0.5$} Our analysis indicates that RX J105453.3+552102 is a very massive galaxy cluster already relaxed at z$\\approx$0.5. This means that $\\approx 6$ Gyr ago this cluster was as massive as the Coma cluster but more dynamically evolved. In a hierarchical structure formation scenario, these very massive, relaxed systems are likely formed at low redshift. In order to understand how common is a cluster like RX J105453.3+552102, we have searched in the Millennium Simulation (Springel et al. \\cite{springel05}; Boylan-Kolchin et al. \\cite{boylan09}) for halos with masses M$_{200}>1\\times10^{15} h_{70}^{-1}$ M$_{\\odot}$ located at z$=0.5$. The total number of such halos was only 9 within the volume covered by the Millennium Simulation (having a box of comoving size of 500 $h_{100}^{-1}$ Mpc). We have also computed the magnitude difference between the brightest and the second brightest galaxies located in each of those halos: only 3/9 have $\\Delta m_{12}>2$. However, the Millennium Simulation has been carried out using a normalization of the power spectrum with $\\sigma_8=0.9$. Using instead a lower normalization, $\\sigma_8=0.8$, in agreement with the most recent CMB and large-scale structure analysis (e.g. Komatsu et al. 2010, and references therein), the number density of such massive clusters at $z=0.5$ drops by about a factor of three. Thus we expect the number of ``fossil'' clusters at least as massive as RX J105453.3+552102 within the Millennium Simulation volume at $z=0.5$ to be of order unity. Clearly, to decide whether the detection of a relaxed massive cluster like RX J105453.3+552102 should be considered as a rare event for a standard $\\Lambda$CDM cosmology, one has to know the volume within which this cluster has been found, i.e. the selection function of the corresponding survey. We postpone the discussion of this issue to a future analysis, which will also be based on a larger statistics of fossil groups/clusters." }, "1101/1101.0849_arXiv.txt": { "abstract": "We present a study of galaxy environments to z$\\sim$2, based on a sample of over 33,000 K-band selected galaxies detected in the UKIDSS Ultra Deep Survey (UDS). The combination of infrared depth and area in the UDS allows us to extend previous studies of galaxy environment to $z>1$ without the strong biases associated with optical galaxy selection. We study the environments of galaxies divided by rest frame $(U-B)$ colours, in addition to `passive' and `star-forming' subsets based on template fitting. We find that galaxy colour is strongly correlated with galaxy overdensity on small scales ($<1$~Mpc diameter), with red/passive galaxies residing in significantly denser environments than blue/star-forming galaxies to $z\\sim 1.5$. On smaller scales ($<0.5$~Mpc diameter) we also find a relationship between galaxy luminosity and environment, with the most luminous blue galaxies at $z\\sim1$ inhabiting environments comparable to red, passive systems at the same redshift. Monte Carlo simulations demonstrate that these conclusions are robust to the uncertainties introduced by photometric redshift errors. ", "introduction": "It has long been known that the properties of galaxies depend on the environment in which they are located. Elliptical, non-star-forming galaxies occupy more dense regions of space than star-forming, disc-dominated galaxies, giving rise to the so-called morphology-density relation (\\citealt{Oemler}; \\citealt{Dressler}). The physical origin of this relation is still subject to debate, with disagreement mainly centering on whether the relation arises due to internal or external processes (nature vs nurture). Most recent low redshift studies (\\citealt{Kauffmann}; \\citealt{Balogh}) utilise the Sloan Digital Sky Survey (SDSS) or the Two-degree-Field Galaxy Redshift Survey (2dFGRS) to conduct statistical investigations of galaxy environments. \\citet{Kauffmann} constrained the specific star formation rate (SSFR) using the 4000\\AA ~break and found that the SSFR (and nuclear activity) depend most strongly on local density, from star-forming galaxies at low densities to predominantly inactive systems at high densities. \\par Studies by \\citet{van Der Wel} and \\citet{Bamford} found that structure, colour and morphology are mainly dependent on galaxy mass but that at fixed mass, colour and, to a lesser extent, morphology are sensitive to environment. Studies of H$\\alpha$ \\citep{Balogh} found that it's strength does not depend on environment but that the fraction of galaxies with equivalent width, $W_0(H\\alpha)>$4\\AA ~is environmentally dependant, decreasing with increasing density. They also noted that emission line fraction appears to depend on both the local environment ($\\sim$ 1Mpc) and on the large scale structure ($\\sim$ 5Mpc). Studies at higher redshifts (z $\\sim$ 1) have used surveys such as DEEP2 \\citep{Davis} and VVDS \\citep{LeFevre}. DEEP2 investigations (\\citealt{Cooper06}; \\citealt{Cooper07}) used the projected third nearest neighbour statistic, studying galaxy properties and the colour-density relation respectively. They concluded that there is a strong dependence on rest frame $(U-B)$ colour, with blue galaxies occupying lower density regions but showing a strong increase in mean local density with luminosity at $z\\sim 1$. This they conclude is consistent with the rapid quenching of star formation by AGN or supernova feedback, as ram pressure stripping, harassment and tidal interactions, which occur preferentially in clusters, would be insufficient to explain these findings. \\cite{Cooper07} also observed that the fraction of galaxies on the red sequence increases with local density, as in the local Universe, but this weakens with redshift and disappears by $z\\sim1.3$. The VIMOS VLT Deep Survey (VVDS ) investigated the redshift and luminosity evolution of the galaxy colour-density relation up to z$\\sim$1.5 \\citep{Cucciati}. In agreement with \\cite{Cooper07} they found that the local colour-density relation progressively weakens and possibly reverses in the highest redshift bin (1.2$<$z$<$1.5). This may imply that quenching of star formation was more efficient in high density regions. The VVDS team also observed that the colour-density relations depend on luminosity and found that at fixed luminosity there is a decrease in the number of red objects as a function of redshift in high density regions. This implied that star formation ends at earlier cosmic epochs for more luminous/massive galaxies, which is consistent with downsizing \\citep{Cowie}. We note, however, that the VVDS survey is based on optical $I$-band selection, and as such will be strongly biased against red, passive galaxies at $z>1$. Conclusions from deep K-selected samples suggest that the galaxy colour bimodality is present to at least z$\\sim$1.5 (e.g. \\citealt{Cirasuolo}) and may be still be present at z$\\sim$2 (\\citealt{Cassata08}; \\citealt{Kriek08}; \\citealt{Williams09}). Furthermore red galaxies have been seen to strongly cluster at z$>$1.5 (\\citealt{Daddi03}; \\citealt{Quadri07}; \\citealt{Hartley08}; \\citealt{Hartley}) which suggests that a colour-density relation may also exist at these higher redshifts. A number of physical processes may be responsible for the observed environmental trends. Mergers or tidal interactions can tear galactic discs apart and are likely to play an important role in forming the most massive galaxies observed today (\\citealt{Toomre}; \\citealt{Farouki}). Other processes such as gas stripping can severely reduce the star formation rate by removing the cold gas from galaxies falling into massive dark-matter halos (\\citealt{Gunn}; \\citealt{Dekel}). Feedback is also thought to play a major role, either from AGN or supernovae (e.g. \\citealt{Benson05}; \\citealt{Springel}). These processes may heat or eject the gas within galaxies and thus effectively terminate any further star formation, which can rapidly lead to the build-up of the galaxy red sequence. Finally infalling cold gas in low mass dark matter halos may fall directly onto the galaxy, whereas in high mass halos the gas is thought to be heated by shocks and therefore remains supported (\\citealt{White}; \\citealt{Birnboim}). A key goal of observational extragalactic astronomy is to disentangle which of these processes are responsible for establishing the bimodal galaxy populations observed in the local Universe. With the recent advent of deep, wide-field infrared imaging in the UKIDSS UDS we can now extend studies of galaxy environments to z$>$1. Selection in the infrared avoids the major biases against dusty and/or evolved stellar populations, allowing us to investigate whether correlations observed at low redshift also occur at high redshift and how these change over time. The large contiguous area of this survey also allows us to probe a wide range of environments using large samples of galaxies. The paper is structured as follows: \\S 2 outlines the data and selection criteria used in this work. \\S 3 discusses the method used to estimate galaxy environments. The results are then presented in \\S 4 and \\S 5, with \\S 6 summarising our conclusions. Throughout this paper we assume a $\\Lambda$CDM cosmology with $\\Omega_{m}$=0.3, $ \\Omega_{\\Lambda}$=0.7 and $H_{0}$=71 km s$^{-1}$ Mpc$^{-1}$. \\section[]{Data and sample selection} \\subsection{The UKIDSS Ultra Deep Survey} \\hspace{5mm} This work has been performed using the third data release (DR3) of the UKIRT (United Kingdom Infra-Red Telescope) Infrared Deep Sky Survey, Ultra-Deep Survey (UKIDSS, UDS; \\citealt{Lawrence}; Almaini et al in prep). The UKIDSS project consists of 5 sub-surveys of which the UDS is the deepest, with a target depth of K=25 (AB) over a single 4-pointing mosaic of the Wide-field camera (WFCAM, \\citealt{Casali}), giving the UDS an area of 0.88 x 0.88 degrees. The 5$\\sigma$, AB depths within 2$''$ apertures for the J, H and K-bands are 23.7, 23.5 and 23.7 respectively for the DR3, making it the deepest near infrared survey over such a large area at the time of release. For details of the stacking procedure, mosaicing, catalogue extraction and depth estimation we refer the reader to Almaini et al. (in prep.) and \\citet{Foucaud}. The field is also covered by deep optical data in the B, V, R, i$^{\\prime}$ and z$^{\\prime}$ -bands with depths of $B_{AB}$=28.4, $V_{AB}$=27.8, $R_{AB}$=27.7, $i'_{AB}$=27.7 and $z'_{AB}$=26.7 from the Subaru-XMM Deep Survey (SXDS) \\citep[$3\\sigma$, $2\\arcsec$ diameter]{Furusawa}. Data from the {\\it Spitzer} Legacy Program (SpUDS, PI:Dunlop) reaching 5$\\sigma$ depths of 24.2 and 24.0 (AB) at 3.6$\\mu$m and 4.5$\\mu$m respectively and U-band data from CFHT Megacam ($U_{AB}$=25.5; Foucaud et al. in prep) are also utilised, which results in a co-incident area of 0.63 deg$^2$ after masking. \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=250pt]{UBvsMk_075z175_HistSubplot_AB.eps} \\caption[$(U-B)$ vs absolute K-band magnitude between $0.7515$) are removed from our sample as these are likely to be unreliable. This removes 4$\\%$ of the galaxy sample, the majority of these are either QSOs (36\\%), cross-talk (26\\%) or the minor members of pairs or mergers (23\\%), with the remainder consisting largely of objects with very low surface brightness. The fraction of otherwise useful objects rejected is therefore $<$0.6$\\%$. \\subsection{Passive Sample} To define a passive galaxy subset with minimal contamination from dusty star forming objects we use a subset of galaxies outlined in \\cite{Hartley}. Templates were used to fit either an instantaneous burst parameterised by an age, or an exponentially decaying star-formation rate parameterised by an age and $\\tau$, the e-folding time in the exponentially declining star-formation rate, such that, \\begin{equation} SFR = SFR_{0} \\times e^{-age/\\tau} \\end{equation} where $SFR$ is the star-formation rate at the time of observation and $SFR_{0}$ was the initial value. We define a conservative passive sample as galaxies that are simultaneously old (age$>$1Gyr) and have ongoing star formation with $SFR\\le 0.1 \\% $ of $SFR_0$, and a star forming sample with $SFR\\geq10\\% $ of $SFR_0$, with 3947 and 22,158 galaxies in each sample respectively.\\par To define the red sequence we performed a $\\chi^2$ minimisation to fit an equation of the form $(U-B) = a \\times M_K + b$~ to the old, burst galaxies, defining the red sample to be all galaxies within 3$\\sigma$ of this fit (see \\cite{Hartley} for a more detailed description). In this work, to separate red and blue galaxies, we use the red-sequence slope from \\citet{Hartley} but choose the division between the two populations to fit the minimum in the overall colour bimodality (as shown in the histogram in Figure \\ref{UBMK}). This leads to a dividing line in the colour-magnitude diagram as follows: \\begin{equation} (U-B)=-7.09 \\times 10^{-3} M_K + 0.52 \\end{equation} This boundary was found to separate the red and blue populations effectively to z$\\sim$1.75. At higher redshift the bimodality in galaxy colours is less clear, which may in part be due to photometric errors. A full examination of this issue and the evolution of the red sequence will be presented in Cirasuolo et al. (in prep). Previous studies have found evidence for an evolution in the location of the red sequence with redshift (e.g \\citealt{Brammer}). For simplicity we choose not to model the red sequence in such detail and instead use the fixed colour selection boundary given above. We note, however, that using an evolving boundary made no significant difference to any of the conclusions presented in this work. Table \\ref{NGal} shows the resulting number of red and blue galaxies assigned to each photometric redshift bin, including the conservative subsamples of passive and actively star-forming galaxies. \\section[]{Environmental Measurement} We used two methods to calculate galaxy environment: Counts in an Aperture and $n$th Nearest Neighbour. In both methods all the galaxies within a photometric redshift bin are collapsed down onto a 2D plane and the redshift information within the bin is not utilised any further. In the aperture method apertures of 1Mpc, 500kpc and 250kpc diameter (physical) are placed on each galaxy and the number of galaxies within that aperture are counted (N$^{Aper}_{g}$). A sample of $\\sim$100,000 random galaxies are then put down in the unmasked regions and the number of randoms within the aperture are counted (N$^{Aper}_{r}$). The number of galaxies within the aperture is then normalised to give the final density measurement, $\\rho/\\rho_r$: \\begin{equation} \\frac{\\rho}{\\rho_r}=\\frac{N^{Aper}_{g}}{N^{Aper}_{r}} \\times \\frac{N^{Tot}_{r}}{N^{Tot}_{g}} \\end{equation} where N$^{Tot}_{r}$ and N$^{Tot}_{g}$ are then total number of random points and galaxies respectively, so that $\\rho/\\rho_r$=1 corresponds to a density consistent with that of a random distribution of galaxies. This method was chosen to be the basis of this work as it is conceptually simple, and as concluded by \\citet{Cooper05}, this technique has a distinct advantage in fields masked by a large number of holes. The nearest neighbour method would require the exclusion of a large fraction of data close to holes and field edges. The $n$th nearest neighbour method was first employed by \\citet{Dressler}, this calculates the distance to the $n$th nearest galaxy, $D_{n}$ in Mpc and is expressed here as a surface density, \\begin{equation} \\Sigma_n = \\frac{n}{\\pi D_n^2} \\end{equation} The surface density, $\\Sigma_n$ is then renormalised such that, \\begin{equation} \\delta_n = {\\frac{\\Sigma_n}{\\bar{\\Sigma}}} \\end{equation} where $\\bar{\\Sigma}$ is the median density of galaxies within the field. To reduce the effect of the edges, the distance to the nearest edge was calculated and if this was less than the distance to the third nearest neighbour then the object was removed from the sample. This method was only used in this work to test the primary findings of the aperture method. The results are presented in the appendix. \\section[]{Results} Below we explore the relationship between galaxy colours and environment as a function of redshift. Relatively broad redshift bins are used to minimise the contamination due to photometric redshift errors. These sources of uncertainty are explored further in section 5. \\begin{figure} \\begin{center} \\includegraphics[angle=0, width=250pt]{Hist_All_1Mpc_smallerlastbin.eps} \\caption{Histograms of the density of galaxies within 1Mpc and 500kpc diameter apertures compared to a random sample for red (dashed line) and blue (thick line) galaxies. The $\\sigma$ values are obtained by performing a KS test, representing the significance in rejecting the null-hypothesis that the samples are drawn from the same underlying population.} \\label{RBHist} \\end{center} \\end{figure} \\begin{figure*} \\begin{minipage}{150mm} \\begin{center} \\includegraphics[angle=0, width=400pt]{MKvdensity_Aper_1Mpc_500kpc_250kpc_newlastbin.eps} \\caption{The average galaxy overdensity as a function of K-band luminosity, displayed in four redshift bins (top to bottom) and using projected apertures of diameter 1Mpc, 500kpc and 250kpc (three columns). Galaxies are displayed in red, blue, passive and star-forming subsets, as defined in Section 2. Note the change in scale for each column. The environments are defined so that $\\rho/\\rho_r$=1 corresponds to a density consistent with that of a random distribution of galaxies.} \\label{RBPass} \\end{center} \\end{minipage} \\end{figure*} \\begin{table*} \\begin{minipage}{150mm} \\begin{center} \\begin{tabular}{|c|c|c|c|c|} \\hline & $0.251.75$. Figure \\ref{RBPass} is a plot of the mean density of red, blue, passive (black) and star forming (cyan) galaxies in bins of absolute K-band magnitude. Error bars are derived from the error on the mean density of galaxies within a given bin. Sources of error are explored further in Section 5. The passive and star forming galaxies are defined in Section 2. As before they are plotted in four redshift bins but with an additional 250kpc aperture. This plot illustrates that red and/or passive galaxies reside in significantly denser environments than blue and/or star-forming galaxies from the present day to $z\\sim 1.5$, and this difference is apparent at all luminosities. This is comparable to what has been found in the local universe by other studies (\\citealt{Kauffmann}; \\citealt{van Der Wel}). Figure \\ref{RBPass} also shows that passive galaxies (shown in black) follow a similar density profile to red galaxies but are on average in slightly denser environments. This supports the conclusion that passive galaxies within the red population are responsible for the enhanced environments compared to blue star-forming objects. The environments of galaxies that were red but not in the strict `passive' sample were also investigated and these were found to lie in-between the red and blue galaxy environments, as would be expected (these are not shown for clarity). The actively star-forming galaxies exhibit the same environmental dependence as the blue galaxies, following the same luminosity-density profile. \\par In addition to the clear separation of red and blue galaxies, we also find a general trend of increasing galaxy density with luminosity, particularly for blue galaxies and on smaller scales. Inspecting the two intermediate redshift bins in Figure \\ref{RBPass}, on scales below $500$kpc we find that the most luminous blue galaxies appear to inhabit environments approaching those of red/passive galaxies. These results are consistent with the findings of \\cite{Cooper07}, who observed a strong increase in local density with luminosity for blue galaxies at $z\\sim 1$. Our results appear to extend these findings to higher redshift, suggesting that the epoch $11.5$ by an order of magnitude, which we expect to dramatically improve the reliability of photometric redshifts and allow us to extend our study of galaxy environments and large-scale structure to the crucial epoch when the galaxy red sequence is first established." }, "1101/1101.0414_arXiv.txt": { "abstract": "We investigate the long-term variability exhibited by the X-ray point sources in the starburst galaxy M82. By combining 9 \\chandra\\ observations taken between 1999 and 2007, we detect 58 X-ray point sources within the $D_{25}$ isophote of M82 down to a luminosity of $\\sim 10^{37}$\\lum. Of these 58 sources, we identify 3 supernova remnant candidates and one supersoft source. Twenty-six sources in M82 exhibit long-term (i.e., days to years) flux variability and 3 show long-term spectral variability. Furthermore, we classify 26 sources as variables and 10 as persistent sources. Among the total 26 variables, 17 varied by a flux ratio of $> 3$ and 6 are transient candidates. By comparing with other nearby galaxies, M82 shows extremely strong long-term X-ray variability that 47\\% of the X-ray sources are variables with a flux ratio of $> 3$. The strong X-ray variability of M82 suggests that the population is dominated by X-ray binaries. ", "introduction": " ", "conclusions": "" }, "1101/1101.5007_arXiv.txt": { "abstract": "We investigate the clustering of HI-selected galaxies in the ALFALFA survey and compare results with those obtained for HIPASS. Measurements of the angular correlation function and the inferred 3D-clustering are compared with results from direct spatial-correlation measurements. We are able to measure clustering on smaller angular scales and for galaxies with lower HI masses than was previously possible. We calculate the expected clustering of dark matter using the redshift distributions of HIPASS and ALFALFA and show that the ALFALFA sample is somewhat more anti-biased with respect to dark matter than the HIPASS sample. ", "introduction": "Measurements of the clustering of galaxies allows one to investigate the relationship between dark and luminous matter. By comparing galaxies selected in different ways one gains understanding of how different galaxies trace the underlying dark matter and also of processes at work in galaxy evolution. This information is important when using galaxies as probes of cosmological parameters. A number of new radio telescopes, such as the MeerKAT\\footnote{www.ska.ac.za} , ASKAP\\footnote{www.atnf.csiro.au/SKA/} and the SKA, are in the pipeline and they will detect huge numbers of galaxies using HI. A reliable measure of the bias of HI-selected galaxies and insight into the evolution of the bias is important for forecasting the capabilities of telescopes which will probe HI at intermediate or high-redshifts. The clustering of HI-selected galaxies has been studied by \\cite{Meyer-07}, \\cite{Basilakos-07} and \\cite{Ryan-Weber-06}. They used data from the HI Parkes All Sky Survey (HIPASS, \\citealt{Meyer-04}), a blind survey for HI of the southern sky which generated a catalogue of 4315 sources, the bulk of which have redshifts below $z\\sim0.02$. They showed that HI-selected galaxies are less clustered than galaxies selected in other ways. \\cite{Meyer-07} investigated clustering of various subsamples of HIPASS galaxies, showing that galaxies with high rotation velocities are more clustered than those with lower rotation velocities. There were indications that galaxies containing more HI are also more clustered but the differences were not as pronounced as in \\cite{Basilakos-07}. The latter work also measures the bias of HIPASS galaxies relative to the expected dark matter distribution. In this paper we measure the clustering of HI-selected galaxies detected with the Arecibo L-band Feed Array (ALFA) and compiled in the partially completed ALFALFA survey (the Arecibo Legacy Fast ALFA survey, \\citealt{Giovanelli-05a}). The results are compared with those obtained for HIPASS. Clustering measurements in HIPASS are limited to large angular scales where the beam-size of $\\sim 15$ arcmins does not cause confusion. The ALFALFA resolution is more than four times better allowing us to probe clustering on smaller scales. The rms noise per ALFALFA beam is about six times smaller, providing a catalogue of sources which spans a wider range of redshifts and includes galaxies with lower HI masses. We are thus able to measure clustering of HI-selected galaxies in regimes that have not yet been explored and to investigate trends seen in HIPASS, using an independent survey. The outline of the paper is as follows: In \\S~\\ref{sec:Data} we give a short introduction to the HIPASS and ALFALFA surveys. The computation of the angular and spatial two-point correlation functions is described in \\S~\\ref{sec:TwoPoint}. The results are presented, discussed and compared with earlier work in \\S~\\ref{sec:Results}. Finally, \\S~\\ref{sec:Conclusions} concludes with a short summary. ", "conclusions": "\\label{sec:Conclusions} We have measured the clustering of HI-selected galaxies using the ALFALFA survey data and compared this with results for HIPASS. Our two methods for determining the real-space correlation function agree well and our results for HIPASS agree with those found by \\cite{Meyer-04}. The real-space clustering in ALFALFA appears to be even lower than in HIPASS, consistent with the idea that ALFALFA probes galaxies with lower HI-masses that are less clustered than their high-mass counterparts. Our measurements of high- and low-mass subsamples in ALFALFA do not provide evidence to support this idea but the uncertainties on the measurements are large. We have calculated the clustering of dark matter expected within a $\\Lambda$CDM model with redshift distributions of HIPASS and ALFALFA. We then calculated the bias of ALFALFA sources over the range $1-10^{\\circ}$, finding a value of 0.62 at $1^{\\circ}$ and an average value of 0.52 over the whole range. The significant anti-bias of galaxies with low HI-mass is important to consider when estimating the signal-to-noise of experiments planned for the SKA and its pathfinders." }, "1101/1101.3883_arXiv.txt": { "abstract": "We report on the Iranian National Observatory (INO) ongoing site characterization studies for INO 3.4m optical telescope under development. Iran benefits from high altitude mountains and a relatively dry climate, thus offer many suitable sites for optical observations. The site selection (2001-2007) studies resulted in two promising sites in central Iran, one of which will host the 3.4m telescope. The studies between 2008 and 2010 aimed at detail characterization of the two sites. This involved measurements of a number of parameters including the wind speed and wind direction, astronomical seeing, sky brightness and microthermal variations. ", "introduction": "\\label{sec:intro} % The present research and training capabilities in observational astronomy in Iran can, by no mean, respond to the growing demand due to the rapid growth in higher education over the past decade. The existing observational facilities consists of a number of small telescopes in various university campus observatories generally used for undergraduate and graduate training. A medium size optical telescope is thought to be a step to facilitate research in astronomy and observational cosmology. The geographic location of Iran, 32N 53E, relative dry climate and high altitude mountains, offer suitable locations for optical telescopes. Site selection study for a proposed 2-4 meter class telescope started few years before the INO project received administrative approval. The study led by S. Nasiri (report in preparation) began by collecting and analysis of weather data, seismic hazard data, accessibility and shinny day statistics over central dry regions of the country. A large number of sites were identified and inspected. When the number of potential sites, mostly scattered around the central desert, was reduced to a manageable number, long term seeing monitoring has also started and continued for two years on 4 different sites with altitudes between 2500m and 3000m. ", "conclusions": "Our studies indicate a relative advantage of the Gargash site in comparison to Dinava site. Gargash site is found to be darker, benefitting from a better astronomical seeing and also higher altitude and therefore less affected by dust." }, "1101/1101.3284_arXiv.txt": { "abstract": "{The origin of the high-energy emission of blazars is still a matter of debate. To investigate the emission mechanism of extragalactic outflows and to pin down the location of the emission, we have constructed a broadband spectral energy distribution (SED) database covering from the radio to the gamma-ray band for the complete MOJAVE sample, which consists of 135 relativistically beamed AGN with well-studied parsec-scale jets. Typically, the broadband SEDs of blazars shows a double-humped profile. It is believed that the lower-energy hump is due to synchrotron emission from the radio jet, and the higher-energy hump is generated by i) inverse-Compton upscattered seed photons (leptonic), ii) proton-induced shower (hadronic). Combining the results of high-resolution VLBI observations and the $\\gamma$-ray properties of the MOJAVE sources, we attempt to reveal the origin of the high-energy emission in relativistic jets, and search for correlations between VLBI and high-energy properties. } \\FullConference{10th European VLBI Network Symposium and EVN Users Meeting: VLBI and the new generation of radio arrays\\\\ September 20-24, 2010\\\\ Manchester UK} \\begin{document} ", "introduction": "Active Galactic Nuclei (AGN) are among the most energetic objects in the Universe. AGN dominate the extragalactic high-energy sky, and are very active at all wavelengths from radio to $\\gamma$-rays. According to the unified model of AGN, it is believed that a super massive black hole is located in the center of host galaxy, and it fuels the whole system with matter accreted around the central engine. In the radio-loud scheme, an energetic jet is launched in the vicinity of the central engine following the warped magnetic field from the pole direction of the accretion disk. Blazars\\footnote{Usually, we use the term blazar to refer to BL Lac objects and flat spectrum radio quasars (FSRQ).} are AGN whose jets are pointing toward us, and dominate the radio and the high-energy sky. After the beginning of operations of the Large Area Telescope (LAT) on-board the \\textit{Fermi} $\\gamma$-ray Space telescope in mid 2008, \\textit{Fermi}/LAT detected 709 AGN in the first 11 months, and 85\\% of them being blazars \\cite{abdo10a}. The high-energy emission location in blazar systems is not yet well-understood. By using the very long baseline interferometry (VLBI) technique, we are able to resolve jet structure and trace component ejection at milli-arcsecond scales \\cite{lister09}. Combining the VLBI with the high-energy observatories (e.g., \\textit{Fermi}/LAT), we might be able to probe the location of the high-energy emission. The \\textbf{M}onitoring \\textbf{O}f \\textbf{J}ets in \\textbf{A}ctive Galactic Nuclei with \\textbf{V}LBA \\textbf{E}xperiments (MOJAVE) program has been monitoring a radio-selected sample since the mid 1990s. The sample contains mostly blazars due to the selection criteria\\footnote{(1) J2000.0 declination$\\leq-$20$^{\\circ}$; (2) galactic latitude |b|$\\leq$2.5$^{\\circ}$; (3) 15\\,GHz VLBI flux density$\\geq$1.5\\,Jy.} used \\cite{lister09a}. In the \\textit{Fermi} one-year AGN catalog \\cite{abdo10b}, 63\\% of the MOJAVE sources were detected. It was found that the LAT-detected MOJAVE sources have higher brightness temperature and higher Doppler-boosting factors than the non-detected ones \\cite{kovalev09}, and the authors suggested that the parsec-scale radio core is likely to be the location of the radio and the $\\gamma$-ray flares. It was also reported that the $\\gamma$-ray bright quasars have faster jets \\cite{lister09b}. In order to investigate the relation between the parsec-scale jets and high-energy emission, we are studying the broadband spectral energy distribution (SED) from the radio to the $\\gamma$-ray of the MOJAVE sample. By comparing the SED properties of the statistical-complete MOJAVE sources with the VLBI parameters, we want to further understand the physical mechanisms ongoing in blazar jets. ", "conclusions": "" }, "1101/1101.1762_arXiv.txt": { "abstract": "The high-resolution ({{\\it{R}}$\\sim$600) {\\it{Spitzer}}/IRS spectrum of the bipolar proto planetary nebula (PN) IRAS 17423$-$1755 is presented in order to clarify the dominant chemistry (C-rich versus O-rich) of its circumstellar envelope as well as to constrain its evolutionary stage. The high quality {\\it{Spitzer}}/IRS spectrum shows weak 9.7 $\\mu$m absorption from amorphous silicates. This confirms for the first time the O-rich nature of IRAS 17423$-$1755 in contradiction to a previous C-rich classification, which was based on the wrong identification of the strong 3.1 $\\mu$m absorption feature seen in the {\\it{Infrared Space Observatory}} (\\it{ISO}}) spectrum as due to acetylene (C$_{2}$H$_{2}$). The high-resolution {\\it{Spitzer}}/IRS spectrum displays a complete lack of C-rich mid-IR features such as molecular absorption features (e.g., 13.7 $\\mu$m C$_{2}$H$_{2}$, 14.0 $\\mu$m HCN, etc.) or the classical polycyclic aromatic hydrocarbon infrared emission bands. Thus, the strong 3.1 $\\mu$m absorption band toward IRAS 17423$-$1755 has to be identified as water ice. In addition, an [NeII] nebular emission line at 12.8 $\\mu$m is clearly detected, indicating that the ionization of its central region may be already started. The spectral energy distribution in the infrared ($\\sim$2$-$200 $\\mu$m) and other observational properties of IRAS 17423$-$1755 are discussed in comparison with the similar post-asymptotic giant branch (AGB) objects IRAS 19343$+$2926 and IRAS 17393$-$2727. We conclude that IRAS 17423$-$1755 is an O-rich high-mass post-AGB object that represents a link between OH/IR stars with extreme outflows and highly bipolar PN. ", "introduction": "IRAS 17423$-$1755 (Hen 3$-$1475) was first suggested by \\citet{ParthaPottasch89} as a possible member of the transition phase from the asymptotic giant branch (AGB) to the planetary nebula (PN) stage due to its unusual IRAS colors. The high values of the [NII]/H$\\alpha$ ratios in the outflowing material detected by \\citet{Riera95} and the low luminosity deduced for the central star allowed them to confirm the classification of this object as an evolved star. {\\it{Hubble Space Telescope}} ({\\it{HST}}) and {\\it{Very Large Array}} ({\\it{VLA}}) observations by \\citet{Bobrowski95} showed the presence of both OH maser emission and highly collimated ionized outflows like those detected in OH/IR stars or very young PN. {\\it{HST}} images revealed a rich and complex morphological structure in the circumstellar material (see Section 4 for details). The outflow is collimated in bipolar jets along several condensations of shock-excited gas that extend about 11 arcsec. The lobes show expansion velocities of about 425 km s$^{-1}$ and a high velocity jet ($\\sim$900 km s$^{-1}$) in the inner part of the lobes \\citep{Riera95}. The nebula displays a remarkable point symmetry that has been interpreted as due to the precession of a central binary system that undergoes episodic events of mass loss. \\citet{Bobrowski95} proposed that the expanding shell has a torus-like structure where the OH emission originates in a high density region, where H$_{2}$O is dissociated and further collimated in the observed jets. Additionally, \\citet{SanchezContreras01} found evidences of ultrafast winds (up to 2300 km s$^{-1}$) highly collimated and located close to the central star which could be a relatively young post-AGB outflow not strongly altered by interaction with the AGB. Stars at the end of the AGB phase are characterized by severe mass loss ($10^{-8}$ to $10^{-4}$ M$_{\\odot}$ yr$^{-1}$), which results in the formation of circumstellar envelopes \\citep{Herwig05}. The spherical symmetry of the envelopes of AGB stars is translated into a variety of shapes in the PN phase by a mechanism or mechanisms not as yet well understood. There is increasing evidence that at least in some instances the shaping starts at the end of the AGB phase \\citep{vanWinckel03}. IRAS 17423$-$1755 is a spectacular example that may represent a link between OH/IR stars with extreme outflows and highly bipolar PN. The spectral energy distribution (SED) of extreme (e.g., highly embedded) OH/IR AGB stars is characterized by the presence of strong and broad amorphous silicate absorption features at 9.7 and 18 $\\mu$m together with crystalline silicate absorption/emission features from 10 to 45 $\\mu$m \\citep{Sylvester99,GarciaH07}. At the end of the AGB phase, the crystalline silicate features become dominant and can be observed in more evolved O-rich PN \\citep{Molster01}. Comparison of the {\\it{Infrared Space Observatory}} ({\\it{ISO}}) observations of O-rich dust shells surrounding evolved stars with laboratory data suggested the presence of several families of crystalline silicates, such as olivines and pyroxenes, and marked the beginning of an emerging discipline: the mineralogy of stellar and other astronomical (i.e., cometary) dust shells. Water ice features at 3.1, 43, and 62 $\\mu$m have been additionally observed in heavily obscured and extremely bipolar sources such as the post-AGB star IRAS 19343$+$2926 (or M1$-$92, see e.g. Dijkstra et al. 2006 and references therein). \\citet{Gauba04} studied the {\\it{ISO}} spectra of seven hot post-AGB stars including IRAS 17423$-$1755. DUSTY models \\citep{Ivezic97} were fitted to optical, near- and far-infrared (IRAS and {\\it{ISO}}) photometry in order to reconstruct the SEDs and to derive physical parameters such as dust temperatures, mass loss rates, angular radii and the inner boundary of the dust envelopes. For the particular case of IRAS 17423$-$1755 they considered a combination of silicates and carbon in the circumstellar environment. They reported the presence of a broad absorption feature at 3.1 $\\mu$m that they identified as due to the presence of C$_{2}$H$_{2}$ and/or HCN in the circumstellar envelope. This identification led these authors to infer a C-rich chemistry for the shell. More recently, \\citet{Cerrigone09} presented observations of a sample of 26 hot post-AGB stars with the Infrared Array Camera and the Infrared Spectrograph (IRS) on board the {\\it{Spitzer Space Telescope}}. These observations were analyzed together with Two Micron All Sky Survey, IRAS and radio centimeter data in order to model the SEDs of the targets. \\citet{Cerrigone09} classified IRAS 17423$-$1755 as a C-rich star on the basis of the \\citet{Gauba04} report of the C$_{2}$H$_{2}$ feature at 3.1 $\\mu$m and in the absence of a strong 9.7 $\\mu$m amorphous silicate absorption/emission feature in their low-resolution (R$\\sim$64$-$128) {\\it{Spitzer}} spectrum. However, they pointed out that the expected polycyclic aromatic hydrocarbon (PAH) features in the 5--12 $\\mu$m region are not detected. It is to be noted here that weak and narrow molecular absorptions from C-based molecules such as 13.7 $\\mu$m C$_{2}$H$_{2}$, 14.0 $\\mu$m HCN, etc., are difficult to detect at the low resolution of their {\\it{Spitzer}} spectrum. The detection of these C-rich molecular absorptions - typical of C-rich AGB/post-AGB stars - requires in most of the cases higher resolution observations such as those provided by the high-resolution modes of {\\it{Spitzer}} ({\\it{R}}$\\sim$600) and {\\it{ISO}} ({\\it{R}}$\\sim$1000) (see e.g., \\citet{Cernicharo99,Cernicharo01, GarciaH09}). The controversial origin of the mid- to far-IR features in IRAS 17423$-$1755 merits a re-analysis of the dust features observed in the high-resolution and higher quality {\\it{Spitzer}} spectrum. In Section 2 we present the new {\\it{Spitzer}} observations together with the construction of the overall SED of the nebula as observed by both {\\it{Spitzer}} and {\\it{ISO}}, while in Section 3 the evidence for an O-rich chemistry is analyzed and discussed. The evolutionary stage of IRAS 17423$-$1755 is discussed, including our new results, in Section 4 while a summary of our main conclusions is presented in Section 5. ", "conclusions": "An O-rich chemistry for the circumstellar envelope around the post-AGB object IRAS 17423$-$1755 is here confirmed, despite a previous classification as C-rich. This result is based on the detection of a weak and broad 9.7 $\\mu$m amorphous silicate absorption in the high-resolution {\\it{Spitzer}}/IRS spectrum. The complete lack of C-rich mid-IR features (in particular C$_{2}$H$_{2}$ at 13.7 $\\mu$m) supports our identification of the strong 3.1 $\\mu$m absorption band seen in the {\\it{ISO}} spectrum as due to water ice as well as our O-rich classification for IRAS 17423$-$1755. IRAS 17423$-$1755, IRAS 19343$+$2926 and IRAS 17393$-$2727 present clear evidences of the presence of a circumstellar disk or torus, where the conditions would be very similar to those found in less evolved and more embedded OH/IR stars. A recent strong mass-loss event has been reported in the case of IRAS 19343$+$2926 \\citep{Alcolea07}, which would favor this scenario. Water ice and crystalline silicates would preferentially form in the outer region of the inner torus, where low temperature conditions and shielding from the central star would allow a favorable rate of crystallization to take place. Both in IRAS 17423$-$1755 and IRAS 19343$+$2926, the ice band at 3.1 $\\mu$m is sharp and presents substructures, and comparison with models of water ice growth around evolved stars \\citep{Smith89, Dijkstra06} allows us to confirm that the ice is mostly in a crystalline state. The morphological properties, detection of OH maser emission, and the {\\it{Spitzer}}/IRS spectra observed in IRAS 17423$-$1755 are similar to those of the O-rich post-AGB stars IRAS 19343$+$2926 and IRAS 17393$-$2727, allowing us to interpret the evolutionary stage of IRAS 17423$-$1755 as belonging to an intermediate stage between those OH/IR stars with extreme outflows and highly bipolar type I PN." }, "1101/1101.0691_arXiv.txt": { "abstract": "Planetary magnetic fields could impact the evolution of planetary atmospheres and have a role in the determination of the required conditions for the emergence and evolution of life (planetary habitability). We study here the role of rotation in the evolution of dynamo-generated magnetic fields in massive Earth-like planets, Super Earths (1-10 $\\ME$). Using the most recent thermal evolution models of Super Earths \\citep{Gaidos10, Tachinami10} and updated scaling laws for convection-driven dynamos, we predict the evolution of the local Rossby number. This quantity is one of the proxies for core magnetic field regime, i.e. non-reversing dipolar, reversing dipolar and multipolar. We study the dependence of the local Rossby number and hence the core magnetic field regime on planetary mass and rotation rate. Previous works have focused only on the evolution of core magnetic fields assuming rapidly rotating planets, i.e. planets in the dipolar regime. In this work we go further, including the effects of rotation in the evolution of planetary magnetic field regime and obtaining global constraints to the existence of intense protective magnetic fields in rapidly and slowly rotating Super Earths. We find that the emergence and continued existence of a protective planetary magnetic field is not only a function of planetary mass but also depend on rotation rate. Low-mass Super Earths ($M\\lesssim 2\\ME$) develop intense surface magnetic fields but their lifetimes will be limited to 2-4 Gyrs for rotational periods larger than 1-4 days. On the other hand and also in the case of slowly rotating planets, more massive Super Earths ($M\\gtrsim 2\\ME$) have weak magnetic fields but their dipoles will last longer. Finally we analyze tidally locked Super Earths inside and outside the habitable zone of GKM stars. Using the results obtained here we develop a classification of Super Earths based on the rotation rate and according to the evolving properties of dynamo-generated planetary magnetic fields. ", "introduction": "\\label{sec:introduction} The number of known exoplanets in the mass range between 1 and 10 $\\ME$ is growing (hereafter these objects will be called ``Super Earths'' or SEs following the classification by \\cite{Valencia06,Valencia07a}). At the time of writing, there are almost 46 confirmed planets in this mass range\\footnote{For updates, please refer to {\\tt http://exoplanet.eu}}\\citep{Rivera05, Beaulieu06, Udry07, Mayor08, Ribas08, Queloz09, Bonfils11, Lissauer11} and more than a few hundred SEs candidates are awaiting further analysis and confirmation \\citep{Borucki11}. These discoveries have increased the interest to model and understand the geophysical properties of this type of planets \\citep{Valencia06, Valencia07a, Valencia07b,Valencia09,Seager07,Kaltenegger10a,Korenaga10}. The habitability of SEs, in particular those similar in composition and structure to the Earth, is an interesting topic in the field and several theoretical works have paid special attention to this particular aspect of SEs properties \\citep{Griebmeier05, Griebmeier09, Griebmeier10, Selsis07a, VanThienen07, VonBloh07, Lammer10}. Models of the interior structure of SEs have been extensively developed over the last 5 years \\citep{Valencia06, Valencia07a, Valencia07b, Fortney07, Seager07, Selsis07a, Sotin07, Adams08, Baraffe08, Grasset09}. Although there are still open issues to be addressed, these models are giving us an understanding of global properties such as the mass-radius relationship and its dependence with planetary composition, as well as different geophysical phenomena such as mantle convection, degassing and plate tectonics \\citep{Olson07,Papuc08,Valencia07c,Valencia09,Korenaga10}. Recently several authors have studied in detail the thermal evolution and magnetic field properties of this type of planets \\citep{Gaidos10,Tachinami10, Driscoll11}. Planetary magnetic fields would likely play a role in planetary habitability \\citep{VonBloh07, Griebmeier05, Griebmeier09, Griebmeier10, VanThienen07, Lammer10}. Understanding the conditions for the emergence and long term evolution of a protective planetary magnetic field (hereafter PMF) is crucial to evaluate the complex conditions for habitability of SEs. The same conditions could also be applied to evaluate the habitability of exomoons around extrasolar giant planets \\citep{Kaltenegger10b}. The current understanding of PMF emergence and evolution in SEs arises from thermal evolution models for the Earth \\citep{Stevenson03, Labrosse03, Labrosse07a, Labrosse07b, Nimmo09a, Aubert09, Breuer10} and scaling laws for convection-driven dynamos obtained from extensive numerical simulations \\citep{Christensen06, Olson06, Aubert09, Christensen09, Christensen10}. Two recent works studied the problem of PMF evolution in SEs by developing detailed models of planetary thermal evolution \\citep{Gaidos10,Tachinami10}. Both works have paid special attention to different but complementary aspects of the problem. \\citet{Gaidos10} uses a model of the structure of the planetary core and its thermal evolution (hereafter the {\\it Core Thermal Evolution} or {\\it \\CTE model}). On the other hand \\citet{Tachinami10} uses the Mixing Length Theory adapted to planetary conditions to model mantle convection with a detailed treatment of its rheological properties (hereafter the {\\it Mantle based Thermal Evolution model} or {\\it \\MTE model}). Although thermal evolution models for the Earth, other terrestrial planets and even SEs have been developed in the past \\citep{Stevenson03, Labrosse03, Labrosse07a, Papuc08, Nimmo09a, Breuer10}, the \\CTE and \\MTE models give the first detailed description aimed at studying dynamo-generated magnetic fields of extrasolar terrestrial planets. We use the results of the \\CTE and \\MTE models to study the role of rotation in the evolution of PMF in SEs. We focus on the evolution of the regime of the core magnetic field (CMF) which can be broadly classified as non-reversing dipolar, reversing dipolar and multipolar. For this purpose we compute the {\\it local Rossby number}, one of the proxies for CMF regime, as a function of the rotation period and planetary mass. Using this property we predict the long term evolution of the surface PMF in rapidly and slowly rotating planets. In section \\ref{sec:thermal-models} we summarize the most important results of the \\CTE and \\MTE models. Section \\ref{sec:dynamo-props} presents the scaling laws for convection-driven dynamos used to predict the properties of the CMF. In section \\ref{sec:rotation-PMF} we present a general procedure to compute the CMF intensity in the dipolar and multipolar regimes including an implicit dependence on rotation rate. In section \\ref{sec:results} we present the results of applying the procedure devised here to predict the maximum dipolar component of the field in SEs with different periods of rotation and thermal histories as predicted by the \\CTE and \\MTE models. Section \\ref{sec:discussion} is devoted to discuss the limitations of our procedure and the implications of our results. A summary, concluding remarks and future prospects are presented in section \\ref{sec:conclusions}. For reference a list of the symbols and the physical quantities used in this work are presented in Table \\ref{tab:symbols}. ", "conclusions": "\\label{sec:conclusions} We studied the role of rotation in the evolution of dynamo-generated magnetic fields in Super Earths. We computed the evolution of the local Rossby number and the volumetric magnetic field strength for core dynamos in SEs. For this purpose we used the results of two recently published thermal evolution models and scaling laws fitted with numerical dynamo experiments. Assuming that the local Rossby number could be used as a proxy to dynamo regime, we estimated the maximum dipolar component of the magnetic field at the CMB, and from it an upper bound to the dipolar part of the field at the planetary surface. We used two properties to characterize the global magnetic properties of SEs: (1) the average of the surface dipolar component of the field, $B_{avg}$ and (2) the total time $T_{dip}$ spent by the dynamo in the dipolar dominated regime (reversing and non reversing). Intense magnetic fields with a strong dipolar component \\citep{Stadelmann10}, are best suited to protect planetary environments from external agents (stellar wind and cosmic rays). Therefore large values of $B_{avg}$, irrespective of the dynamo regime, are consistent with planetary habitability. The long-term preservation of water and other volatiles in a planetary atmosphere and the development of life, would require long-lived protective PMFs, i.e. large values of $T_{dip}$. Intense and protective magnetic fields in the early phases of planetary and stellar evolution will be also suited for the preservation of an atmosphere or its volatiles. However a planet that achieves to preserve its atmosphere during the harsh conditions of the early active phases of stellar evolution but lacks of a protective magnetic field soon after this period will leave emerging forms of life to an integrated effect of galactic and stellar cosmic rays induced damages. We found that the PMF properties depend strongly on planetary mass and rotation period. Rapidly rotating SEs $P\\approx O(1)$ day, with mass $\\approx O(1) \\ME$, have the best potential to develop long-lived and intense PMFs. More massive planets develop weaker magnetic fields but they have dipolar dominated dynamos in a slightly larger range of rotational periods $P\\sim 1-3$ days. SEs with rotation periods larger than $3-10$ days (depending on their mass) will spend the majority of the dynamo lifetime in a multipolar state. In order to summarize our results we have introduced a rotation-based classification. Using the \\CTE thermal evolution model, SEs could be rapid rotators $P\\lesssim 1.5-4$ days, slow rotators $4\\lesssim P\\lesssim 10-20$ days and otherwise, very slow rotators. Planets in the HZ of low mass stars $\\Ms<0.6$ that will be tidally locked in less than 1 Gyr, will fall between the slow and very slow rotator types. Unlocked planets could be any of the types described before, according to their primordial period of rotation and the effects that could dampen it. More theoretical and observational efforts should be undertaken to address the problem of direct or indirect detection of PMF around low mass planets. The detection and measurement of such planetary magnetic fields will help us to constrain thermal evolution and dynamo models. The role of magnetic fields in planetary habitability is another problem that deserves close attention. Recent works have tackled this problem in detail \\citep{Griebmeier10} but their PMF models are too simplistic. Although the effect of rotation rates is considered in those models and they have focused on tidally locked planets, their models do not include the effects of thermal evolution on the PMF properties and their treatment of the dependence of these properties of the rotation rate is also limited." }, "1101/1101.5761_arXiv.txt": { "abstract": "We investigate the effects of a coupled Dark Energy (cDE) scalar field on the alignment between satellites and matter distributions in galaxy clusters. Using high-resolution N-body simulations for $\\Lambda$CDM and cDE cosmological models, we compute the probability density distribution for the alignment angle between the satellite galaxies and underlying matter distributions, finding a difference between the two scenarios. With respect to $\\Lambda$CDM, in cDE cosmologies the satellite galaxies are less preferentially located along the major axis of the matter distribution, possibly reducing the tension with observational data. A physical explanation is that the coupling between dark matter and dark energy acts as an additional tidal force on the satellite galaxies diminishing the alignments between their distribution and the matter one. Through a Wald test based on the generalized $\\chi^{2}$ statistics, the null hypothesis that the two probability distributions come from the same parent population is rejected at the $99\\%$ confidence level. It is concluded that the galaxy-matter alignment in clusters may provide a unique probe of dark sector interactions as well as the nature of dark energy. ", "introduction": "The concordance cosmological model -- based on a family of non-relativistic massive particles possibly weakly interacting with the standard sector of particle physics (Cold Dark Matter particles, CDM hereafter) and on a cosmological constant $\\Lambda $ -- very successfully accounts for a large wealth of observational data, but poses significant theoretical puzzles and requires an extreme fine-tuning of its basic parameters. In this context, alternative scenarios have been explored, as i.e. the possibility that the cosmological constant $\\Lambda $ be replaced by a Dark Energy (DE) component represented by a classical dynamical scalar field $\\phi $ evolving in a self interaction potential \\citep[][]{Wetterich_1988,Ratra_Peebles_1988}. A further step in the exploration of alternative cosmological models has been recently proposed, speculating about a possible direct interaction between the DE and the CDM sectors of the Universe \\citep[][]{Wetterich_1995,Amendola00,Amendola_2004}. Such coupled Dark Energy models (cDE) have been widely investigated in recent years concerning their background evolution and their effects on structure formation \\citep[e.g. by][]{Mainini_Bonometto_2006,Pettorino_Baccigalupi_2008,abdalla-etal09,Wintergerst_Pettorino_2010} by means of specifically designed N-body algorithms aimed at exploring the nonlinear evolution of cosmic structures within these alternative scenarios \\citep[][]{Maccio-etal04,baldi-etal10,Baldi_2010,Li_Barrow_2010}. These studies have led to the highlight of some distinctive features in the evolved matter density fields of cDE models and have shown how the properties of nonlinear collapsed objects can be significantly affected by the dark interactions in a potentially observable way \\citep[][]{BV10,Baldi_Pettorino_2010}. Therefore, one of the most important tasks in the present investigation of these alternative models consists in linking the predictions obtained for the CDM density distribution to directly observable quantities. \\citet{lee10} very recently brought up a speculative idea, namely that the dark sector interaction might be probed by measuring the alignment between galaxy and matter distributions in triaxial clusters. In $\\Lambda$CDM cosmology, the spatial correlations of the large-scale tidal fields yield the strong alignments of the satellite galaxies in the clusters with the cluster dark matter distributions \\citep[e.g.,][]{altay-etal06}. \\citet{lee10} claimed that the alignment of the cluster galaxies with the underlying CDM distribution would be weaker in the cDE models than in $\\Lambda$CDM since the spatial correlations of the tidal fields inside the clusters would be less strong due to the existence of a fifth force. If this is really the case, the alignment between galaxy and matter distributions in triaxial clusters would provide a unique probe of the cDE scenarios, especially because the alignment angle is insensitive to the other cosmological parameters. To back up the speculative idea of \\citet{lee10}, it is essential to examine numerically whether or not the cDE scalar field truly makes a detectable difference on the cluster galaxy-matter alignment. The goal of this Paper is to use high-resolution N-body simulations to study how the alignment between satellites and the underlying dark matter distribution in galaxy clusters changes in cDE models as compared to $\\Lambda$CDM \\citep{wmap7}. ", "conclusions": "Using data from high-resolution N-body simulations, we have shown that the $\\Lambda$CDM and cDE cosmologies yield different strengths of the alignment between satellite galaxy and matter distributions in triaxial clusters. As speculated by \\citet{lee10}, it is found that the alignment is less strong in cDE cosmologies than in the standard $\\Lambda$CDM scenario. The null hypothesis that the two cosmologies give the same alignment tendency is rejected at the $98.9\\%$ confidence level through a Wald test based on the generalized $\\chi^{2}$ statistics. The differences in the satellite distributions between $\\Lambda$CDM and cDE models may be understood as follows. There are two possible mechanisms for explaining the alignments between the dark matter and galaxy distributions \\citep{altay-etal06}. The first one is the spatial correlations between the large scale external tidal fields and the internal tidal fields due to the cluster potentials which affect he non-spherical shapes of dark matter components of galaxy clusters and the anisotropic distributions of their satellite galaxies, respectively. The second one is the filamentary merging/accretion of matter and galaxies along the large-scale filaments, which also induce the alignments between the dark matter and galaxy distributions in clusters. In the $\\Lambda$CDM cosmology, the tidal fields evolve nonlinearly only via gravity. Whereas in the cDE cosmology the coupling between dark energy and dark matter generate additional (more isotropic) tidal effects in the clusters, which weaken the spatial correlations between the external and internal tidal fields and also redistribute previously accreted satellites that were initially in a flattened distribution. Furthermore, the ``modified inertia\" characterizing cDE models would contribute to an effective enhancement of the internal cluster tidal field. Recently \\citet{oguri-etal10} compared the observed distribution of satellite galaxies in clusters with the shape of the cluster dark matter distribution determined from weak lensing, finding a very weak alignment \\citep{lee10}. Due to the small sample (composed of only $13$ clusters), it is difficult to say at the moment whether or not this observational signal supports the cDE scenario. In any case, a crucial implication of our results is that the alignment between galaxy and matter distributions in clusters is in principle a unique probe of dark sector interactions." }, "1101/1101.2377_arXiv.txt": { "abstract": "Accretion disk winds are revealed in {\\it Chandra} gratings spectra of black holes. The winds are hot and highly ionized (typically composed of He-like and H-like charge states), and show modest blue-shifts. Similar line spectra are sometimes seen in ``dipping'' low-mass X-ray binaries, which are likely viewed edge-on; however, that absorption is tied to structures in the outer disk, and blue-shifts are not typically observed. Here we report the detection of blue--shifted He-like Fe XXV ($3100\\pm 400$~km/s) and H-like Fe XXVI ($1000\\pm 200$~km/s) absorption lines in a {\\it Chandra}/HETG spectrum of the transient pulsar and low-mass X-ray binary IGR J17480$-$2446 in Terzan 5. These features indicate a disk wind with at least superficial similarities to those observed in stellar-mass black holes. The wind does not vary strongly with numerous weak X-ray bursts or flares. A broad Fe K emission line is detected in the spectrum, and fits with different line models suggest that the inner accretion disk in this system may be truncated. If the stellar magnetic field truncates the disk, a field strength of $B = $0.7--4.0$\\times 10^{9}~ {\\rm G}$ is implied, which is in line with estimates based on X-ray timing techniques. We discuss our findings in the context of accretion flows onto neutron stars and stellar-mass black holes. ", "introduction": "IGR J17480$-$2446 was discovered on 10 October 2010 in an INTEGRAL monitoring observation of the Galactic bulge (Bordas et al.\\ 2010). The source was found to be consistent with the position of the globular cluster Terzan 5, which hosts the better-known transient and Type-I X-ray burst source EXO 1745$-$248. An affiliation with EXO 1745$-$248 was apparently strengthened with the detection of Type-I bursts (Chevenez et al.\\ 2010; Strohmayer \\& Markwardt 2010). Analysis of a prior {\\it Chandra} image of Terzan 5, and additional {\\it Swift} observations revealed, however, that the source was not EXO 1745$-$248, but rather a new transient source (Heinke et al.\\ 2010; also see Pooley et al.\\ 2010). The source was subsequently given the name IGR J17480$-$2446, in recognition of its discovery with INTEGRAL (Ferrigno et al.\\ 2010). Detailed timing analysis shows that IGR J17480$-$2446 is an 11~Hz pulsar (Strohmayer \\& Markwardt 2010). It also appears that IGR J17480$-$2446 evolved from an ``atoll'' into a ``Z'' source, based on the path it traces in an X-ray color--color diagram and its rapid variability components (including the possible detection of a kHz quasi-periodic oscillation, or QPO, at 815~Hz; see Altamirano et al.\\ 2010). Previously, only XTE J1701$-$462 had been observed to evolve in this manner (Homan et al.\\ 2010). Slower, mHz QPOs were also detected from IGR J17480$-$2446, caused by recurrent weak bursts (Linares et al.\\ 2010). Neutron stars in globular clusters are of particular importance because the distance to globular clusters is often known precisely. This can eliminate a major source of uncertainty when estimating a blackbody emission radius. The distance to Terzan 5 is likely 5.5~kpc (Ortolani et al.\\ 2007; also see Ransom 2007). Motivated by the rare opportunity to study a bursting neutron star in a globular cluster, we proposed a {\\it Chandra} Director's Time observation of IGR J17480$-$2446. ", "conclusions": "We have analyzed a {\\it Chandra}/HETG observation of the transient X-ray pulsar and low-mass X-ray binary IGR J17480$-$2446 in Terzan 5. The Fe K band can be decomposed into a broad emission line consistent with He-like emission excited at inner edge of a radially-truncated accretion disk, and blue-shifted He-like and H-like Fe absorption lines that likely arise in an X-ray disk wind. The implications of a truncated inner accretion disk in IGR J17480$-$2446 and a disk wind are of particular interest, and they are explored in this section. At 11~Hz, IGR J17480$-$2446 is a much slower pulsar than e.g. SAX J1808.4$-$3658, which has a frequency of 401~Hz (Wijnands \\& van der Klis 1998), and in which relativistic line spectroscopy and X-ray timing both imply a smaller inner disk radius and smaller stellar magnetic field (Cackett et al.\\ 2009b; Papitto et al.\\ 2009, Hartman et al.\\ 2008). In SAX J1808.4$-$3658, then, it is not surprising that the source shows kHz QPOs, which are often associated with orbits in the inner disk, and connections between the disk and the star. At a flux level similar to that measured in this {\\it Chandra} observation (after accounting for a distance disparity), Altimirano et al.\\ (2010) reported the detection of QPOs at 48~Hz, 173~Hz, and 815~Hz. If the fastest QPO signals a Keplerian orbital frequency, a radius of approximately $9~ {\\rm GM}/{\\rm c}^{2}$ is implied. If the inner disk was still truncated at $20(2)~ {\\rm GM}/{\\rm c}^{2}$, then the kHz QPOs would have to be produced interior to the disk truncation radius. Fourier-resolved spectroscopy has shown that the variable part of such spectra may originates in the boundary layer (Revnivtsev \\& Gilfanov 2006); thus, it is possible that kHz QPOS originate in the boundary layer. A more likely explanation is that the disk simply extended closer to the neutron star when kHz QPOs were detected. The disk wind found in IGR J17480$-$2446 may be the clearest detection of such a flow in a neutron star system. Iron absorption lines from highly ionized gas are clearly detected in a number of ``dipping'' neutron star systems, but this absorption is not typically blue-shifted, and likely occurs in the outer accretion disk of these nearly edge-on sources (see, e.g., Diaz Trigo et al.\\ 2006). An ionized outflow is clearly detected in the persistent spectrum of GX 13$+$1 (Ueda et al.\\ 2004); however, this source shows frequent dips, and this casts some doubt on the nature of the absorption. Similarly, an ionized outflow is detected in Circinus X-1 (Brandt \\& Schulz 2000), but it is difficult to rule out a massive companion wind as the source of the outflow. An ionized X-ray wind was detected in a {\\it Chandra} spectrum of the pulsar 1A 0535$+$262 (Reynolds \\& Miller 2010). The flow is likely a disk wind, but it is difficult to entirely rule out absorption in the wind of the massive B companion star in 1A 0535$+$262. It is interesting to explore why disk winds are not regularly detected in the relatively large sample of persistent ``Z'' and ``atoll'' neutron stars. A recent survey of {\\it Chandra}/HETG spectra does not find evidence for ionized outflows (Cackett et al.\\ 2009); in some cases the sensitivity of the spectra may have simply been insufficient. However, even in deep observations of Cygnus X-2, blue-shifted absorption is not reported in the non-dip spectra (Schulz et al.\\ 2009). An intriguing possibility is that the outflow observed in IGR J17480$-$2446 may be a different kind of disk wind than those observed in stellar-mass black holes. The winds seen in stellar-mass black holes originate close to the black hole (e.g. within $1000~GM/c^{2}$; see Miller et al.\\ 2006a, 2006b, 2008; Kubota et al.\\ 2007; Neilsen \\& Lee 2009), and can carry away a good fraction of the accreting gas. Thus, the winds in stellar-mass black holes may be at least partially driven by magnetic pressure (see Miller et al.\\ 2008). In contrast, X-ray absorption in the disk wind found in IGR J17480$-$2446 originates two orders of magnitude further from the central engine (measured in units of $GM/c^{2}$). The high ionization parameter of the wind in IGR J17480$-$2446 makes it unlikely that radiation pressure can drive the wind (e.g. Proga et al.\\ 2000), but Compton heating of the outer disk could plausibly drive such a wind (e.g. Begelman, McKee, \\& Shields 1983). \\vspace{0.2in} We thank Harvey Tananbaum, Belinda Wilkes, and Andrea Prestwich for executing this observation. We acknowledge David Pooley and Jeroen Homan for helping to coordinate different observations of IGR J17480$-$2446. We thank Tim Kallman and Cole Miller for helpful discussions. Finally, we thank the anonymous referee for a helpful review. JMM ackhowledges support from the {\\it Chandra} Guest Observer program." }, "1101/1101.0144_arXiv.txt": { "abstract": "Recent observations of Gamma-Ray Bursts (GRBs) by the Fermi Large Area Telescope (LAT) revealed a power law decay feature of the high energy emission (above 100 MeV), which led to the suggestion that it originates from a (probably radiative) external shock. We analyze four GRBs (080916C, 090510, 090902B and 090926A) jointly detected by Fermi LAT and Gamma-ray Burst Monitor (GBM), which have high quality lightcurves in both instrument energy bands. Using the MeV prompt emission (GBM) data, we can record the energy output from the central engine as a function of time. Assuming a constant radiative efficiency, we are able to track energy accumulation in the external shock using our internal/external shell model code. By solving for the early evolution of both an adiabatic and a radiative blastwave, we calculate the high energy emission lightcurve in the LAT band and compare it with the observed one for each burst. The late time LAT light curves after $T_{90}$ can be well fit by the model. However, due to continuous energy injection into the blastwave during the prompt emission phase, the early external shock emission cannot account for the observed GeV flux level. The high energy emission during the prompt phase (before $T_{90}$) is most likely a superposition of a gradually enhancing external shock component and a dominant emission component that is of an internal origin. ", "introduction": "The Large Area Telescope (LAT; Atwood et al. 2009) aboard the Fermi Gamma Ray Space Telescope (Fermi) has recently detected nearly 20 GRBs (e.g. Abdo et al. 2009a,b; Abdo et al. 2010; Ackermann et al. 2010, see Zhang et al. 2011 for a synthetic study). Among them, several bright GRBs (e.g. GRBs 080916C, 090510, 090902B and 090926A) have well sampled long-term LAT-band lightcurves. In logarithmic space, these GRBs have count rates that rise, peak and begin decaying before the MeV prompt emission is over, i.e. peaking at a time smaller than $T_{90}$ defined in the Gamma-ray Burst Monitor (GBM; Meegan et al. 2009) detector energy band. The post peak lightcurve typically has a decay slope steeper than $-1$ (e.g. ranging from $-1.3$ to $-2$, Ghisellini et al. 2009; Zhang et al. 2011). The simple temporal behavior (a broken power law lightcurve) of LAT emission led to the suggestion that GRB GeV emission is of an external forward shock origin (Kumar \\& Barniol Duran 2009, 2010; Ghisellini et al. 2009), possibly from a highly radiative blastwave. A simple broken power law lightcurve is expected from the blastwave evolution of an instantaneously injected fireball with fixed explosion energy. Such an approximation is valid if the analyzed time scale is much longer than $T_{90}$, the duration of the prompt gamma-ray emission itself. However, for the early blastwave evolution, especially during the epoch when the central engine is still active (as is the case for the LAT GRBs discussed in this paper), one would not expect a simple lightcurve evolution, since the energy output from the central engine is continuously injected into the blastwave. The high quality spectral and temporal data of GRBs co-detected by Fermi LAT and GBM allow us to track the energy output from the central engine as a function of time. Recently we have developed a shell code to model the internal and external shock development for arbitrary central engine activities (Maxham \\& Zhang 2009). By processing the spectral and temporal evolution data of Fermi GRBs using the method described in Zhang et al. (2011), we can model the early development of the external shock based on first hand data. \\section[]{Data Analysis} We study four bright LAT GRBs (080916C, 090510, 090902B, and 090926A). GBM and LAT data reduction was carried out using the data analysis script introduced in Zhang et al. (2011). This code uses the public Fermi data and extracts time-resolved spectral information derived from a joint GBM/LAT fit. For the GBM data, the background spectrum is extracted using the CSPEC data, while the source spectrum is extracted using the event (TTE) data. The LAT background is different since only a few photons are detected by LAT for most GRBs, so on-source region data long after the GBM trigger when the photon counts merge into a Poisson noise are used to derive the LAT lightcurve background. The GBM and LAT data are then used to make dynamically time-dependent spectral fits. The code refines the number of time slices as necessary to preserve adequate statistics in each bin, and a spectral fit is chosen among a list of spectral models, such as a single power-law, a power-law with exponential cut-off, a Band function, a black body or a combination of these. Chi square statistics are performed to determine which fits are the best, and Ockham's Razor chooses the simplest spectral model between two statistically reasonable fits (Zhang et al. 2011). For the 4 bright GRBs in our sample, we adopt the following models (for details, see Zhang et al. 2011). For GRB 080916C and 090926A we adopt the Band function model throughout the burst, with the spectral parameters evolving with time. GRB 090902B shows a blackbody thermal component plus a non-thermal single power law component, and the short burst GRB 090510 is best-fit with a cutoff power law plus power law component. Similar to Ghisellini et al. (2009), we found that the long-term LAT light curves decay before the end of $T_{90}$ with a slope steeper than -1. \\section[]{External shock modeling} \\subsection[]{Blastwave evolution} We model a GRB as an explosion of many matter shells with some mass and Lorentz factor (Rees \\& M\\'esz\\'aros 1994). As the first matter shell moves outward into the ambient medium, it slows down when sweeping up this medium (M\\'esz\\'aros \\& Rees 1993). As time goes by, more and more trailing shells pile up onto the leading decelerating shell (Rees \\& M\\'esz\\'aros 1998). For an instantaneous explosion with constant energy, the motion of this decelerating ejecta along with the medium collected along the way, known as the ''blastwave\", is governed by three differential equations (Chiang \\& Dermer 1999; Huang et al. 2000): radius changing with time, $\\f{d R}{dt} = \\beta c = \\f{\\sqrt{\\gamma^2 - 1}}{\\gamma}c$, a statement of conservation of energy and momentum across the blastwave $\\f{d \\gamma}{dm} = \\f{-(\\gamma^2 -1)}{M}$ (Blandford \\& McKee 1976), and the amount of medium swept up as a function of radius $\\f{dm}{dR} = 4 \\pi R^2 \\rho$. Here $t$ is the time since explosion in the rest frame of the central engine, $R$ is the distance from the central engine, $\\rho$ is the density of the ambient medium, $\\gamma$ is the Lorentz factor of the shell, $m$ is the swept-up mass, and $M=M_0+\\gamma m$ is the total effective mass including the internal energy of the blastwave, where $M_0$ is the initial mass of the ejecta. As a result, one has another differential equation, \\begin{equation} \\f{dm}{dM} = \\f{1}{(1-\\epsilon)\\gamma + \\epsilon}~, \\label{Mandm} \\end{equation} where the value of $0\\leq \\epsilon \\leq 1$ determines the efficiency of the radiation, with 0 representing the purely adiabatic case and 1 representing the fully radiative condition. The above set of differential equations can be solved analytically. The adiabatic solution was presented as Eq.(14) in Maxham \\& Zhang (2009). Since the LAT lightcurves decay with a slope steeper than -1 (typical value for an adiabatic blastwave), e.g. in the range of -1.3 and -2 (Zhang et al. 2011), it may be reasonable to assume a completely radiative blastwave (Ghisellini et al. 2009). By adopting a value of $\\epsilon = 1$, one can get a purely radiative solution for the blastwave, which reads \\begin{equation} \\gamma = \\f{9 (M_0 \\gamma_0)^2 + 12 \\pi \\rho M_0 R^3 (1+\\gamma_0) + 8 \\pi^2 \\rho^2 R^6 (1+\\gamma_0) }{9 M_0^2 + 12 \\pi \\rho M_0 R^3 (1+\\gamma_0) + 8 \\pi^2 \\rho^2 R^6 (1+\\gamma_0)}. \\label{bwsolution} \\end{equation} In the deceleration regime, one has $\\gamma \\propto R^{-3}$, and $F_\\nu \\propto t^{(2-6p)/7}$ for $\\nu > {\\rm max}(\\nu_m, \\nu_c)$ (which is relevant for LAT band), which is $F_\\nu \\propto t^{-1.6}$ for $p=2.2$ (e.g. Sari et al. 1998). This is consistent with the rapid decay observations. \\subsection{Energy injection into the blastwave} During the prompt emission phase (i.e. $T < T_{90}$), the central engine continuously injects energy into the blastwave. So the solution should take into account the progressively increasing total energy in the blastwave. We apply the shell code developed and laid forth in Maxham \\& Zhang (2009) to this problem. The code, which originally generated randomized matter shells with different mass, Lorentz factor and ejection time, is here modified to use input values for these parameters which are taken from the data as follows. The most important parameter affecting blastwave evolution is the total injection energy. In principle the injected energy during each episode is the kinetic energy of the ejecta after energy dissipation during the prompt emission phase. Lacking a direct measure of this energy, we hereby assume that the emitted $\\gamma$-ray energy is a good proxy of the kinetic energy, so that $E_k = \\xi E_\\gamma$. In other words, we assume a constant radiative efficiency throughout the burst. We take $\\xi=1$ as the nominal value (i.e. 50\\% radiative efficiency, which may be achieved for efficient magnetic energy dissipation, Zhang \\& Yan 2011). In order to fit the data, we also allow $\\xi > 1$ for the GRBs, which corresponds to a less efficient dissipation mechanism (e.g. in internal shocks, Panaitescu et al. 1999; Kumar 1999; Maxham \\& Zhang 2009). To evaluate $\\gamma$-ray energy $E_\\gamma$ as a function of time, we divide the lightcurve into multiple time bins for each burst. For each time bin (with uneven duration denoted as $\\Delta T_i$ for $i$-th bin), we record its average flux $F_i$ in the GBM band, along with other useful information such as spectral parameters and the maximum photon energy. The total gamma-ray energy released in this time bin ($i$-th) is therefore \\begin{equation} E_{\\gamma,i} = \\f{4 \\pi d_L^2 F_i \\Delta T_i}{1 + z}. \\label{Energystart} \\end{equation} where $z$ is the redshift (see Table 1 for values of each burst), $d_L$ is the luminosity distance of the source, and the concordance cosmology with $\\Omega_{\\Lambda}=0.7$ and $\\Omega_m = 0.3$ is adopted in the calculation. Adopting $E_{k,i}=\\xi E_{\\gamma,i}$, we then progressively increase the total energy in the blastwave $E_k = \\Sigma E_{k,i}$ by adding $E_{k,i}$ in each step. For each time step, we calculate the lightcurve giving the available $E_k$. This results in a series of lightcurve solutions. The final lightcurve is then derived by jumping to progressively higher level solutions due to additional energy injections in each time step (see also Maxham \\& Zhang 2009). This would result in a series of ``glitches'' in the lightcurves, each representing injection of energy from $i$-th shell into the blastwave. Besides the energy, we also derive the (lower limit) Lorentz factor $\\gamma_i$ of each shell. This parameter is important, especially for early shells, since it determines the deceleration time of a certain shell. This is particularly relevant for the first shell. The Lorentz factors of later shells are also relevant fot two reasons. First, they can be used to calculate the effective Lorentz factor of a ``merged'' shell after adding energy to an existing shell. This is needed to calculate the deceleration time of the blastwave solutions. Second, since the observed time for a late energy injection is defined by (Maxham \\& Zhang 2009) \\begin{equation} t_{\\oplus,col}=t_{ej}+\\frac{(t_{col}-t_{ej})}{2\\gamma^2}~, \\label{collisiontime} \\end{equation} where $t_{ej}$ and $t_{col}$ are the times of ejection and collision measured in the rest frame of the central engine. The effect of $\\gamma$ becomes progressively less important, since at large $t_{ej}$'s, the second term in Eq.(\\ref{collisiontime}) becomes negligible so that the observed collision time is essentially defined by the ejection time. In any case, we derive the constraints on $\\gamma$ for each time bin using the pair opacity argument as described below. To derive a constraint on the Lorentz factor, we have collected the spectral parameters and the observed maximum photon energy $E_{\\rm{\\oplus,max},i}$ for each time bin. One can then derive the maximum photon energy in the cosmological local frame, i.e. $E_{\\rm{max},i}=E_{\\rm{\\oplus,max},i}(1+z)$. Requiring the pair production optical depth to be less than unity for $E = E_{\\rm{max},i}$, we can write a general constraint in the parameter space of $R$ and $\\gamma$ (where $R$ is the distance of the emission region from the central engine, Gupta \\& Zhang 2008; Zhang \\& Pe'er 2009), i.e. \\begin{equation} R(\\gamma) > \\sqrt{\\f{C(\\beta) \\sigma_T d_z^2}{-1-\\beta f_0} \\left(\\f{E_{\\rm{max}}}{511 \\rm keV^2}\\right )^{-1-\\beta} \\left(\\f{\\gamma}{1+z}\\right )^{2+2\\beta}}, \\label{Solveme} \\end{equation} where $\\sigma_T$ is the Thompson cross section, $\\beta$ represents the slope of the power law component for GRBs 090902B and 090510 and the Band function high energy spectral parameter for GRBs 080916C and 090926A, and $f_0$ (in units of $\\rm{ergs} \\cdot \\rm{cm}^{-2} \\cdot s^{-1}$) can be written as $f_0 =A \\cdot \\Delta T \\left[ \\f{E_p (\\alpha - \\beta)}{2 + \\alpha}\\right]^{\\alpha - \\beta} \\rm{exp}(\\beta - \\alpha)(100 \\quad \\rm{keV})^{-\\alpha}$ for the Band function model, and $f_0=K \\cdot \\Delta T (100 \\quad \\rm{keV})^{-\\beta}$ for the simple power law model, where $A$ and $K$ are normalization factors (both normalized to 100 keV). The approximation $\\rm{C}(\\beta) \\simeq (7/6)(- \\beta)^{5/3}/(1-\\beta)$ (Svensson 1987) is adopted to perform the calculation. In order to further constrain $\\gamma$, one needs to make an assumption about $R$. Without other independent constraints, we apply the conventional assumption of internal shocks, so that $R(\\gamma)=\\gamma^2 c \\f{\\delta t}{1+z}$, where $\\delta t$ is the observed minimum variability time scale. Combining Eq.(\\ref{Solveme}), the lower limit for $\\gamma$ is derived for each time bin of each burst (see also Lithwick \\& Sari 2001, Abdo et al. 2009). In our calculation, we generally adopt $\\gamma_i$ as the derived lower limit. This is because the derived Lorentz factors of other GRBs using the afterglow deceleration constraint (Liang et al. 2010) or photosphere constraint (Pe'er et al. 2011) are all below or consistent with these lower limits derived from the opacity constraints (Abdo et al. 2009a,b, 2010; Ackermann et al. 2010). \\subsection{Model results} Feeding this data into our shell model code, letting each shell be ejected with energy $E_{k,i}$ and Lorentz factor $\\gamma_i$ at time equal to that of the beginning of the bin time, we can calculate the early blastwave evolution and LAT band (integrated over $>100$ MeV) lightcurve for the four GRBs. To match the observed steep decay (with slope $\\sim -1.5$), we adopt a radiative fireball solution or an adiabatic fireball solution with steep electron energy index. Even though each solution (for a fixed kinetic energy) has a steep decay slope, the overall lightcurve shows a shallower decay due to piling up of successive shells ejected later, with glitches introduced by jumping among the solutions. As an example, the radiative model lightcurve of GRB 080916C as compared with observation is presented in Fig.1. The top panel shows the long term evolution, while the bottom panel is the zoomed-in early afterglow lightcurve. The dotted lines denote the blastwave solutions with progressively increasing total energy. The lowest one corresponds to the first time bin, the second lowest corresponds to adding the energy of the second time bin, etc. Since the lightcurve is chopped into discrete time bins, the blastwave energy is added in discrete steps. This introduces some artificial glitches in the lightcurve. Such an approximation is more realistic for GRBs with distinct emission episodes. For GRB 080916C, the lightcurve is more appropriately approximated as a continuous wind with variable luminosity. The artificial glitches should appear to be more smeared. For this reason, we have smoothed the glitches to make more natural transitions between solutions. The model afterglow parameters (the fraction of electron energy $\\epsilon_e$, the fraction of magnetic energy $\\epsilon_B$, and the number density $n$) are presented in Table 1. These are in general consistent with the parameter constraints derived by Kumar \\& Barniol Duran (2009, 2010). In general, the model lightcurve of GRB 080916C cannot fit the early LAT data. Making the model suitable to fit the late-time steep decay, the early model lightcurve level is too low to account for the observed data. Alternatively, one can make the early model lightcurve match the observed flux level. Then inevitably the late time afterglow level exceeds the observed level significantly due to the continuous energy injection. We believe that if the LAT band emission after $T_{90}$ originates from the external shock, then the LAT emission during the prompt emission phase {\\em cannot} be solely interpreted by the external shock model. The external shock contribution is relatively small, especially during early epochs when energy in the blastwave is small. As a result, the GeV emission during the prompt phase must be of an internal origin. This is consistent with the fact that the entire GBM/LAT emission during the prompt phase can be well fit by a single Band-function spectral model in all the time bins (Abdo et al. 2009a; Zhang et al. 2011). We have also modeled GRBs 090510, 090902B and 090926A. The model parameters (for both radiative and adiabatic solutions) are listed in Table 1, and the results for radiative solution are shown in Fig.\\ref{fluxlightcurves}. In all cases, the slope and flux level of the data are matched in the latter part of the curve only. During the prompt emission phase, the data points rise above the flux prediction of the external shock model, suggesting that GeV emission is a superposition of external and internal components during the prompt emission phase ($T100$ MeV lightcurve of GRB 080916C for a radiative blastwave solution (yellow line) as compared with the data (blue points). Successive lightcurves that correspond to different total blastwave kinetic energy are shown as dashed lines. The top panel shows the global lightcurve, while the bottom panel shows a zoom view where the flux deficit at early times can be clearly seen.} \\label{chopbw} \\end{figure} \\begin{figure} \\includegraphics[scale=0.8]{090510zoomrad.eps} \\includegraphics[scale=0.8]{090902Bzoomrad.eps} \\includegraphics[scale=0.8]{090926Azoomrad.eps} \\caption{Model predictions of $>100$ MeV lightcurve (for a radiative blastwave solution) vs. observed data for GRBs 090510, 090902B, and 090926A. The conventions are similar to Fig.1, but without successive solutions specifically plotted.} \\label{fluxlightcurves} \\end{figure} ", "conclusions": "Using the first-hand Fermi data, we have tracked the energy output from the central engine and modeled the early blastwave evolution of four bright LAT GRBs. The predicted $>100$ MeV lightcurve is found unable to account for the observed LAT emission during the prompt emission phase. The main reason is that during the phase when the central engine is still active, the forward shock is continuously refreshed by late energy injection, so that the afterglow decays much slower than the case predicted by an instantaneously ejected constant energy fireball. This suggests that at least during the prompt emission phase, the LAT band emission is not of external forward shock origin. This is in contrast to the suggestion of Ghisellini et al. (2009), Kumar \\& Barniol Duran (2009) and Feng \\& Dai (2010), who did not consider the energy accumulation during the prompt emission phase and interpreted the entire GeV emission as due to the external shock origin. Our conclusion is based on the assumption that GRB radiative efficiency is essentially a constant throughout the burst. In order to interpret the entire afterglow as due to the external forward shock origin, one needs to ``artificially\" assume that the GRB efficiency increases with time, so that the late time central engine activity, even though producing bright $\\gamma$-ray emission, adds little kinetic energy into the blastwave. We believe that such an assumption is contrived. Our conclusion is consistent with some independent arguments. From data analysis, Zhang et al. (2011) showed that during the prompt emission phase the GeV emission and MeV emission traces each other well. For GRB 080916C, the entire GBM/LAT emission can be modeled by a single Band function component in all the time bins (see also Abdo et al. 2009a). For GRB 090902B, even though GeV emission belongs to a distinct spectral component, its flux seems to track the flux of the MeV component nicely, suggesting a connection in the physical origin (see Pe'er et al. 2011 for modeling). A more definite argument in favor of an internal origin of GeV emission in GRB 080916C is that the GeV lightcurve peak coincides the second peak in the GBM lightcurve, suggesting that GeV emission is the spectral extension of MeV emission to higher energies (Zhang et al. 2011). Also individual case studies of GRB 090902B (Pe'er et al. 2011; Liu \\& Wang 2011) and GRB 090510 (He et al. 2011) all suggest that the external shock model cannot interpret the prompt GeV data. In general, our modeling suggests that it is possible to use the external shock model to interpret GeV emission after the prompt emission phase, but not during the prompt emission phase (see also Kumar \\& Barniol Duran 2010). Our conclusion also has implications for understanding GRB prompt emission physics, in particular, the composition of the GRB outflow. The internal origin of GeV emission in GRB 080916C makes it essentially impossible to interpret the entire Band spectrum with the photosphere model (e.g. Beloborodov 2010; Lazzati \\& Begelman 2010). The lack of photosphere emission then demands a Poynting-flux-dominated outflow at least for this burst (Zhang \\& Pe'er 2009; Fan 2010), and new models in the Poynting flux dominated regime (e.g. Zhang \\& Yan 2011) are called for. \\smallskip This work is partially supported by NSF AST-0908362 and NASA NNX09AT66G, NNX10AD48G. We thank Xue-Feng Wu for helpful discussion." }, "1101/1101.0002_arXiv.txt": { "abstract": "We perform 3-D dust radiative transfer (RT) calculations on hydrodynamic simulations of isolated and merging disk galaxies in order to quantitatively study the dependence of observed-frame submillimeter (submm) flux density on galaxy properties. We find that submm flux density and star formation rate (SFR) are related in dramatically different ways for quiescently star-forming galaxies and starbursts. Because the stars formed in the merger-induced starburst do not dominate the bolometric luminosity and the rapid drop in dust mass and more compact geometry cause a sharp increase in dust temperature during the burst, starbursts are very inefficient at boosting submm flux density (e.g., a $\\ga16$x boost in SFR yields a $\\la 2$x boost in submm flux density). Moreover, the ratio of submm flux density to SFR differs significantly between the two modes; thus one cannot assume that the galaxies with highest submm flux density are necessarily those with the highest bolometric luminosity or SFR. These results have important consequences for the bright submillimeter-selected galaxy (SMG) population. Among them are: 1. The SMG population is heterogeneous. In addition to merger-driven starbursts, there is a subpopulation of galaxy pairs, where two disks undergoing a major merger but not yet strongly interacting are blended into one submm source because of the large ($\\ga 15$'', or $\\sim 130$ kpc at $z = 2$) beam of single-dish submm telescopes. 2. SMGs must be very massive ($M_{\\star} \\ga 6 \\times 10^{10} \\msun$). 3. The infall phase makes the SMG duty cycle a factor of a few greater than what is expected for a merger-driven starburst. Finally, we provide fitting functions for SCUBA and AzTEC submm flux densities as a function of SFR and dust mass and bolometric luminosity and dust mass; these should be useful for calculating submm flux density in semi-analytic models and cosmological simulations when performing full RT is computationally not feasible. ", "introduction": "Submillimeter-selected galaxies (SMGs; \\citealt{Smail:1997,Barger:1998,Hughes:1998,Eales:1999}; see \\citealt{Blain:2002} for a review) are extremely luminous \\citep[bolometric luminosity $\\lbol \\sim 10^{12} - 10^{13} L_{\\odot}$; e.g.,][]{Kovacs:2006}, high-redshift \\citep{Chapman:2005} galaxies powered primarily by star formation rather than AGN \\citep{Alexander:2005,Alexander:2005b,Alexander:2008,Valiante:2007,Menendez:2007,Menendez:2009,Pope:2008MIR, Younger:2008phys_scale,Younger:2009EGS}. Because of their high dust content, SMGs emit almost all of their luminosity in the IR. As the name suggests, a galaxy is defined as an SMG if it is detected in the submm (historically, 850 \\micron ~flux density $S_{850} \\ga 3-5$ mJy; the nature of the population is sensitive to the adopted flux density cut, so we define an SMG as a source with $S_{850} > 3$ mJy), which requires $\\lir \\ga 10^{12} \\lsun$ \\citep{Kovacs:2006,Coppin:2008}, so SMGs are typically ultra-luminous infrared galaxies (ULIRGs). Locally, ULIRGs are almost exclusively merging galaxies \\citep{Sanders:1996,Lonsdale:2006}, so one might expect that at least some SMGs are also merging galaxies. Indeed, many observations support a merger origin for SMGs \\citep[e.g.,][]{Ivison:2002,Ivison:2007,Ivison:2010,Chapman:2003, Neri:2003,Smail:2004,Swinbank:2004,Greve:2005,Tacconi:2006,Tacconi:2008, Bouche:2007,Biggs:2008,Capak:2008,Younger:2008phys_scale,Younger:2010,Iono:2009,Engel:2010}. Furthermore, in \\citet[][hereafter N10]{Narayanan:2010smg} we combined hydrodynamic simulations and radiative transfer (RT) calculations to show that major mergers can reproduce the full range of submm flux densities and typical UV-mm spectral energy distribution (SED) of SMGs (cf. \\citealt{Chakrabarti:2008SMG}; \\citealt{Chakrabarti:2009}). Semi-analytic models also predict that the SMG population is dominated by merger-induced starbursts rather than quiescent star formation (\\citealt{Baugh:2005,Fontanot:2007,Swinbank:2008,LoFaro:2009,Fontanot:2010,Gonzalez:2011}; but cf. \\citealt{Granato:2004}). However, because of the much greater rate of gas supply onto galaxies at high redshift \\citep[e.g.,][]{Keres:2005,Dekel:2009nature}, gas fractions \\citep{Erb:2006,Tacconi:2006,Tacconi:2010,Daddi:2010} and star formation rates \\citep{Daddi:2007,Noeske:2007b,Noeske:2007a} of galaxies at fixed galaxy mass increase rapidly with redshift. Thus, at $z \\sim 2-3$ even a ``normal'' star-forming galaxy can reach ULIRG luminosities \\citep[e.g.,][]{Hopkins:2008cosm_frame1,Hopkins:2010IR_LF,Daddi:2005,Daddi:2007,Dannerbauer:2009}. Furthermore, roughly estimating submm counts using estimates of high-redshift major merger rates and the short duty cycle of merger-induced starbursts suggests that there may not be enough major mergers to account for the SMG population \\citep{Dave:2010}. This motivates the view that, instead, typical SMGs may be massive, gas-rich disks quiescently forming stars and fueled by continuous gas supply from mergers and smooth accretion (\\citealt{Carilli:2010}, but cf. \\citealt{Daddi:2009smgb}). The mode of star formation responsible for the majority of the SMG population is still a matter of debate, as it is difficult to discriminate between the two scenarios given the currently available data. A better understanding of the submm galaxy selection can clarify the nature of the SMG population. Since SMGs have redshifts $z \\sim 1-4$ \\citep{Dannerbauer:2002,Chapman:2005,Younger:2007high-z_SMGs,Younger:2008phys_scale,Capak:2008,Greve:2008,Schinnerer:2008, Daddi:2009smga,Daddi:2009smgb,Knudsen:2010}, the observed submm flux density traces rest-frame $\\sim 150 - 400$ \\micron, longward of the peak of the IR SED. Thus the observed submm flux density is sensitive to both the total IR luminosity and the ``dust temperature''\\footnote{As is convention, we will use the term ``dust temperature'' to denote the temperature derived from a single-temperature modified blackbody fit to the SED. This is simply a parameterization of the SED shape rather than a physical temperature. In our simulations dust grains have a continuum of temperatures, depending on both grain size and the local radiation field heating the dust.} of the SED, which depend on the luminosity from stars and AGN absorbed by the dust, the mass and composition of the dust, and the spatial distribution of stars, AGN, and dust. Galaxies do not have identical SED shapes, so the dependence on dust temperature implies that galaxies with the highest submm flux density are not necessarily those with the highest bolometric luminosity. Furthermore, because star formation histories are more complicated than an instantaneous burst, the luminosity and instantaneous SFR are not necessarily linearly proportional. Thus the relationship between submm flux density and SFR is potentially more complicated than the relationship between submm flux density and bolometric luminosity. We therefore cannot say \\emph{a priori} that the galaxies with the highest submm flux densities are the most rapidly star-forming or most luminous bolometrically. Indeed, it has already been observationally demonstrated that submm selection does not select all the brightest galaxies in a given volume, as there are galaxies with luminosities and redshifts comparable to those of SMGs that are undetected in the submm because of their relatively hot SEDs \\citep{Chapman:2004,Chapman:2010,Casey:2009,Casey:2010,Hwang:2010dust_T_evolution, Magdis:2010dust_T,Magnelli:2010}. A submm galaxy selection is clearly biased toward cold galaxies; however, the details of the selection bias are yet to be understood. Despite the basic physical reasons that one does not expect a simple relation between submm flux density and SFR, a linear relation between submm flux and SFR has been used explicitly (and, even more frequently, implicitly) to infer SFR from observed submm flux densities \\citep[e.g.,][]{Chapman:2000,Peacock:2000,Blain:2002,Scott:2002,Webb:2003,vanKampen:2005,Tacconi:2008,Wang:2011}, typically because the data sets do not have enough photometric data points to precisely constrain the IR SED shape (\\emph{Herschel} data are already helping greatly in this regard; e.g., \\citealt{Chapman:2010,Dannerbauer:2010,Magnelli:2010}). Furthermore, some theoretical studies \\citep{Dave:2010} have assumed that SMGs are the most rapidly star-forming galaxies in order to identify SMGs in cosmological simulations without performing RT. If SFR and submm flux density are not simply related this approach is problematic. It is clear that a better understanding of the relationship between submm flux density and SFR and, more generally, what galaxy properties a submm galaxy selection selects for, is needed. In other work we have combined hydrodynamic simulations and dust RT to show that major mergers of massive, gas-rich disk galaxies can reproduce the 850 \\micron ~flux densities (N10), CO properties \\citep{Narayanan:2009}, number densities \\citep[][C. Hayward et al.~2011, in preparation]{Hayward:2011num_cts_proc}, and intersection with the dust-obscured galaxy (DOG) population \\citep{Narayanan:2010dog} of SMGs. Motivated by the success of our simulations in reproducing a variety of SMG properties, we use them here to quantify how submm flux density depends on SFR, $\\lbol$, dust content, and geometry. The aim of this study is to clarify for what galaxy properties a submm selection criterion selects and to provide a discriminant among the different modes of star formation that could power SMGs. ", "conclusions": "We have combined high-resolution 3-D hydrodynamic simulations of high-redshift isolated and merging disk galaxies and 3-D Monte Carlo dust radiative transfer calculations to study the submillimeter galaxy selection, focusing on the relationships among submm flux density, star formation rate, bolometric luminosity, and dust mass. Our main conclusions are the following: \\begin{enumerate} \\item The relationship between SFR and submm flux density differs significantly for quiescent and starburst star formation modes. Starbursts produce significantly less submm flux density for a given SFR, and the scaling between submm flux density and SFR is significantly weaker for bursts than for quiescent star formation. Bursts are a very inefficient way to boost submm flux density (e.g., a starburst that increases SFR by $\\ga 16$x increases submm flux density by $\\la 2$x). Another consequence is that the galaxies with highest submm flux density are not necessarily those with highest SFR or bolometric or infrared luminosity. \\item The submm flux density of our simulations can be parameterized as a power law in SFR and dust mass ($\\lbol$ and dust mass) to within $\\sim 0.1 (0.15)$ dex. The scaling derived from the commonly used optically thin modified blackbody model systematically overpredicts the submm flux density by $\\ga 2$x because numerous assumptions of the model (optical thinness in the FIR, $\\lir \\propto$ SFR, $\\lbol \\approx \\lir$) do not hold. The fitting functions we provide (Equations \\ref{eq:sfr_fitting_functions} and \\ref{eq:L_fitting_functions}) should be useful for calculating the flux density in semi-analytical models and cosmological simulations when full radiative transfer cannot be performed and for interpreting observations. \\item Mergers create SMGs through another mechanism besides the strong starburst induced at coalescence---they cause the two infalling disks to be observed as one submm source because both disks will be within the large ($\\sim 15$'', or 130 kpc at $z = 2$) beam of the single-dish submm telescopes used to identify SMGs during much of the infall stage. For major mergers, this effect boosts the submm flux density by 2x. To achieve the same boost in submm flux density one would have to boost the SFR of a quiescent disk by $\\sim 6$x or induce a starburst that boosts the SFR by $\\ga 16$x. This implies that the SMG population is heterogeneous: it is composed of both late-stage major mergers and two (or more) infalling disks observed as a single submm source (``galaxy-pair SMGs''). The largest quiescently star-forming galaxies may also contribute. Thus, unlike local ULIRGs, SMGs are a mix of quiescent and starburst sources. \\item SMGs must be very massive: to reach $S_{850} \\ga 3$ mJy, stellar mass of at least $6 \\times 10^{10} \\msun$ is required, and typical values are higher. \\item The submm duty cycles of our simulated galaxies are a factor of a few longer than what one would expect if all SMGs were merger-driven bursts because the relatively gentle decline in SFR, $\\lbol$, and dust mass during the galaxy-pair phase results in a longer duty cycle for the galaxy-pair phase than for the starburst. The duty cycle of the latter is limited because the peak in luminosity is narrow and the dust temperature increases sharply during the burst. \\item Fitting the SEDs of SMGs with an optically thin modified blackbody tends to yield significantly lower dust temperatures than when the full opacity term is used because the effective optical depths can be $\\sim 1$ out to rest-frame $\\sim 200 ~\\micron$, both for our simulated SMGs and observed SMGs. Therefore, one should be cautious when interpreting effective dust temperatures derived via fitting an optically thin modified blackbody to the FIR SED, especially when comparing SMGs to galaxies for which optical thinness in the IR may be a reasonable approximation. \\end{enumerate} Future work will include predictions of submm number counts from our model, an investigation of the observational signatures and physical implications of the proposed SMG bimodality, and an improved method for fitting IR SEDs of galaxies." }, "1101/1101.3976_arXiv.txt": { "abstract": "The fractal shape and multi-component nature of the interstellar medium together with its vast range of dynamical scales provides one of the great challenges in theoretical and numerical astrophysics. Here we will review recent progress in the direct modelling of interstellar hydromagnetic turbulence, focusing on the role of energy injection by supernova explosions. The implications for dynamo theory will be discussed in the context of the mean-field approach. Results obtained with the test field-method are confronted with analytical predictions and estimates from quasilinear theory. The simulation results enforce the classical understanding of a turbulent Galactic dynamo and, more importantly, yield new quantitative insights. The derived scaling relations enable confident global mean-field modelling. ", "introduction": "Apart from stars, the baryonic matter within the Galaxy is in the form of an extremely dilute, turbulent plasma known as the interstellar medium (ISM). The multitude of physical processes within the ISM entails a rich heterogeneous structure \\citep{1978ppim.book.....S}. Approximating radiative processes by a simplified cooling prescription, and restricting the computational domain to a local patch, the turbulent ISM is now routinely modelled by means of three-dimensional fluid simulations \\citep[e.g.,][]{1999ApJ...514L..99K,2004A&A...424..817M,2005MNRAS.356..737S,% 2006ApJ...653.1266J,2006ApJ...638..797D}. One main focus of these simulations has been to obtain filling factors of the different ISM phases and compare them to the classical predictions as well as observations \\citep[e.g.,][]{1992FCPh...15..143D}. Further topics of interest include turbulent mixing \\citep{2005ApJ...634..390B}, thermodynamic distribution functions \\citep{2005ApJ...626..864M}, and line-of-sight integrated column densities \\citep{2005ApJ...634L..65D}. \\subsection{The small-scale dynamo} While various simulations \\citep{1999A&A...350..230K,2005A&A...436..585D,2005ApJ...626..864M} discuss the influence of magnetic fields on the ISM morphology, little is said about the actual mechanism of field amplification. \\citet{2004ApJ...617..339B} have addressed this question by means of unstratified simulations of SNe turbulence. The authors relate the growth of small-scale magnetic fields to vorticity production in supernova shocks \\citep{2001ApJ...563..800B}, and chaotic field line-stretching \\citep{2005ApJ...634..390B}. The fact that vorticity production by colliding shells is almost inevitable in a clumpy and highly structured ISM has first been pointed out by \\citet{1999A&A...350..230K}. The issue has then been investigated for the simplified case of driven expansion waves by \\citet{2006MNRAS.370..415M} and, more recently, by Del Sordo \\& Brandenburg (this volume). Considering turbulence driven by non-helical transverse waves, \\citet{2004MNRAS.353..947H} have shown that the small-scale dynamo becomes harder to excite in the super-sonic regime, albeit the critical Reynolds number for the onset of dynamo action only seems to depend weakly on the Mach number. Because the eddy turnover time is short at small scales, a dynamo based on chaotic field line stretching will be fast. This is in-line with observations \\citep[see][and this volume]{1996ARA&A..34..155B}, which exhibit dominant turbulent fields. Open issues remain with respect to the mechanism governing the saturation of the small-scale dynamo. Therefore, it is currently unclear whether equipartition field strengths can be obtained by a non-helical dynamo alone. Alternatively, the turbulent field might be explained as a ``shredded'' coherent field, i.e., as the by-product of a helical mean-field dynamo. ", "conclusions": "Because current numerical simulations are limited to very moderate Reynolds numbers, it is important to understand how efficient the observed mechanisms remain under realistic conditions. To achieve this, it has proven fruitful to study simplified scenarios and run multiple parameter sets \\citep[see][for a comprehensive review]{2005PhR...417....1B}. Fortunately, the emerging physical effects are dominated by the outer scale of the turbulence, i.e., as soon as a rudimentary scale separation is achieved, the turbulent quantities should become independent of the actual micro scale. The growing complexity of models challenges the distribution of computing time: increasing physical realism leaves little margin for the variation of key parameters, let alone convergence checks or running multiple representations of a single parameter set. Dedicated studies remain mandatory to segregate artificial trends from genuinely physical ones \\citep{2009A&A...498..335H,2009IAUS..259...81G}. The dilemma becomes even more apparent when looking at the recent trend to performing ``resolved'' global simulations. For these, convergence checks are a rare exception. While an ``enhanced'' diffusivity may be sufficiently approximated by the numerical truncation error on the grid scale, the diamagnetic pumping term certainly is not. Yet vertical transport has profound implications on the emerging dynamo modes and growth rates \\citep[see, e.g.,][]{2001A&A...370..635B}. In conclusion, we advocate a strategy that has been applied with great success in the design of aircraft, namely the concept of large eddy (or mean-field) simulations. Global fluid simulations currently cannot guarantee scale separation for all relevant physical scales. To obtain quantitatively correct results, we therefore believe that a sub-grid scale model is inevitable. Present local box simulations are a valuable means to provide a rigorous framework for the calibration of such a model." }, "1101/1101.1973_arXiv.txt": { "abstract": "We present new 1--1.25 $\\mu m$ (z and J band) Subaru/IRCS and 2 $\\mu m$ (K band) VLT/NaCo data for HR 8799 and a rereduction of the 3--5 $\\mu m$ MMT/Clio data first presented by \\citet{Hinz2010}. Our VLT/NaCo data yields a detection of a fourth planet at a projected separation of $\\sim$ 15 AU -- ``HR 8799e\". We also report new, albeit weak detections of HR 8799b at 1.03 $\\mu m$ and 3.3 $\\mu m$. Empirical comparisons to field brown dwarfs show that at least HR 8799b and HR8799c, and possibly HR 8799d, have near-to-mid IR colors/magnitudes significantly discrepant from the L/T dwarf sequence. Standard cloud deck atmosphere models appropriate for brown dwarfs provide only (marginally) statistically meaningful fits to HR 8799b and c for unphysically small radii. Models with thicker cloud layers not present in brown dwarfs reproduce the planets' SEDs far more accurately and without the need for rescaling the planets' radii. Our preliminary modeling suggests that HR 8799b has log(g) = 4--4.5, T$_{eff}$ = 900K, while HR 8799c, d, and (by inference) e have log(g) = 4--4.5, T$_{eff}$ = 1000--1200K. Combining results from planet evolution models and new dynamical stability limits implies that the masses of HR 8799b, c, d, and e are 6--7 M$_{J}$, 7--10 M$_{J}$, 7--10 M$_{J}$ and 7--10 M$_{J}$. ''Patchy\" cloud prescriptions may provide even better fits to the data and may lower the estimated surface gravities and masses. Finally, contrary to some recent claims, forming the HR 8799 planets by core accretion is still plausible, although such systems are likely rare. ", "introduction": "The HR 8799 planetary system is the first directly imaged multiplanetary system \\citep{Marois2008}. Along with Fomalhaut and $\\beta$ Pic, it is also the only imaged system with companion mass ratios and separations reasonably close to the giant planets in the Solar System \\citep{Kalas2008,Lagrange2009,Lagrange2010}\\footnote{Here, we consider 1RXJ1609.1-210524b discovered by \\citet{Lafreniere2008a} to be a more complicated case as its mass ratio and separation are continuous with brown dwarf companions (see Discussion Section)}. After the initial detection of HR 8799bcd, one or more planets were recovered in prior datasets \\citep{Lafreniere2009, Fukagawa2009, Metchev2009}. Recently, \\citet{Marois2011} imaged a fourth planet -- HR 8799e -- which we independently detected (see \\S 2). Mass estimates based on cooling models yield 5--11 M$_{J}$ for HR 8799b and 7--13 M$_{J}$ for the other planets \\citep{Marois2008,Marois2011}. Dynamical constraints placed by HR 8799bcd imply that the companions likely have masses below the deuterium-burning limit \\citep[][]{Spiegel2010} and are kept stable by resonant interactions \\citep{Fabrycky2009,MoroMartin2010}. Including the fourth planet, \\citet{Marois2011} argue that the planets most likely have masses at the low end of the range allowed by cooling models. With masses of $\\approx$ 5--13 M$_{J}$, the HR 8799 planets then bridge the gap between the solar system's gas giants/Jupiter-mass planets detected by radial velocity surveys \\citep[e.g.][]{Howard2010} and low-mass brown dwarf companions to nearby stars such as GJ 758B and PZ Tel \\citep{Thalmann2009,Currie2010a,Biller2010}. Recent studies complicate our understanding of the relationship between brown dwarfs, the gas giants detected in RV surveys, and the HR 8799 planets. The planets' masses are significantly larger than most planets detected by radial velocity and transit methods. \\citet{Marois2008} noted that the planets appear slightly redder than the distribution of H/H-K$_{s}$ colors for old field brown dwarfs. The K-band spectrum of HR 8799b is not well matched by typical L and T-type brown dwarf spectra \\citep{Bowler2010}. Comparisons between the HR 8799 planet photometry/spectroscopy and atmosphere models reveal additional difficulties in understanding their properties within the theoretical framework of standard, cloud deck models that track the field L/T dwarf sequence. In the discovery paper, \\citet{Marois2008} briefly mention a discrepancy between temperatures derived from atmosphere models and those estimated from more simple, and presumably most accurate, cooling model estimates. More recently, \\citet{Bowler2010} provide a detailed comparison between the HR 8799b spectra and 1.1--4.1 $\\mu m$ photometry and predictions from standard atmosphere models. They show that the `best-fit' temperatures derived from modeling are inconsistent with cooling model estimates. They also explicitly show that the implied radii for best-fit models are well below the 1.1--1.3 R$_{J}$ range allowed by standard cooling models (e.g. 0.3--0.5 R$_{J}$). To interpret these modeling difficulties, \\citet{Bowler2010} argue that a different atmospheric structure, namely atmospheres with stronger cloud coverage, may better explain the HR 8799b SED. Since atmospheric dust entrained in clouds absorbs more efficiently at shorter wavelengths, photometry for HR 8799b at wavelengths shortward of J band would provide a crucial test of the planet's level of cloud coverage \\citep[cf.][]{Burrows2006}. The \\citet{Bowler2010} study also found difficulty in reconciling their model fitting of detections from \\citet{Marois2008} with 3--5 $\\mu m$ upper limits from \\citet{Hinz2010}. More sensitive photometry at these wavelengths would then provide better modeling constraints. In this study, we investigate the atmospheres and dynamics of the HR 8799 planets using new observations obtained at the Subaru Telescope and VLT and a rereduction of MMT data presented by \\citet{Hinz2010}. Combined with photometry presented by \\citet{Marois2008}, our data yield nine photometric points spanning 1--5 $\\mu m$ for a detailed comparison to the properties of field brown dwarfs. This wavelength range also provides a sensitive probe of the effects of surface gravity, temperature, (non)equilibrium chemistry, metallicity, and cloud coverage. We compare the planets' SEDs to atmosphere models exploring a phase space defined by these effects. By quantifying the model fits, we determine the range of parameter space that fails to characterize the planets' SEDs and identify the subset of models that more accurately reproduce the data and may better represent their atmospheres' physical properties. These results will then be used to more thorougly and accurately probe the planets' atmospheric properties in a companion paper (Madhusudhan et al. 2011, in prep.). Our study is structured as follows. \\S 2 describes our observations, image processing, and detections for each dataset. The first part of \\S 3 compares the HR 8799 planet photometry to the L/T dwarf sequence and the IR properties of other very low-mass objects (M $<$ 25 M$_{J}$). The rest of \\S 3 presents preliminary comparisons between the HR 8799 planet SEDs and planetary atmosphere models. \\S 4 describes simple dynamical modeling of the system to identify the range of masses for dynamically stable orbits. \\S 5 summarizes our results, discusses our work within the context of previous studies of HR 8799 and planet imaging in general, discusses how our results fit within the context of planet evolution models, and comments on the plausible formation mechanism(s) for the planets. ", "conclusions": "Our primary result in this paper is that the atmospheres of at least two and potentially all of the HR 8799 planets do not easily fit within the empirical IR color sequence for L/T type brown dwarfs of similar temperatures, nor can they be well fit by standard atmosphere models used to infer the properties of brown dwarfs. Adopting realistic assumptions about planet radii, all atmosphere model fits to data for HR 8799b and c are far poorer than any meaningful threshold identifying models consistent with the data. The models primarily fail by underpredicting the 3.3 $\\mu m$ flux and badly overpredicting flux at 1--1.3 $\\mu m$. Our analysis suggests that having ``thicker\" clouds --ones with larger vertical extents -- is key to reproducing the planets' SEDs. Compared to cloud structures assumed in standard L/T dwarf atmosphere models, these clouds are more optically thick at a given T$_{eff}$, so they are visible (in the photosphere) at a lower T$_{eff}$ even though the cloud base is located far below at much higher pressures. Adopting a thick cloud layer prescription, we succeed in identifying models for each planet that quantitatively are good-fitting models. Moreover, the temperatures of these models are consistent with simpler, presumably more accurate cooling model estimates. \\subsection{Comparisons with Previous Studies of HR 8799} The most direct comparison to this work is the recent analysis of the HR 8799b K-band spectrum and 1.1--4.1 $\\mu m$ photometry from \\citet{Bowler2010} whose modeling formalism we largely follow. \\citet{Bowler2010} also finds difficulties in using standard atmosphere models to fit HR 8799b's SED and interpret its properties \\citep[see also][]{Marois2008}. Likewise, they find that temperatures inferred from standard atmosphere models disagree with cooling model predictions and that the former require unphysically small radii. Our results indicate that including Y/z band data only exacerbates the already serious disagreement between standard cloud deck model predictions and the planet's SED. Our analysis confirms \\citet{Bowler2010}'s inference that HR 8799b's atmosphere is exceptionally dusty compared to field brown dwarfs. Our results extend this inference, indicating that HR 8799c and, plausibly, HR 8799d are also dusty compared to field brown dwarfs. \\citet{Janson2010} noted that while standard atmosphere models -- the COND models in their case -- can reproduce the mean brightness of HR 8799c's L'-band spectrum they incorrectly predict the spectral slope from 3.9 $\\mu m$ to 4.5 $\\mu m$. They cite greater atmospheric dust absorption and, especially, non-equilibrium carbon chemistry as features that may bring the models into better agreement. \\citet{Hinz2010} argue that incorporating non-equilibrium chemistry is necessary to reproduce the mid-IR photometry of HR 8799bcd since the chemical equilibrium models they use \\citep{Saumon2006} predict M-band fluxes larger than the upper limits they report. Non-equilibrium carbon chemistry has little effect on the near-IR portion of the SED \\citep[e.g.][]{HubenyBurrows2007}. Thus, our analysis indicates that thicker clouds -- and, by implication, stronger atmospheric dust absorption -- are far more important than non-equilibrium chemistry in reproducing the HR 8799 planet 1--5 $\\mu m$ SEDs. Nevertheless, the HR 8799 planet atmospheres are plausibly not in local chemical equilibrium. Since departures from chemical equilibrium alter the spectral structure at 4--5 $\\mu m$, non-equilibrium chemistry incorporated into thick or ``patchy\" cloud models may yield better fits to 1--5 $\\mu m$ photometry and mid-IR spectroscopy of the planets. Higher signal-to-noise L' band spectra and detections/more stringent upper limits at M will better identify evidence of non-equilibrium chemistry in the planets' atmospheres. \\subsection{Comparisons with Planet Evolution Models and Implied Masses} Within the context of the \\citet{Burrows1997} planet cooling models, a particular combination of log(g) and T$_{eff}$ defines an object with a mass M and age t. Taking the gravity and temperature range implied by our modeling at face value, we can then identify the mass and age range implied. Our modeling efforts succeed in yielding planets with physically realistic radii. However, if our range of log(g) and T$_{eff}$ were to imply wildly discrepant masses compared to cooling model estimates and dynamical stability requirements or widely varying ages our analysis would have solved one problem only to create comparably serious ones. Here, we combine all modeling results to identify the range of best-fit parameters and implied parameters -- mass and age -- from atmosphere models that we consider. We then determine whether the atmospheric and dynamical modeling constraints are consistent and, if so, what mass and age range they imply. \\begin{itemize} \\item {HR 8799b} -- The minimum $\\chi^{2}$ value for HR 8799b for thick cloud models is 27.6 if we allow the radius to vary by up to 10\\% from the \\citet{Burrows1997} values and 48.9 if we don't. For the ``patchy\" cloud approximation, the corresponding $\\chi^{2}$ minima are 20.6 and 51.4. Considering the best-fit models passing the $\\Delta \\chi^{2}$ threshold in each case, this range covers log(g) = 4--4.5 and T$_{eff}$ = 800--1000K. Thus, our modeling yields log(g) = 4--4.5, T$_{eff}$ = 800--1000K. Using the \\citet{Burrows1997} evolutionary models, this implies a mass and age range of M, t = 5 M$_{J}$, 30 Myr to 15 M$_{J}$, 300 Myr. \\item {HR 8799c, d, and e} -- The minimum $\\chi^{2}$ values here for thick cloud models are 43.5 and 60.7 for c and 5.7 and 5.3 with and without radius rescaling. For the ``patchy\" cloud approximation, the corresponding $\\chi^{2}$ minima are 14--14.1 for c and 2.8--7.4 for d. For HR 8799c, the range of models passing the $\\Delta \\chi^{2}$ threshold for the thick and patchy cloud prescriptions cover log(g) = 4--4.5 and T$_{eff}$ = 1000K--1200K. This yields a mass/age range of 7 M$_{J}$, 30 Myr to 15--17.5 M$_{J}$ at 150--300 Myr. For HR 8799d, the range is log(g) = 3.75--4.5, T$_{eff}$ = 1000-1200K, yielding 5 M$_{J}$ at 10 Myr to 15--17.5 M$_{J}$ at 150--300 Myr. Since HR 8799e likely has a bolometric luminosity and K-L colors comparable to HR 8799c and d, its range of masses is plausibly consistent with those derived for HR 8799c and d. \\end{itemize} Dynamical constraints require that HR 8799b is less than 7 M$_{J}$ and HR 8799cde are less than 10 M$_{J}$ \\citep[Section 4 of this work;][]{Marois2011}. The 5 M$_{J}$ mass estimate for HR 8799d can be ruled out because the primary star is on the main sequence and thus cannot be 10 Myr old. Coupled with the range in surface gravities and temperatures, the implied range in masses are then 6--7 M$_{J}$ for HR 8799b, 7--10 M$_{J}$ for HR 8799c, and 7--10 M$_{J}$ for HR 8799 d. If HR 8799e's atmospheric properties mirror those of c and d, its plausible range of masses is also 7--10 M$_{J}$. Conversely, for these ranges of masses, the surface gravities of HR 8799bcde should be no greater than log(g) $\\approx$ 4.25. These estimates are consistent with cooling model estimates from \\citet{Marois2008,Marois2011}. For the lower end of the mass ranges, the system age corresponding to these models is $\\approx$ 30 Myr and puts HR 8799's age on the low end of the 30--160 Myr range quoted by \\citet{Marois2008}. The (disfavored) high end of the mass range corresponds to $\\sim$ 100 Myr-old objects. Despite our success in arriving at self-consistent answers for the planets' masses and ages, we strongly caution against overinterpreting these results. Our results do not \\textit{prove} that, above the cloud base, the vertical density/pressure profile of clouds follows that of the gas as a whole (e.g. s$_{1}$ = 0), as opposed to being truncated at higher pressures. Neither do our results prove that other models with slightly different assumptions about the clouds, grain particles, atmospheric chemistry, etc. provide better fits to the data. In particular, slight modifications to our models may improve the fit at L' band, the datapoint responsible for much of the $\\chi^{2}$ contribution for HR 8799b. Even within the context of our adopted physical models, our sampling in temperature and gravity is also too coarse to precisely estimate best-fit atmosphere parameters. On the other hand, our analysis provides compelling evidence for thick clouds, motivates future modeling work to test how different assumptions about thick clouds affect model fits to planetary atmospheres, and encourages further observations of substellar objects to test these models. Madhusudhan et al. (2011) will develop and better assess model fits for varying cloud strengths and more precisely and accurately determine temperatures and gravities for the HR 8799 planets and other planetary-mass objects. \\subsection{Constraints On The Formation of the HR 8799 Planetary System} The planets' large masses and wide orbits make them a particularly interesting probe of planet formation. The favored theory invoked to explain the formation of gas giant planets is \\textit{core accretion} \\citep[e.g.][]{Mizuno1980,Pollack1996,KenyonBromley2009, Chambers2010}, where cores that have grown to $\\approx$ 5--10 M$_{\\oplus}$ rapidly accrete much more massive gaseous envelopes. Alternatively, planets could form by disk instability \\citep[][and later papers]{Boss1997}, where the protoplanetary disk is massive and gravitationally unstable, forming multiple self-gravitating clumps of gas that coalesce into bound, planet-mass objects. HR 8799's planets are often described as confounding either planet formation model \\citep[e.g.][]{Marois2011} or being clear examples of disk instability-formed planets, as claimed by \\citet{DodsonRobinson2009}. They find that cores at distances characterizing the HR 8799 planets cannot reach $\\sim$ 10 M$_{\\oplus}$ in mass to undergo runaway gas accretion \\textit{even under the most favorable conditions}. They claim that planet-planet scattering cannot create stable, wide-orbit systems like HR 8799's. They conclude that massive, wide-separation gas giants like HR 8799bcd form by disk instability and ''can certainly rule out core accretion\". Critical to \\citeauthor{DodsonRobinson2009}'s conclusion is their treatment of the core growth rate. The growth rate strongly depends upon the planetsimal approach velocity, which they fix at v$_{a}$ = $\\Omega$R$_{hill}$. They claim this velocity yields an ``optimistically high\" growth rate. Their formalism implicitly assumes that planetesimals have an isotropic velocity dispersion (v$_{a}$ $\\sim$ v$_{z}$), which is valid as long as the scale height of planetesimals accreted by cores (v$_{z}$/$\\Omega$) is larger than the core's impact parameter, R$_{core}$$\\sqrt{(1+\\theta)}$ \\citep[][]{Rafikov2004}, where $\\theta$ is the Safranov number. However, if the planetesimals are dynamically cold such that v$_{z}$ $\\le$ $\\sqrt{p}$$\\Omega$R$_{Hill}$ (where p = R$_{core}$/R$_{Hill}$), this condition is violated \\citep{Dones1993,Rafikov2004}. The core can then accrete the entire vertical column of planetesimals at a vastly higher rate since accretion is now essentially two-dimensional \\citep{Rafikov2004}. As a result, \\citet{DodsonRobinson2009} catastrophically underestimate the maximum growth rate by a factor of p$^{-1/2}$, or up to 114, 85, and 68 at the positions of HR 8799b, c, and d (cf. Equations 78, 80, and 82 in Rafikov 2004; see also Rafikov 2010)\\footnote{At first glance, Equation (17) in \\citet{Rafikov2010} appears to imply that the limiting distance for core accretion in shear-dominated growth is comparable to \\citeauthor{DodsonRobinson2009}'s estimate (44 AU vs. their 20--35 AU). However, \\citeauthor{Rafikov2010}'s result of 44 AU is valid for a Minimum Mass Solar Nebula case \\citep{Hayashi1981}. Adopting initial assumptions more comparable to those that \\citeauthor{DodsonRobinson2009} assumes -- e.g. a disk more massive than the Minimum Mass Solar Nebula or a longer-lived one with $\\tau_{disk}$ = 5 Myr instead of 3 Myr-- implies that gas giants can in some cases form by core accretion at separations comparable to HR 8799c and b.}. Detailed numerical simulations confirm that this rapid growth phase can be reached if collisional fragmentation and gas drag are properly treated \\citep{KenyonBromley2009}. The \\citeauthor{DodsonRobinson2009} planet-planet scattering simulations also were conducted assuming gas free, planetesimal-free conditions and assumed that planets could not further grow after scattering. However, gas drag and dynamical friction from planetesimals are critically important as they promote orbit circularization and stability \\citep[e.g.][]{Goldreich2004,FordChiang2007}\\footnote{In fairness, they clearly acknowledge that their study does not consider planet-planet scattering \\textit{in a gaseous disk}, which may result in a more favorable outcome for core accretion.}. Cores with masses sufficient for rapid gas accretion can circularize after being scattered to the outer disk \\citep[][S. Kenyon 2010, pvt. comm.]{BromleyKenyon2011}. Simulations by Thommes et al. (in prep.) show that the HR 8799 planet cores could acquire most of their gas \\textit{after} scattering. The mass ratio and semimajor axis distribution of wide planets and low-mass brown dwarfs may help constrain the formation mechanism for HR 8799's planets \\citep{Kratter2010}. Core accretion preferentially forms planets with smaller masses and orbital separations, while disk instability has difficulty producing lower-mass gas giants and forming them close to the star \\citep[e.g.][]{Rafikov2005,Kratter2010}. Therefore, if HR 8799bcde formed by core accretion (disk instability), they should comprise the high-mass extrema (low-mass tail) of a population continuous with radial-velocity detected planets (brown dwarf companions). Using our new results for the masses of the HR 8799 planets, we update \\citeauthor{Kratter2010}'s plot comparing planet and brown dwarf distributions. We also add the planet-mass companions to 1RXJS1609.1-210524, and 2M J044144b \\citep[5--10 M$_{J}$, 15 AU]{Todorov2010}; the planet/brown dwarf companion to GSC 06214-00210B \\citep[14 M$_{J}$, $\\sim$ 300 AU][]{Ireland2010}; and the low-mass brown dwarf companion GJ 758B \\citep[25--40 M$_{J}$, 44 AU][]{Currie2010a}. As shown by Figure \\ref{kratterplot}, the revised masses for the HR 8799 planets and the addition of HR 8799e expand the space between them and brown dwarf companions (asterisks). Visually, they join with the distribution of closer-separation planets plausibly formed by core accretion. The other new companions (red triangles) are continuous with brown dwarfs that may form by disk fragmentation. While core accretion -- especially when coupled with planet-planet scattering -- may form the HR 8799 planetary system, HR 8799-like systems are still plausibly uncommon. The Gemini Deep Planetary Survey of 85 nearby, young (mostly solar-mass) stars was typically sensitive to 2 M$_{J}$ planets at 40--200 AU yet failed to detect any \\citep{Lafreniere2007b}. Similarly, non-detections from the deep (M $<$ 1 M$_{J}$) survey from \\citet{Kasper2007} showed that the giant planet populations detected at small separations (a $\\lesssim$ 4 AU) by RV surveys cannot extend to separations larger than $\\sim$ 30 AU. More massive stars like HR 8799 likely have more massive disks, which aid gas giant planet formation. However, their disks also dissipate much more rapidly \\citep{CurrieLada2009}: even if critical core masses are reached, the leftover mass of gas may be small. Moreover, rapid core growth results from being able to fragment and then dynamical cool the surrounding planetesimal population. The current state-of-the-art simulations show that this requires Pluto-mass cores to start with \\citep[e.g.][]{KenyonBromley2009}, yet the formation time for Pluto-mass objects at wide separations may be long \\citep[e.g.][]{Rafikov2010}. Thus, forming HR 8799-like systems by core accretion is difficult, though \\textit{not} impossible, and probably happens infrequently. \\subsection{Implications for the Atmospheres of Other Substellar Companions: A Possible Fundamental Difference Between Planetary-Mass Objects and M $>$ 15--20 M$_{J}$ Brown Dwarfs} In some ways, the difficulty in reproducing the IR SEDs of the HR 8799 planets mirror difficulties in modeling other planetary-mass objects detected prior to HR 8799bcde. In particular, 2M 1207b also appears discrepant compared to standard atmosphere models as noted in \\citet{Mohanty2007} and discussed in this work. Like HR 8799b, 2M 1207b is noticeably underluminous ($\\sim$ 2.5 mags) in the near-IR \\citep[][ this work]{Mohanty2007}. To explain 2M 1207b's properties, \\citet{Mohanty2007} propose that the object is occulted by an edge-on disk with large, gray dust grains. Alternatively, \\citet{Mamajek2007} propose that 2M 1207b's properties can be explained as resulting from a recent protoplanet-protoplanet collision. Comparing high-resolution spectra of 2M 1207b to the DUSTY atmosphere models from \\citet{Allard2001}, \\citet{Patience2010} identify a problem similar to that noted for modeling HR 8799b from \\citet{Bowler2010} and this work. Namely, allowing the object radius to freely vary yields best-fit radii far smaller ($\\approx$ 0.5 R$_{J}$) than is physically plausible \\citep[cf.][]{Burrows1997}. \\citet{Patience2010} also conclude that extinction from an edge-on disk comprised of gray dust grains is also a viable scenario. For the same reasons -- underluminosity/red colors -- a disk origin also has been proposed to explain the IR SED of HR 8799b and (to a lesser extent) c and d \\citep{Marois2008}. However, \\citet{Marois2008} consider the chance alignment of an edge-on circumplanetary disk to be unlikely, especially given that the system is viewed nearly face on. Even more unlikely is the chance that circumplanetary disks or recent protoplanet collisions explain the near-IR properties of two to four separate planets in two systems with very different ages and primary star properties. Given the success in better reproducing HR 8799bcd's SEDs with thick cloud models and the similarity between HR 8799b and 2M 1207b, it is more plausible that the latter's near-IR spectrum is likewise explained by thick clouds. If this is generally true of planetary-mass objects, thicker clouds may constitute the primary difference between the atmospheres of massive planets and brown dwarfs, at least over the gravity and temperature range enclosed by the HR 8799 planets and 2M 1207b (e.g. log(g) = 3.75--4.5, T$_{eff}$ =900-1600K). Since thicker clouds affect the color-magnitude positions of substellar objects it is quite possible the Model A 'thick cloud' sequence extending to HR 8799b and 2M 1207b from the nominal L/T dwarf boundary continues on to even cooler temperatures (e.g. T$_{eff}$ $\\sim$ 700--900K). Since thick clouds present reshape the spectral structure at $\\sim$ 1.6 $\\mu m$ (e.g. in the methane band), they may also affect the L/T dwarf transition, which is already known to be dependent upon surface gravity \\citep[e.g.][]{Metchev2006,Luhman2007}. \\subsection{Future Work} Our study motivates the development of a suite of new atmosphere models with clouds intermediate in thickness between the Model E cloud deck and Model A thick cloud layer prescriptions. Adopting these models as fiducial models, we can revisit the (secondary) effects of surface gravity, metallicity, and non-equilibrium chemistry on the atmospheres of planetary mass objects, complementing similar investigations for brown dwarfs \\citep[][]{Allard2001,Marley2002,Burrows2006,HubenyBurrows2007}. These models will be developed and applied to HR 8799bcde and other planetary-mass objects in upcoming papers (Madhusudhan et al. 2011, in prep.) and may provide a useful comparison to planet parameters derived from cooling models \\citep[e.g.][]{Burrows1997, Baraffe2003,Fortney2007,Fortney2008}. New observations at 1--5 $\\mu m$ will provide better constraints on the HR 8799 planet atmospheres. In addition to more sensitive data at Y band and [3.3], Figure \\ref{modseq} (lower-left panel) implies that thick-cloud atmospheres may have \\textit{far} stronger emission at $\\sim$ 2.3 $\\mu m$ and 3.0 $\\mu m$ than standard models predict. This wavelength range can be probed for at least HR 8799bcd by current ground-based facilities such as VLT/NaCo, Keck/NIRC2, and MMT/Clio. Integral field spectrographs on \\textit{Gemini Planet Imager} \\citep[GPI][]{MacIntosh2008} and SPHERE \\citep{Beuzit2008} will sample the 1--2.5 $\\mu m$ SED region with exceptional sensitivity and thus provide a detailed comparison between observed and predicted atmospheric properties of all planets. Finally, ongoing collaborations such as the IDPS survey (Marois et al., in progress) and Gemini/NICI \\citep{Liu2010} will better probe the frequency of wide, massive ($\\sim$ 5--13 M$_{J}$, $>$ 30 AU) around nearby stars. GPI and SPHERE will probe 1--5 M$_{J}$ planets at even smaller separations (e.g. 5--30 AU). These surveys will produce a far more complete census of Jupiter-mass planets to better determine their ubiquity and constrain how the formation of planets like HR 8799's compare to that expected for lower-mass planets at smaller separations and wide-separation, low-mass brown dwarfs." }, "1101/1101.4598_arXiv.txt": { "abstract": "{ We performed three dimensional resistive magnetohydrodynamic simulations to study the magnetic reconnection using an initially shearing magnetic field configuration (force free field with a current sheet in the middle of the computational box). It is shown that there are two types of reconnection jets: the ordinary reconnection jets and fan-shaped jets, which are formed along the guide magnetic field. The fan-shaped jets are much different from the ordinary reconnection jets which are ejected by magnetic tension force. There are two driving forces for accelerating the fan-shaped jets. The one is the Lorentz force which dominates the motion of fluid elements at first and then the gas pressure gradient force accelerates the fluid elements in the later stage. The dependence on magnetic reconnection angle and resistivity value has also been studied. The formation and evolution of these jets provide a new understanding of dynamic magnetohydrodynamic jets. ", "introduction": "Magnetic reconnection plays a very important role in solar flares, coronal mass ejections, and other solar activities. The problem of flare energy release has been puzzling people for many years before the development of magnetic reconnection mechanism. The first pioneers in this field of magnetic reconnection are \\cite{Sweet1958} and \\cite{Parker1957}, who developed a theory well known as Sweet-Parker mechanism. However, it can not explain the solar flares. The biggest problem in their theory is that the time scale of energy release is much longer than the realistic one of flares. Several years later, \\cite{Petschek1964} improved this work by considering a pair of slow mode shocks outside a very small diffusion region, which greatly increases the reconnection rate. The fast reconnection model is successful, though some people do not believe that the large energy release process has such a small diffusion region. The famous CSHKP model (\\citealt{Carmichael1964, Sturrock1966, Hirayama1974, Kopp1976}) based on magnetic reconnection for flares has been developed since 1960s. Besides observations, magnetohydrodynamic (MHD) simulations play a key role in studying the magnetic reconnection (\\citealt{Ugai1977, Forbes1983, Forbes1990, Forbes1991, Yokoyama1994, Magara1996, Ugai1996, Chen1999}). With the development of computer and observation technology, the CSHKP flare model has been extended. More and more simulations show that the plasmoid ejection may greatly increase the magnetic reconnection rate (\\citealt{Shibata1995, Shibata1996, Shibata1997, Magara1997, Nishida2009}). The plasmoids can form in an anti-parallel magnetic field. Before ejection magnetic energy is restored in the diffusion region, and simultaneously the plasmoids merge with each other. Finally, the restored magnetic energy is released in a very short time scale after the ejection of the big plasmoid. Recently, the three dimensional (3D) simulation is one of the hottest points in astrophysics. In this case, the reconnecting process is much more complicated than that in two dimensional (2D) case. The 3D simulations can provide us very complex and realistic structures (\\citealt{Yokoyama1995, Isobe2005, Ugai2005, Isobe2006, Shimizu2009}). It is also a good tool to explain the mechanism of many small scale solar activities (\\citealt{Cirtain2007, Shibata2007, Katsukawa2007}) which have been observed by Hinode (\\citealt{Kosugi2007}). Some of these observed jetlike features are believed to be a result of magnetic reconnection in the so-called interlocking-comb like magnetic configuration in sunsport penumbra (i.e., penumbral jets). The motivation of this paper is to study the effect of the guide field using 3D numerical simulations. In our new simulation, we found that some parts of the reconnection jets which are ejected from the diffusion region can move along the magnetic guide field lines in the interlocking-comb like magnetic configuration. Moreover, the reconnection jets which move along the guide field lines are almost perpendicular to the ambient magnetic field. In order to distinguish the jets moving in different directions, we call the jets moving along the ambient magnetic field as ordinary reconnection jets and the jets moving along the guide field as fan-shaped jets (we will see the shape of the jets is similar to a fan in the next section). Besides our simulations, two very recent paper by~\\cite{Pontin2005} and~\\cite{Ugai2010} also discussed the effect of guide field. Their simulations showed that the guide field can distort the propagation of 3D plasmoids and the plasma is almost ejected along the guide field. It seems that the guide field plays a very important role for forming these jets in the 3D magnetic reconnection. In our first paper (\\citealt{Jiang2011}) we briefly reported the main results of our 3D simulation, while in this paper the detailed analysis and discussion of fan-shaped jets are presented. This paper is organized as follows. A detailed description of the magnetohydrodynamics (MHD) equations and the initial condition are given in the next sections. In Section~\\ref{typical} we show the 3D, 2D and 1D structures of our typical simulation results. Moreover, the analysis of the Lagrangian fluid elements, reconnection rate and slow mode shock are also given. Section \\ref{dependence} shows dependence on different reconnection angles (defined in the initial conditon) and different values of resistivity in the diffusion region. Finally, discussion and summary are given in the last Section. ", "conclusions": "\\label{discussion} We perform three dimensional resistive MHD simulations to study the magnetic reconnection using an initially shearing magnetic field configuration, which is similar to the papers by~\\cite{Pontin2005} and~\\cite{Ugai2010}. It is shown that there are two types of reconnection jets: the ordinary reconnection jets which move along the ambient magnetic field lines, and the fan-shaped jets which move along the guide field lines. In this paper we analyse the fan-shaped jets in detail. This may provide a new understanding of the dynamic phenomena in astrophysics. One of the possibility is that it can be applied to interpret the solar sunsport penumbral jets (\\citealt{Katsukawa2007}) or anemone jets (\\citealt{Shibata2007}) in the solar chromosphere. Actually, before the fan-shaped jets, there is a slow mode wave propagating along the magnetic field lines. Because of its fast speed this wave can be seen as the front of the fan-shaped jets. But it is very hard to distinguish which is the slow mode wave and which are the jets. There is no clear boundary between them. The part which has gas pressure increasement but with small velocity is the slow mode wave as shown in Figure~\\ref{fig02}. Of course, the ordinary jets connect the fan-shaped jets, the slow mode wave also exists at the front of ordinary reconnection jets. The driving force for the fan-shaped jets is another interesting point. As there is no gravity in our simulations, the analysis of the Lagrangian fluid elements is easy to understand. Before Lagrangian fluid elements move along the magnetic field lines, both the gas pressure gradient force and the Lorentz force can drive the elements. The key point is that the magnetic tension force dominates the element movement first and then the element is affected by the gas pressure gradient force. In case 1 listed in Table~\\ref{cases} with a small total grid points, we incidentally found that a fluid element is only accelerated by the Lorentz force. That is due to the initial position of that element being at the upper edge (outflow side) of the diffusion region. It goes out of the diffusion region soon, so that only the Lorentz force affects this element. The dependence of the simulation variabes (namely density, pressure, temperature and velocity) on resistivity and reconnection angle is reasonable. The measurement of reconnection rate is simple in our paper and needs to be improved. However, from this method, people can get a general understanding about the 3D magnetic reconnection rate. It is roughly but useful, at least we can know when the magnetic reconnection reaches the maximum. Another disadvantage in our simulation is the limitation of the total grid points. A finer grid mesh or a higher order finite difference scheme or algorithm is under consideration. If the effect of gravity is considered, the density stratification effect in the solar atmosphere can greatly increase velocity amplitude of the upward fan-shaped jets when it propagates up to the low density region (like in the solar upper chromosphere and the corona) and eventually becomes a shock or jets (\\citealt{Shibata1982a, Shibata1982b}). The extended development of this study with gravity considered is in progress. In summary, we give the conclusions as follows: 1. From 3D MHD simulations we have studied a type of jets which move along the guide magnetic field. Because of the rotation of the initial magnetic field these jets have a fan-shaped structure. Of course, the ordinary reconnection jets ejected by the magnetic tension force also exist. 2. The driving forces of these fan-shaped jets are Lorentz force and gas pressure gradient force. The magnetic pressure gradient force drives the Lagrangian fluid elements first, and then the force on the elements is dominated by the gas pressure gradient in the later stage. 3. The pressure, density, temperature and velocity do not sensitively depend on the initial resistivity value. 4. The ratio between the velocity of the fan-shaped jets and the velocity of the ordinary reconnection jets is almost a constant being about 0.5 for different reconnection angles." }, "1101/1101.4551_arXiv.txt": { "abstract": "{Solar explosive events are commonly explained as small scale magnetic reconnection events, although unambiguous confirmation of this scenario remains elusive due to the lack of spatial resolution and the statistical analysis of large enough samples of this type of events.} {In this work, we propose a sound statistical treatment of data cubes consisting of a temporal sequence of long slit spectra of the solar atmosphere. The analysis comprises all the stages from the explosive event detection to its characterization and the subsequent sample study.} {We have designed two complementary approaches based on the combination of standard statistical techniques (Robust Principal Component Analysis in one approach and wavelet decomposition and Independent Component Analysis in the second) in order to obtain least biased samples. These techniques are implemented in the spirit of letting the data speak for themselves. The analysis is carried out for two spectral lines: the C~{\\sc iv} line at 1548.2 \\AA~ and the Ne~{\\sc viii} line at 770.4 \\AA.} {We find significant differences between the characteristics of the line profiles emitted in the proximities of two active regions, and in the quiet Sun, most visible in the relative importance of a separate population of red shifted profiles. We also find a higher frequency of explosive events near the active regions, and in the C~{\\sc iv} line. The distribution of the explosive events characteristics is interpreted in the light of recent numerical simulations. Finally, we point out several regions of the parameter space where the reconnection model has to be refined in order to explain the observations.} {} ", "introduction": "Explosive events are localised energy release episodes detected mainly as broad emission lines in solar transition region lines. They were first discovered and classified by \\cite{1983ApJ...272..329B} using the {\\sl High Resolution Telescope and Spectrograph}, HRTS. Their properties, were summarized by \\cite{1989SoPh..123...41D} and \\cite{1994AdSpR..14...13D}. \\citep{1995SoPh..162..189W} has provided a wealth of detailed observations in a wavelength range overlapping that of the HRTS, no update of the statistical picture has been carried out based on the new data. SUMER observations have been used in combination with other SOHO and Earth-based instruments to explore the relationship of explosive events with the magnetic field evolution \\citep{1998ApJ...497L.109C, Teriaca:2004a, 2003A&A...403..731M}; to provide a coherent picture of explosive events in relation with other transient events such as blinkers and/or surges \\citep{2009ApJ...701..253M, 2005A&A...432..307B, 2003A&A...403..731M, 2003A&A...403..287P, 2000ApJ...528L.119C, 1998ApJ...504L.123C}; and to explore specific aspects of the explosive event phenomenon, like the timing and variations in lines of different formation temperatures in \\citet{2005A&A...431..339M}, and the comparison of the signatures of explosive events in various lines of ions with similar formation temperatures in \\citet{2005A&A...439.1183D}. The theoretical work developed in the past decade in order to explain the observed non-gaussianity of the explosive events line profiles has converged in the framework of magnetic reconnection \\citep{1997Natur.386..811I}. Recent examples of numerical simulations of explosive events in this scenario can be found in \\citet{2010A&A...510A.111D}, \\citet{2009ApJ...702....1H}, \\citet{2009A&A...495..953L}, and \\citet{2006SoPh..238..313C}. In this paper, we propose an automatized procedure for the detection and analysis of explosive events. We aim at studying their properties in sufficiently large samples, and compare them with predictions from the models, in the hope that, by pointing at the discrepancies between observations and the numerical simulations, we can help refine the models of magnetic reconnection. The outline of this paper is as follows: in Sect. \\ref{obs} we describe the observations used to test the validity of the techniques proposed in Sect. \\ref{sec:methodology} for the detection and analysis of the explosive events line profiles; in Sect. \\ref{sec:results} we describe the results of applying these techniques to the SUMER data, and in Sect. \\ref{discussion}, we discuss the general properties of the explosive events samples thus obtained, and the match between the observed properties and the simulations. ", "conclusions": "In this work we have presented a totally automated processing of homogeneous series of spectral images taken by the SUMER spectrograph on board SOHO. As a result, we have several samples of explosive events in the two lines of C~{\\sc iv} and Ne~{\\sc viii}, in regions of the quiet Sun and in the outer parts of two active regions. The main objective of this work was to advance in the analysis of explosive events by looking at the general properties of the samples rather than studying individual cases. The main conclusions from our work are that \\begin{itemize} \\item even though we have significantly enlarged the sample size of explosive events by automating its detection, no clear and unique picture emerges, capable of describing the variety of line profiles encountered and described in Sect. \\ref{propee}; \\item current numerical simulations fail to explain the existance of explosive events line profiles with blue components much brighter than their red counterparts, and/or with maximum Doppler velocities in the red wing much larger than in the blue one; \\item the characteristics of explosive events near active regions as compared to their quiet Sun counterparts, and in particular their maximum velocities, cannot be explained only as the result of enhanced magnetic flux available for reconnection. \\end{itemize} This work leaves many unanswered questions. In particular, we have not analysed the time evolution of the explosive events, nor have we correlated the position of explosive events near the two active regions in the C~{\\sc iv} and Ne~{\\sc viii} lines. We intend to explore these issues in the future, together with the latitude dependence of the sample properties or the impact of the time resolution by comparing this dataset with others available in the SOHO data archive." }, "1101/1101.3371_arXiv.txt": { "abstract": "We used the Nobeyama 45-m telescope to conduct a spectral line survey in the 3-mm band (85.1--98.4 GHz) toward one of the nearest galaxies with active galactic nucleus NGC 1068 and the prototypical starburst galaxy NGC 253. The beam size of this telescope is $\\sim$ 18$^{\\prime\\prime}$, which was sufficient to spatially separate the nuclear molecular emission from the emission of the circumnuclear starburst region in NGC 1068. We detected rotational transitions of C$_{2}$H, cyclic-C$_{3}$H$_{2}$, and H$^{13}$CN in NGC 1068. These are detections of carbon-chain and carbon-ring molecules in NGC 1068. In addition, the C$_{2}$H {\\it N} = 1--0 lines were detected in NGC 253. The column densities of C$_{2}$H were determined to be 3.4 $\\times$ 10$^{15}$ cm$^{-2}$ in NGC 1068 and 1.8 $\\times$ 10$^{15}$ cm$^{-2}$ in NGC 253. The column densities of cyclic-C$_{3}$H$_{2}$ were determined to be 1.7 $\\times$ 10$^{13}$ cm$^{-2}$ in NGC 1068 and 4.4 $\\times$ 10$^{13}$ cm$^{-2}$ in NGC 253. We calculated the abundances of these molecules relative to CS for both NGC 1068 and NGC 253, and found that there were no significant differences in the abundances between the two galaxies. This result suggests that the basic carbon-containing molecules are either insusceptible to AGN, or are tracing cold ($T_{\\rm rot} \\sim$ 10 K) molecular gas rather than X-ray irradiated hot gas. ", "introduction": "To date, about 40 molecular species have been identified in nearby external galaxies (e.g. list available in CDMS\\footnote{The Cologne Database for Molecular Spectroscopy (CDMS) can be accessed at http://www.astro.uni-koeln.de/cdms/.}; M\\\"{u}ller et al. 2005). As a result, it is possible to study molecular abundances and chemical reactions in external galaxies. For example, a difference in molecular abundance has been reported between the nearby starburst galaxies NGC 253 and M 82 (e.g. Mauersberger \\& Henkel 1991, Takano et al. 1995), and the reason for this difference has been discussed. They suggested that the temperature in NGC 253 is higher than that in M 82 and/or that the size of the high-temperature region in NGC 253 is much larger than that in M 82. In addition, Takano et al. (2002) pointed out the peculiarity of the molecular abundance in M 82 among nearby starburst galaxies. Some groups have indicated that the specificity of the molecular abundances of M 82 is because of the chemistry of photon-dominated regions (PDRs) (e.g. Garc\\'{i}a-Burillo et al. 2002; Mart\\'{i}n et al. 2009). Different galaxies have different properties and activities. These include those with small or large amounts of gas, having a starburst and/or active galactic nucleus (AGN), and with interaction, etc. Molecular line observations of these different galaxies allow us to study the effects of these different properties/activities on the molecular medium. In fact, some groups have suggested that it is possible to diagnose power sources in dusty galaxies using molecular line ratios (e.g. Kohno et al. 2001, Usero et al. 2004, Kohno 2005, Imanishi et al. 2007, Kohno et al. 2008, Krips et al. 2008). The observation of the molecular gas chemistry of the AGN toward NGC 1068, one of the nearest galaxies with an AGN, has been reported by Usero et al. (2004), P\\'{e}rez-Beaupuits et al. (2009), and Garc\\'{i}a-Burillo et al. (2010). They observed molecular lines of HCO, H$^{13}$CO$^{+}$, SiO, HCO$^{+}$, HOC$^{+}$, and CN, and concluded that the circumnuclear disk of NGC 1068 is a giant X-ray dominated region (XDR). There are theoretical studies on the XDR model by Lepp \\& Dalgarno (1996), Maloney et al. (1996), and Meijerink et al. (2005; 2007). However, further systematic observations of molecular lines are indispensable to study the impact of AGN on the interstellar medium. Nevertheless, no systematic unbiased scans of millimeter/submillimeter molecular lines toward AGN have been published to date, although such scans do exist for two nearby starburst galaxies (NGC 253, Mart\\'{i}n et al. 2006; M 82, Naylor et al. 2010). Therefore, we started a project to conduct a line survey in the 3-mm band of NGC 1068 using the 45-m telescope at Nobeyama Radio Observatory (NRO) in Japan. The beam size of this telescope (18$^{\\prime\\prime}$ at 86 GHz) is smaller than the size of the circumnuclear starburst ring in NGC 1068 (d $\\sim$ 30$^{\\prime\\prime}$; e.g. Planesas et al. 1991) , and it is therefore essential to study the impact of the AGN on the surrounding molecules; this will enable us to mitigate the contamination of the molecular lines from the circumnuclear starburst region in NGC 1068. We also observed NGC 253 to compare the effects of AGN on molecular abundance. This is an on-going project, and we present here the initial results of the survey. In this study, we report detections of the rotational transition lines of C$_{2}$H, cyclic-C$_{3}$H$_{2}$, and H$^{13}$CN in NGC 1068, and detection of the C$_{2}$H {\\it N} = 1--0 transition in NGC 253 as well as other previously reported lines. The C$_{2}$H molecule was first detected in interstellar clouds by Tucker et al. (1974). However, extragalactic detections are limited to just a few objects (M82, Henkel et al. 1988; NGC 4945, Henkel et al. 1990; NGC 253, Mart\\'{i}n et al. 2006). The C$_{2}$H molecule is important to study the formation and the characteristics of carbon-chain molecules, because it is related to the carbon-chain growth. The cyclic-C$_{3}$H$_{2}$ molecule has been observed in a large number of galactic sources. In external galaxies, it has been detected in Cen A (Seaquist \\& Bell 1986; Bell \\& Seaquist 1988), NGC 253, and M82 (Mauersberger et al. 1991). Information on the abundances of basic carbon-containing molecules is important to understand the carbon chemistry in NGC 1068. The detection of H$^{13}$CN in the external galaxies has only been reported toward NGC 253 by Mauersberger \\& Henkel (1991). It is interesting that the isotope of HCN was detected in NGC 1068, because the isotope ratio gives the constraint on the optical depth of the HCN line; this will help us to unveil the nature of the over-luminous HCN emission at the center of NGC 1068. ", "conclusions": "We started line surveys in the 3-mm band toward NGC 1068 and NGC 253. The main results include detections of C$_{2}$H, cyclic-C$_{3}$H$_{2}$, and H$^{13}$CN toward NGC 1068 and the C$_{2}$H {\\it N} = 1--0 transition toward NGC 253. We calculated the relative abundances of C$_{2}$H and cyclic-C$_{3}$H$_{2}$ with respect to CS in both galaxies and found no significant differences in the results between NGC 1068 and NGC 253. Thus, it is concluded that these basic carbon-containing molecules are insusceptible to AGN and/or these molecules exist in a cold gas away from AGN. We continue an unbiased line survey toward NGC 1068 and NGC 253 and examine the abundances of other molecules in order to study the effects of AGN on the interstellar medium. A clear separation of the circumnuclear disk from the starburst ring is important to understand the abundances of the molecular gas around AGN. With the advent of ALMA, we will be able to study the AGN chemistry with much higher spatial resolution." }, "1101/1101.1360_arXiv.txt": { "abstract": "A specific class of flat Emergent Universe (EU) is considered and its viability is tested in view of the recent observations. Model parameters are constrained from Stern data for Hubble Parameter and Redshift ($H(z)$ vs. $z$) and from a model independent measurement of BAO peak parameter. It is noted that a composition of Exotic matter, dust and dark energy, capable of producing an EU, can not be ruled out with present data. Evolution of other relevant cosmological parameters, viz. density parameter ($\\Omega$), effective equation of state (EOS) parameter ($\\omega_{eff}$) are also shown. ", "introduction": "It is known from observational cosmology that our universe is passing through a phase of acceleration. Unfortunately, the present phase of acceleration of the universe is not clearly understood. Standard Big Bang cosmology with perfect fluid assumption fails to accommodate the observational fact. However, an accelerating universe is permitted if a small cosmological constant ($\\Lambda$) be included in the Einstein's gravity . There is, however, no satisfactory theory that explains the origin of $\\Lambda$ which is required to be unusually small. Moreover, Standard Big Bang model without a cosmological constant is inevitably pleagued with a time like singularity in the past. The Big Bang model is also found to be entangled with some of the observational features which do not have explantion in the framework of perfect fluid model. Consequently an inflationary epoch in the early universe is required \\citep{alt4} to resolve the outstandinng issues in cosmology. It is not yet understood when and how the universe entered the phase. However, the concept of inflation is taken up to build a consistent scenario of the ealy universe. Inflation may be realized in a semiclassical theory of gravity where one requires an additional inputs like existence of a scalar field which describes the matter in the universe. An alternative approach is also followed where gravitational sector of the Einstein field equation is modified by including higher order terms in the Einstein-Hilbert action \\citep{b19}. To address the present accelerating phase of the universe once again attempts are made where theories with a modification of the gravitational sector taking into account higher order terms that are relevant at the present energy scale are considered. There are other approaches generally adopted considering modification of the matter sector by including very different kind of matter known as exotic matter namely, Chaplygin gas and its variations \\citep{b3,b4}, models consisting one or more scalar field and tachyon fields \\citep{b18}. While most of these models address dark energy part of the universe, other models based on non-equilibrium thermodynamics and Boltzmann formulation, which do not require any dark energy \\citep{b15,b16,b17}, are also considered suitable for describing late universe. A viable cosmological model should accommodate an inflationary phase in the early universe with a suitable accelerating phase at late time. An interesting area of cosmology is to consider models which are free from the initial singularity also. Emergent Universe (EU) scenario is one of the well known choices in this field. EU models are proposed in different framework like Brans-Dicke theory \\citep{b151}, brane world cosmology \\citep{b1,b2,b5}, Gauss-Bonnet modified gravity \\citep{b13}, loop quantum cosmology \\citep{alt3} and standard General Relativity (GR) \\citep{b12}. Some of these models are implemented in a closed universe \\citep{alt2} while others in a flat universe \\citep{b12}. If EU be developed in a consistent way it might solve some of the well known conceptual problems not understood in the Big-Bang model. An interesting class of EU model in the standard GR framework has been obtained by \\citet{b12} considering a non-linear equation of state in a flat universe. The EU model evolves from a static phase in the infinite past into an inflationary phase and finally it admits an accelerating phase at late time. The universe is free from initial singularity and large enough to begin with so as to avoid quantum gravity effects. The non-linear equation of state is the input of the model which permits different composition of matter in addition to normal matter as cosmic fluid. The model has been explored in a flat universe as such universe is supported by recent observations. The EOS considered in obtaining EU model by \\citet{b12} is \\begin{equation} \\label{eos1} p=A\\rho-B\\rho^\\frac{1}{2}, \\end{equation} where $A$ and $B$ are unknown parameters with $B>0$ always. Different values of $A$ and $B$ corresponds to different composition of matters in the EU model. In the literature \\citep{w2}, similar kind of non-linear EOS has been considered as a double component dark energy model and analyzed to obtain acceptable values of model parameters. The EOS given by eq. (\\ref{eos1}) is a special form of a more general EOS, $p=A\\rho -B\\rho^{\\alpha}$; which permits Chaplygin gas as a special case (with $\\alpha <0$) \\citep{b3,b4}. Chaplygin gas is considered widely in recent times to build a consistent cosmological model. It interpolates between a matter dominated phase and a de Sitter phase. Later various modified forms of Chaplygin gas were proposed \\citep{b10} to track cosmological evolution. For example models like Modified Chaplygin gas interpolates between radiative era and $\\Lambda$CDM era. \\citet{w2} showed in their work that such interpolation is permissible even with $\\alpha>0$ and a string specific configuration may be phenomenologically realized with an EOS considered by \\citet{b12}. Recently using eq. (1) for an EU model proposed by \\citet{b12}, we determined various constraints that are imposed on the EOS parameters from observational data namely, SNIa data, BAO peak parameter measurement and CMB shift parameter measurement \\citep{alt1}. It was noted that an EU model is permitted with $A<0$. It is found that the possibility of $A=0$ case is also permitted when we probe the contour diagram of $A-B$ plane with 95 $\\%$ confidence. The case $A=0$ corresponds to a composition of dust, exotic matter and dark energy in the universe which is certainly worth exploring. In this paper a specific EU model is taken up where the matter energy content of the universe comprises of dust, exotic matter and dark energy. Using Stern data (Table. 1), the admissibility of model parameters are determined from $H(z)$ vs. $z$ \\citep{b14} and using measurement of model independent BAO peak parameter $\\mathcal A$. We also plotted evolution of cosmologically relevant parameters in our model. The paper is organized as follows : in section 2 field equations for the model are discussed, in section 3 and 4 we the model is constrained with Stern data and Stern+BAO data respectively. Finally in section 5 the findings are summarized with a discussion. \\begin{table} \\begin{minipage}{140mm} \\caption{Stern Data ($H(z) vs. z$)} \\begin{tabular}{l|c|r} \\hline {\\it z Data} & $H(z)$ & $\\sigma$ \\\\ \\hline 0.00 & 73 & $ \\pm $ 8.0\t \\\\ 0.10 & 69 & $ \\pm $ 12.0 \\\\ 0.17 & 83 & $ \\pm $ 8.0 \\\\ 0.27 & 77 & $ \\pm $ 14.0 \\\\ 0.40 & 95 & $ \\pm $ 17.4 \\\\ 0.48 & 90 & $ \\pm $ 60.0 \\\\ 0.88 & 97 & $ \\pm $ 40.4 \\\\ 0.90 & 117 & $ \\pm $ 23.0 \\\\ 1.30 & 168 & $ \\pm $ 17.4 \\\\ 1.43 & 177 & $ \\pm $ 18.2 \\\\ 1.53 & 140 & $ \\pm $ 14.0 \\\\ 1.75 & 202 & $ \\pm $ 40.4 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table} ", "conclusions": "\\begin{figure} \\label{effeos} \\includegraphics[width=240pt,height=200pt]{effeos.eps} \\caption{(Colour Online) Evolution of $\\omega_{eff}$ with best fit values and values within different confidence level} \\end{figure} \\begin{figure} \\label{rhop} \\includegraphics[width=240pt,height=200pt]{rho.eps} \\caption{(Colour Online) Evolution of the matter-energy density in a asymptotically de Sitter universe} \\end{figure} In this paper considering a very specific model of flat EU, we determine the observational constraints on the model parameters. For this recent observational data namely, Stern data, measurement of BAO peak parameter are used. The specific form of EOS given by eq. (1) to obtain EU scenario in a flat universe is employed here for the purpose. We set $A=0$ in the eq. (1) to begin with. $A$ equal to zero represents a universe with a composition of exotic matter only. This kind of EOS has been considered in \\citet{noz}. In our previous work \\citep{alt1} on EU model it is noted that a small non zero value of $A$ (although zero is not ruled out) is permitted. As a result the analysis was done with non zero $A$. In this paper since we are interested in a specific composition of matter energy content of the universe corresponding to $A=0$ anlysis is carried out for EOS given by (1) with $A=0$ only. As suggested by \\citep{b12}, it corresponds to the content of the universe which is a composition of dark energy, dust and exotic fluid> The above composition is reasonable to obtain a viable scenario of the universe considering the observational facts. It seems that the exotic part of the EOS may also contribute in the budget of dark energy content of the universe. We found that the observationally favoured amount of dark energy present in universe today $\\Omega_{\\Lambda} \\approx 0.72$ is permitted in our model within $68.3 \\%$ confidence level. However, it may be mentioned here that the model may be extended even if $ |A|<<1$ and so that we can write $A+1 \\approx 1$. We found that a composition of dust, exotic matter and dark energy may produce an EU model within the framework of Einstein's gravity with a non-linear equation of state. It is also noted that this kind of model can accommodate many other composition of matter energy depending on the value of $A$. The viability for those will be taken up elsewhere." }, "1101/1101.1967_arXiv.txt": { "abstract": "We investigate the rotation periods of fully convective very low mass stars (VLM, $M<0.3\\,M_{\\odot}$), with the aim to derive empirical constraints for the spindown due to magnetically driven stellar winds. Our analysis is based on a new sample of rotation periods in the main-sequence cluster Praesepe (age 600\\,Myr). From photometric lightcurves obtained with the Isaac Newton Telescope, we measure rotation periods for 49 objects, among them 26 in the VLM domain. This enlarges the period sample in this mass and age regime by a factor of 6. Almost all VLM objects in our sample are fast rotators with periods $<2.5$\\,d, in contrast to the stars with $M>0.6\\,M_{\\odot}$ in this cluster which have periods of 7-14\\,d. Thus, we confirm that the period-mass distribution in Praesepe exhibits a radical break at $M\\sim 0.3-0.6\\,M_{\\odot}$. Our data indicate a positive period-mass trend in the VLM regime, similar to younger clusters. In addition, the scatter of the periods increases with mass. For the $M>0.3\\,M_{\\odot}$ objects in our sample the period distribution is probably affected by binarity. By comparing the Praesepe periods with literature samples in the cluster NGC2516 (age $\\sim 150$\\,Myr) we constrain the spindown in the VLM regime. An exponential rotational braking law $P \\propto \\exp{(t/\\tau)}$ with a mass-dependent $\\tau$ is required to reproduce the data. The spindown timescale $\\tau$ increases steeply towards lower masses; we derive $\\tau \\sim 0.5$\\,Gyr for 0.3$\\,M_{\\odot}$ and $>1$\\,Gyr for 0.1$\\,M_{\\odot}$. These constraints are consistent with the current paradigm of the spindown due to wind braking. We discuss possible physical origins of this behaviour and prospects for future work. ", "introduction": "\\label{intro} The spin of stars is a strong function of stellar mass and age. The age-dependence for main-sequence F-K-type stars has been empirically established in the seminal paper by \\citet{1972ApJ...171..565S} as $\\omega \\propto t^{1/2}$. Originally found from rotational velocities in the Pleiades and Hyades, this relation still holds asymptotically on the main-sequence when evaluated with the large sets of rotation periods in open clusters that is currently available. Most recent tests of the Skumanich law tend to give slightly higher power law exponents of 0.56 (Collier Cameron et al. 2009) or 0.52 \\citep{2007ApJ...669.1167B}. From the theory side, the Skumanich law has been reproduced in the prescription provided by \\citet{1988ApJ...333..236K}, which is based on the analytical wind model by \\citet{1984LNP...193...49M}. Under plausible assumptions (linear dynamo, magnetic field a mixture between dipolar and radial), the Kawaler expression simplifies to $dJ/dt \\propto \\omega^3$, which gives the desired $\\omega \\propto t^{-1/2}$ behaviour. Recent numerical work, however, indicates that the Kawaler-type wind parameterisation may not be an adequate explanation for the empirically found Skumanich law \\citep{2008ApJ...678.1109M}. F-K-type stars exhibit a well-studied rotation-mass relation on the main-sequence. For example, at the age of the Hyades the rotation periods increase steadily towards later spectral types, from 5\\,d for late F-stars to 12\\,d for late K-stars \\citep{1987ApJ...321..459R}. This relation is remarkably tight and it seems possible to explain the few outliers as tidally locked binaries or as objects with specific spot configurations resulting in a wrong period measurement. Thus, for these objects mass and age essentially fix the rotation rate, which allows for the possibility of 'gyrochronology', i.e. measuring ages from rotation periods \\citep{2007ApJ...669.1167B}. Observations have not been able yet to establish similarly robust age/mass-rotation dependencies for the very low mass stars in the M-type regime. It is clear that the F-K-type period-mass relation breaks down in the early-M regime, corresponding to a mass threshold of 0.3-0.5$\\,M_{\\odot}$. This is most readily seen from the M-dwarfs periods in Praesepe \\citep{2007MNRAS.381.1638S}, which are 1-3\\,d, much shorter than in the K-type regime, and from the rotational velocity data, which indicates a significant increase in the rotation rate between early to mid M-types \\citep[e.g.,][]{1998A&A...331..581D,2009ApJ...704..975J}. Similarly, the Skumanich-type rotational braking does not hold anymore for VLM objects with $M<0.3\\,M_{\\odot}$. While angular momentum losses occur in this mass regime as well, the stars tend to maintain high rotation rates over Gyrs, which is not consistent with the $\\omega \\propto t^{-1/2}$ spindown. Most commonly the VLM spindown is empirically described with an exponential braking law $\\omega \\propto \\exp{(-t/\\tau)}$. This exponential behaviour is primarily motivated by the theoretical framework by \\citet{1988ApJ...333..236K}, see above. In a modification suggested by \\citet{1995ApJ...441..865C}, stars above a critical threshold $\\omega >\\omega_{\\mathrm{crit}}$ are treated with $dJ/dt \\propto \\omega_{\\mathrm{crit}}^2 \\omega$, which results in an exponential spindown law. To be able to match the period data in open clusters, $\\omega_{\\mathrm{crit}}$ has to be assumed to be a function of mass \\citep[e.g.][]{1997ApJ...480..303K,2007MNRAS.377..741I}. Empirically, however, the form of the spindown law is poorly constrained. For recent reviews on these subjects, see \\citet{2009AIPC.1094...61S} and \\citet{2009IAUS..258..363I}. In this paper we set out to investigate the period-mass and period-age relation for fully convective very low mass stars based on a new set of rotation periods measured for members of the open cluster Praesepe. Praesepe, at an age of $\\sim 600$\\,Myr, is an important cluster to constrain the spindown law, because the effect of wind braking can be studied in isolation. So far, however, only a very small sample of 4 periods was available for evolved VLM stars in open clusters \\citep{2007MNRAS.381.1638S}. Our goal here is to provide a quantitative measurement of the VLM spindown law based on a significantly larger sample of periods. ", "conclusions": "The analysis in Sections \\ref{rotmass} and \\ref{rotage} can be summarised as follows: \\begin{enumerate} \\item{At ages of 600\\,Myr, VLM objects with masses below $0.3\\,M_{\\odot}$ are almost exclusively fast rotators with periods of $<2.5$\\,d. This is in stark contrast to higher mass stars (0.6-1.2$\\,M_{\\odot}$) which have periods of 7-14\\,d and results in a sharp break in the period-mass relation at 0.3-0.6$\\,M_{\\odot}$.} \\item{In the VLM regime, the periods as well as the scatter in the periods increases with mass. The scatter is significantly larger than for 0.6-1.2$\\,M_{\\odot}$ stars.} \\item{Between 100 and 600\\,Myr VLM objects experience angular momentum losses. However, a single spindown law cannot explain the evolution of upper and lower period limit simultaneously. Instead, the fast rotators exhibit less rotational braking than the slow rotators. The exponential spindown timescale increases steeply from $\\sim $500\\,Myr for 0.3$\\,M_{\\odot}$ to several Gyrs at 0.1$\\,M_{\\odot}$.} \\end{enumerate} These results are in line with the currently used models for the rotational evolution. As outlined in Sect. \\ref{intro}, most recent papers on this subject use a Skumanich-type $P \\propto t^{0.5}$ law for the slow and an exponential law $P \\propto \\exp{(t)}$ for the fast rotators. This approach is supported by our new data. In particular, the periods in Praesepe show that VLM objects at 600\\,Myr rotate much faster than F-K type stars and have not arrived yet to the Skumanich-type spindown tracks. The physical origin of the empirical findings outlined above is a matter of debate. The breakdown of the Skumanich law around $\\sim 0.3-0.6\\,M_{\\odot}$ can possibly be understood as a consequence of interior structure. Going from solar-mass stars down to the VLM regime, the convective zone deepens until the objects become fully convective around 0.35$\\,M_{\\odot}$ \\citep{1997A&A...327.1039C}. Assuming that magnetic field generation and properties are a function of interior structure, this could qualitatively explain why VLM objects spend longer on the exponential track than more massive stars. The mass-dependence of the spindown timescale might actually be a dependence on $T_{\\mathrm{eff}}$, as argued in \\citet{2004PhDT.........4S}. With decreasing temperature, the electrical conductivity of the photospheric gas drops as well and the coupling between gas and magnetic field becomes less efficient. This might affect the mass load of the flux tubes and the efficiency of the stellar wind, and explain the strong increase in the spindown timescale at very low masses. It would also provide an explanation for the universally fast rotation of the ultracool L dwarfs \\citep{2008ApJ...684.1390R}. The positive period-mass trend at very low masses is already seen in very young clusters \\citep[e.g.][]{2005A&A...429.1007S}, albeit with much more scatter than in Praesepe. Thus, this feature is likely a remnant of the initial conditions. More detailed theoretical work on the magnetic field generation, wind physics, and their connection to angular momentum loss is required to substantiate this interpretation. Observationally, the current database still has two major flaws that need to be addressed: a) Too few clusters have been monitored with good time sampling and depth. As a result, we are lacking observational constraints in the substellar regime and cannot fully exclude to be affected by bias in the period range and possible environmental effects on rotation, as recently reported by \\citet{2010MNRAS.403..545L}. b) Follow-up observations are required for cluster objects with known rotation periods, to exclude contaminating field stars and obtain complementary information about the stellar and magnetic field properties. \\begin{figure*} \\includegraphics[width=3.0cm,angle=-90]{pp1.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp2.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp3.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp4.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp5.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp6.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp7.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp8.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp9.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp10.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp11.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp12.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp13.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp14.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp15.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp16.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp17.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp18.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp19.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp20.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp21.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp22.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp23.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp24.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp25.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp26.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp27.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp28.ps} \\\\ \\caption{Phased lightcurves for the 49 objects with periods in the order as listed in Table \\ref{periods}, part 1. Ids from \\citet{2007AJ....134.2340K} and adopted periods are indicated. The most robust periods (flag $\\ge 4$) are plotted with bold symbols. \\label{f6}} \\end{figure*} \\begin{figure*} \\includegraphics[width=3.0cm,angle=-90]{pp29.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp30.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp31.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp32.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp33.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp34.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp35.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp36.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp37.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp38.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp39.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp40.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp41.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp42.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp43.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp44.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp45.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp46.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp47.ps} \\hfill \\includegraphics[width=3.0cm,angle=-90]{pp48.ps} \\\\ \\includegraphics[width=3.0cm,angle=-90]{pp49.ps} \\\\ \\caption{Phased lightcurves for the 49 objects with periods in the order as listed in Table \\ref{periods}, part 2. Ids from \\citet{2007AJ....134.2340K} and adopted periods are indicated. The most robust periods (flag $\\ge 4$) are plotted with bold symbols. \\label{f7}} \\end{figure*}" }, "1101/1101.4635_arXiv.txt": { "abstract": "We study how the addition of on-board optical photometric bands to future space-based weak lensing instruments could affect the photometric redshift estimation of galaxies, and hence improve estimations of the dark energy parameters through weak lensing. Basing our study on the current proposed Euclid configuration and using a mock catalog of galaxy observations, various on-board options are tested and compared with the use of ground-based observations from the Large Synoptic Survey Telescope (LSST) and Pan-STARRS. Comparisons are made through the use of the dark energy Figure of Merit, which provides a quantifiable measure of the change in the quality of the scientific results that can be obtained in each scenario. Effects of systematic offsets between LSST and Euclid photometric calibration are also studied. We find that adding two ($U$ and $G$) or even one ($U$) on-board optical band-passes to the space-based infrared instrument greatly improves its photometric redshift performance, bringing it close to the level that would be achieved by combining observations from both space-based and ground-based surveys while freeing the space mission from reliance on external datasets. ", "introduction": " ", "conclusions": "" }, "1101/1101.2155_arXiv.txt": { "abstract": "Mildly mixed coupled models include massive $\\nu$'s and CDM--DE coupling. We present new tests of their likelihood {\\it vs.}~recent data including WMAP7, confirming it to exceed $\\Lambda$CDM, although at $\\sim 2$--$\\sigma$'s. We then show the impact on the physics of the dark components of $\\nu$--mass detection in $^3$H $\\beta$--decay or $0 \\nu \\beta \\beta$--decay experiments. \\vspace{1pc} ", "introduction": "Cosmological data (apart $^7$LI abundance) are nicely fitted by $\\Lambda$CDM, a model which however has severe {\\it fine tuning} and {\\it coincidence} problems. Here we therefore discuss an alternative easing these problems: that, symoultaneously, neutrinos ($\\nu$) have mass, and DE is a scalar field $\\phi$ self--interacting and interacting with Cold Dark Matter (CDM). To our knowledge, this is the only alternative whose likelihood, although marginally, exceeds $\\Lambda$CDM. An energy transfer from CDM to Dark Energy (DE) causes significant distorsions of $C_l$ and $P(k)$ spectra in respect to $\\Lambda$CDM, but allows DE to be a significant cosmic component since ever; distortions are also caused by $\\nu$ masses, in the range $M_\\nu = \\sum_i m_i \\sim 1$~eV. These two distorsions tend however to compensate and compensation allows to fit data better than $\\Lambda$CDM (Figure \\ref{cmb} shows this for CMB anisotropy spectrum). This yields models including a slight amount of Hot Dark Matter (typically $\\Omega_h \\sim 0.01$); they are then {\\it Mildly Mixed} and {\\it Coupled} (MMC) models. \\begin{figure}[h!] \\begin{center} \\includegraphics[height=6.cm,width=7.5truecm]{bin2.eps} \\end{center} \\vskip-1.truecm \\caption{$\\Lambda$CDM and MMC $C_l$ compared. In the upper (lower) plot $C_l$ are normalized to $\\Lambda$CDM (the best fitting SUGRA model including coupling and massive $\\nu$'s). $C_l$ obtained with either coupling or massive $\\nu$'s only are also shown. The error bars are a sampling of WMAP7 $C_l$ data.} \\label{cmb} \\end{figure} ", "conclusions": "" }, "1101/1101.2680_arXiv.txt": { "abstract": "We report on the aperiodic X-ray timing and color behavior of the accreting millisecond X-ray pulsar (AMXP) IGR J17511--3057, using all the pointed observations obtained with the {\\it{Rossi X-ray Timing Explorer}} Proportional Counter Array since the source's discovery on 2009 September 12. The source can be classified as an atoll source on the basis of the color and timing characteristics. It was in the hard state during the entire outburst. In the beginning and at the end of the outburst, the source exhibited what appear to be twin kHz quasi periodic oscillations (QPOs). The separation $\\Delta\\nu$ between the twin QPOs is $\\sim$ 120 Hz. Contrary to expectations for slow rotators, instead of being close to the 244.8 Hz spin frequency, it is close to half the spin frequency. However, identification of the QPOs is not certain as the source does not fit perfectly in the existing scheme of correlations of aperiodic variability frequencies seen in neutron star low mass X-ray binaries (NS LMXBs), nor can a single shift factor make it fit as has been reported for other AMXPs. These results indicate that IGR J17511-3057 is a unique source differing from other AMXPs and could play a key role in advancing our understanding of not only AMXPs, but also NS LMXBs in general. ", "introduction": "\\label{sec-intro} Low-mass X-ray binaries (LMXBs) are neutron star (NS) or black hole systems with low-mass ({\\it{M}} $\\leq$ 1$M_\\odot$) companion stars. Out of the nearly 200 LMXBs known so far \\citep{Liu2007}, 13 are accreting millisecond X-ray pulsars (AMXPs), i.e., they have shown coherent millisecond pulsations \\citep{Patruno2010}. The neutron stars in LMXBs are believed to be spun up by accretion to millisecond periods \\citep[see, e.g.,][]{Bhattacharya1991}, but why only some LMXBs appear as AMXPs is still an open question. These systems can be studied through the spectral (color-color diagram - CD) and timing (Fourier analysis) properties of their X-ray emission. Based on the paths traced on the CD and associated variability, NS LMXBs are classified either as {\\it{Z}} or atoll source \\citep{Hasinger1989}. Correlated with the position of the source in the CD, the Fourier power spectra of the X-ray flux variations exhibit different variability components. Apart from coherent pulsations, the power spectra also exhibit aperiodic phenomena: broad components (noise components) and narrow components (quasi periodic oscillations; QPOs) \\citep[see, e.g.,][for a review]{2006}. \\\\ \\\\ The coherent pulsations at the NS spin frequency are thought to be due to hot spots formed by magnetically channeled accreted matter \\citep[see, e.g.,][]{Pringle1972}. The origin of the aperiodic phenomena is poorly understood. They are presumably mostly associated with inhomogeneities in the matter moving in Keplerian orbits in the accretion disk, or in the boundary layer. The timescale for matter orbiting in the strong gravity region very near to the compact object is of the order of milliseconds (the dynamical timescale $\\tau = (r^3/GM)^{1/2} \\sim$ 0.1 ms at distance {\\it{r}} = 10 km from a compact object of mass {\\it{M}} = $1.4M_\\odot$). This strong gravity region can therefore be probed by the QPOs with millisecond timescales \\citep[see, e.g.,][]{2006}. The first millisecond phenomena were found with {\\it{Rossi X-ray Timing Explorer}} (RXTE): in Sco X-1, twin QPOs in the kHz range \\citep{vdk1996} and in 4U 1728-34 similar QPOs and also burst oscillations (oscillations near the spin frequency $\\nu_s$ that occur during type I X-ray bursts; see, e.g., \\citeauthor{Strohmayer2006} 2006 for a review). In many LMXBs, $\\Delta\\nu$, the difference between the twin kHz QPO frequencies, was at the frequency of burst oscillations $\\nu_{burst}$. A beat mechanism and $\\nu_{burst}$ = $\\nu_s$ were suggested by \\cite{Strohmayer1996b}. This led to the sonic-point beat-frequency model \\citep*{Miller1998}. \\\\ \\\\ Observations of variable $\\Delta\\nu$ in Sco X-1 \\citep{vdk1996} and later in other sources were inconsistent with the sonic-point beat-frequency model and led to a modified version \\citep{Lamb2001}. Also, the relativistic precession model \\citep{Stella1999} was proposed, in which $\\nu_s$ plays no direct role in the formation mechanism of QPOs. Observations of LMXBs like 4U1636--53 \\citep{1996a} which exhibited $\\Delta\\nu$ $\\sim$ $\\nu_{burst}$/2, clinched by the discovery of kHz QPOs in SAX J1808.8--3654 with $\\Delta\\nu$ $\\sim$ $\\nu_s$/2 by \\cite{Wijnands2003}, led to the proposal of new models, involving resonances. The relativistic resonance model \\citep{Kluzniak} and spin-resonance model \\citep{Lamb2003} were proposed which allowed $\\Delta\\nu$ = $\\nu_s$ and/or $\\Delta\\nu$ = $\\nu_s$/2. For historical accounts see \\citeauthor{2006} (2006, 2008) and \\citet{M'endez2007}. A model which can explain all the observations is still awaited. \\\\ \\\\ IGR J17511--3057 was discovered on 2009 September 12 during galactic bulge monitoring by INTEGRAL \\citep{Baldovin2009}. It showed X-ray pulsations at 244.8 Hz during {\\it{RXTE}} Proportional Counter Array (PCA) pointed observations which established its nature as an AMP \\citep{Markwardt2009}. It also exhibited type I X-ray bursts and burst oscillations (\\citeauthor{Watts2009a} 2009a, \\citeauthor{Altamirano2010} 2010a). A minimum companion mass estimate (assuming the NS mass = 1.4$M_\\odot$ and orbital inclination = $90^\\circ$) is 0.13$M_\\odot$ \\citep{Markwardt2009} and the upper limit to the distance is 6.9 kpc \\citep{Altamirano2010}. In this paper, we discuss the aperiodic variability of the AMP IGR J17511--3057. \\\\ ", "conclusions": "\\label{sec-discussion} The behavior of IGR J17511--3057 in the CD and the power spectra at first sight appears similar to what has been observed in other atoll sources and AMXPs \\citep[see, e.g,][]{2006}. From the color diagrams and the shape of power spectra of groups 3--6, the source appears to be an atoll source in the EIS. The other AMXPs have also been classified as atoll sources (see \\citeauthor{Wijnands2006} et al. 2006, \\citeauthor{Watts2009b} 2009a, \\citeauthor{Linares2008} 2008, \\citeauthor{Kaaret2003} 2003, \\citeauthor{Reig2000} 2000). However, closer study of the power spectral components indicates that this source is peculiar and does not fit well in the scheme defined by other sources. As discussed in Section \\ref{sec-identification}, all the possible scenarios for fitting our source in this scheme have their own shortcomings. Therefore, none of the components can be identified with certainty.\\\\ \\\\ The scenario that appears to require the least number of additional assumptions is scenario 1, where the only peculiarity is that twin kHz QPOs appear in what is otherwise an ordinary EIS; this was reported once before, in 4U 1728--34 \\citep{Migliari2003}. We note that if the two high frequency components are twin kHz QPOs, the measured differences $\\Delta\\nu$ between the centroid frequencies of high frequency QPOs in groups 1, 2 and 7 are $112^{+19}_{-8}$ , $144.9^{+18.6}_{-18.2}$ and $104.2^{+17.2}_{-12.9}$ Hz, respectively, and $119^{+22}_{-15.6}$ and $150^{+18.1}_{-15.8}$ Hz in group 1a and group 2a, respectively. These values are inconsistent with being close to the spin frequency $\\nu_{s}$ of 244.8 Hz as would have been expected for this so-called slow rotator ($\\nu_{s}$ $<$ 400 Hz; \\citeauthor{Miller1998} 1998). Instead, they are all consistent with half the spin frequency, which otherwise has only been seen in fast rotators ($\\nu_{s}$ $>$ 400 Hz). This can be seen in Figure \\ref{fig:deltanuvsnu} which shows the plot of $\\Delta\\nu$/$\\nu_s$ as a function of $\\nu_s$ for AMXPs and other atoll sources (the spin frequency is inferred from burst oscillations for these systems). The step function shows the historical distinction between the slow and fast rotators. It has been suggested that $\\Delta\\nu$ and $\\nu_s$ are (nearly) independent (\\citeauthor{Yin2007} 2007, \\citeauthor{M'endez2007} 2007). The curved line represents a constant $\\Delta\\nu$ of 300 Hz. To make the AMXPs SAX J1808.4--3658 and XTE J1807--294 fit this curve, they would have to be shifted up by a factor of $\\sim$1.5 \\citep{M'endez2007}. The points of IGR J17511--3057 for groups 1, 2 and 7 would require a factor $\\sim$2.5 to fall on this curved line. However, groups 1 and 2 do not require this same factor to fit in the scheme of correlations seen in Figure \\ref{fig:oldnunu}, but rather a factor 2.05. We applied no shifts to any data in this paper as there is no single factor. Note that all four AMXPs in Figure \\ref{fig:deltanuvsnu} are consistent with either $\\Delta\\nu$ = $\\nu_s$ or $\\Delta\\nu$ = $\\nu_{s/2}$ without shifts. \\\\ \\\\ Our results favor models like the sonic-point and spin-resonance model \\citep{Lamb2003} and the relativistic resonance model \\citep{Kluzniak} which predict that either $\\Delta\\nu$ = $\\nu_s$ or $\\Delta\\nu$ = $\\nu_s$/2. The relativistic precession model \\citep{Stella1999} predicts that $\\Delta\\nu$ should decrease when $\\nu_u$ increases as well as decreases (see Figure 2.14. in \\citeauthor{2006} 2006). Our results do suggest a low $\\Delta\\nu$, however the value is almost a factor of two lower than the value expected from the model at the observed $\\nu_u$ of IGR J17511--3057. \\\\ In conclusion, IGR J17511--3057 is indeed a very peculiar and interesting source. If scenarios 1 and 2 apply, the properties of the source can be summarized as follows: a) It exhibits kHz QPOs while in the EIS, b) In spite of being a slow rotator, kHz QPO frequency separation is $\\Delta\\nu$ $\\sim$ $\\nu_s$/2 and c) It requires different shift factors to fall on the frequency correlations of LMXBs at different times. Clearly this source could play a very important role in testing the existing models for the origin of QPOs. More observations of this source with {\\it{RXTE}} or perhaps ASTROSAT \\citep{Agrawal2002}, when it goes into an outburst again, are necessary to understand the nature of the different components. If we could observe the source in different spectral states, the components exhibited and their frequency evolution would help in establishing their nature. Observations of the high frequency QPOs and their evolution as a function of time and spectral state are key to their reliable identification." }, "1101/1101.1295_arXiv.txt": { "abstract": "The mystery of dark energy suggests that there is new gravitational physics on long length scales. Yet light degrees of freedom in gravity are strictly limited by Solar System observations. We can resolve this apparent contradiction by adding a Galilean-invariant scalar field to gravity. Called Galileons, these scalars have strong self-interactions near overdensities, like the Solar System, that suppress their dynamical effect. These nonlinearities are weak on cosmological scales, permitting new physics to operate. In this Letter, we point out that a massive gravity inspired coupling of Galileons to stress energy gravity can have a surprising consequence: enhanced gravitational lensing. Because the enhancement appears at a fixed scaled location for a wide range of dark matter halo masses, stacked cluster analysis of weak lensing data should be able to detect or constrain this effect. ", "introduction": " ", "conclusions": "" }, "1101/1101.3290_arXiv.txt": { "abstract": "Luminous compact blue galaxies (LCBGs) are a diverse class of galaxies characterized by high luminosity, blue color, and high surface brightness that sit at the critical juncture of galaxies evolving from the blue to the red sequence. As part of our multi-wavelength survey of local LCBGs, we have been studying the \\HI\\ content of these galaxies using both single-dish telescopes and interferometers. Our goals are to determine if single-dish \\HI\\ observations represent a true measure of the dynamical mass of LCBGs and to look for signatures of recent interactions that may be triggering star formation in LCBGs. Our data show that while some LCBGs are undergoing interactions, many appear isolated. While all LCBGs contain \\HI\\ and show signatures of rotation, the population does not lie on the Tully-Fisher relation nor can it evolve onto it. Furthermore, the \\HI\\ maps of many LCBGs show signatures of dynamically hot components, suggesting that we are seeing the formation of a thick disk or spheroid in at least some LCBGs. There is good agreement between the \\HI\\ and H$\\alpha$ kinematics for LCBGs, and both are similar in appearance to the H$\\alpha$ kinematics of high redshift star-forming galaxies. Our combined data suggest that star formation in LCBGs is primarily quenched by virial heating, consistent with model predictions. ", "introduction": "When the universe was 4.6 Gyr old, the galaxy population was dominated by blue, star-forming galaxies. Up to 40\\% of these galaxies were luminous compact blue galaxies (LCBGs) which contribute significantly to the global star formation rate density at that time \\cite{guzman97}. Today, the population of galaxies is roughly evenly divided between a red and a blue population and the star formation rate density has dropped by an order of magnitude. Similarly, LCBGs are an order of magnitude less common \\cite{werk04} and contribute negligibly to the global star formation rate \\cite{guzman97}. LCBGs are a diverse class of galaxies characterized by their high luminosities (M$_B \\le$-18.5 mag), compact sizes (SBe(B)$\\le$21 mag arcsec$^{-2}$, equivalent to r$_{eff}\\le$4 kpc), and blue colors ($B-V \\le$0.6 mag); they have the highest star formation rate per unit mass for high mass galaxies \\cite{gildepaz00}. Typical stellar masses of LCBGs are $\\sim$5$\\times$10$^{10}$\\msun, \\cite{guzman03} placing them near the maximal stellar mass of the blue sequence \\cite{kauffmann03}. Above this mass limit, all galaxies are red so some process must quench the star formation in galaxies as they grow. There have been numerous theories as to what quenching mechanisms operate in galaxies on the blue sequence. These include the shock heating of gas to the virial temperature \\cite[(Cattaneo et al. 2006)]{cattaneo06}, or heating by starbursts by supernovae- or AGN-driven winds or some combination of multiple processes \\cite[(Hopkins et al. 2006, and references therein)]{hopkins06}. Since LCBGs reside at the high mass end of the blue sequence, they are poised to have their star formation quenched in the near future and, therefore, represent an ideal population to study viable quenching mechanisms that could also be responsible for the emergence of a red sequence in the past 8 Gyr. We are conducting a multi-wavelength survey, spanning the ultraviolet through the radio, of the rare, local LCBGs to constrain the viable mechanisms for quenching star formation in blue galaxies and the future evolutionary paths of LCBGs. Therefore, we have selected our LCBGs from the Sloan Digital Sky Survey (SDSS) within D$\\le$200 Mpc to have the same properties, listed above, as LCBGs at high redshift. This yields a total of 2359 LCBGs out of over 800,000 galaxies in the SDSS DR4. Of these, we have collected single-dish \\HI\\ observations of 163 LCBGs. The distribution of properties for all LCBGs are shown in Figure~\\ref{fig1}. \\begin{figure}[b] \\begin{center} \\includegraphics[width=0.4\\textwidth]{vsig_mb_bv.eps} \\includegraphics[width=0.4\\textwidth]{vsig_mb_sbe.eps} \\caption{Left: $B-V$ vs. M$_B$ for all LCBGs selected from the SDSS DR4 (grey dots) and all LCBGs with single-dish \\HI\\ data (circles). The filled circles are those LCBGs with companions within the beam, the open circles are isolated. The triangles represent those LCBGs with GMRT or VLA data with the size of the triangle inversely proportional to V$_{rot}$/$\\sigma$. Right: Same as the left but for SBe(B) vs. M$_B$.} \\label{fig1} \\end{center} \\end{figure} ", "conclusions": "The signature of ongoing spheroid formation in some LCBGs is consistent with the idea that star formation in these galaxies is being quenched via virial heating, but this is not a unique explanation. Figure~\\ref{fig1} shows that those LCBGs with the smallest values of $V_{rot}/\\sigma$ are the most compact, bluest, and highest luminosity systems. This could also indicate that quenching from heating due to the intense central starburst or its associated supernovae is a possibility. This is supported by the results of optical spectroscopy by \\cite{perez-gallego11} who found that while only 5\\% of LCBGs have an AGN, 27\\% have signatures of supernovae- driven winds. The remaining LCBGs could then be quenched via virial heating. In the future, we will be expanding our \\HI\\ mapping to study additional LCBGs with a wider range of properties and we will use multi-wavelength data to search for signatures of active quenching in LCBGs." }, "1101/1101.1540_arXiv.txt": { "abstract": "Sunward-flowing voids above post-coronal mass ejection (CME) flare arcades were first discovered using the soft X-ray telescope (SXT) aboard \\textit{Yohkoh} and have since been observed with \\textit{TRACE} (extreme ultra-violet (EUV)), \\textit{SOHO}/LASCO (white light), \\textit{SOHO}/SUMER (EUV spectra), and \\textit{Hinode}/XRT (soft X-rays (SXR)). Supra-arcade downflow (SAD) observations suggest that they are the cross-sections of thin flux tubes retracting from a reconnection site high in the corona. Supra-arcade downflowing loops (SADLs) have also been observed under similar circumstances and are theorized to be SADs viewed from a perpendicular angle. Previous studies have presented detailed SAD observations for a small number of flares. In this paper we present a substantial SADs and SADLs flare catalog. We have applied semi-automatic detection software to several of these events to detect and track individual downflows thereby providing statistically significant samples of parameters such as velocity, acceleration, area, magnetic flux, shrinkage energy, and reconnection rate. We discuss these measurements, how they were obtained, and potential impact on reconnection models. ", "introduction": "Introduction} Long duration flaring events are often associated with downflowing voids and/or loops in the supra-arcade region (see Figure~\\ref{sads_sadls_example} for example images) whose theoretical origin as newly reconnected flux tubes has been supported by observations (\\citeauthor{mckenzie-hudson_1999}~\\citeyear{mckenzie-hudson_1999}; \\citeauthor{mckenzie_2000}~\\citeyear{mckenzie_2000}; \\citeauthor{innes-mckenzie-wang_2003a}~\\citeyear{innes-mckenzie-wang_2003a}; \\citeauthor{asai_2004}~\\citeyear{asai_2004}; \\citeauthor{sheeley-warren-wang_2004}~\\citeyear{sheeley-warren-wang_2004}; \\citeauthor{khan-bain-fletcher_2007}~\\citeyear{khan-bain-fletcher_2007}; \\citeauthor{reeves-seaton-forbes_2008}~\\citeyear{reeves-seaton-forbes_2008}; \\citeauthor{mckenzie-savage_2009}~\\citeyear{mckenzie-savage_2009}; \\citeauthor{savage_2010}~\\citeyear{savage_2010}). \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=.55\\textwidth]{sads_sadls_example.pdf} \\caption{(a) Example image from the 2002 April 21 TRACE flare showing supra-arcade downflows (SADs) enclosed within the white box. (b) Example image from the 2003 November 4 flare with supra-arcade downflowing loops (SADLs) indicated by the arrows. The left panel of each set is the original image. The right panel has been enhanced for motion via run-differencing and scaled for contrast.} \\label{sads_sadls_example} \\end{center} \\end{figure} The downflowing voids, (a.k.a. supra-arcade downflows (SADs) -- Figure~\\ref{sads_sadls_example} (a)), differ in appearance from downflowing loops (a.k.a. supra-arcade downflowing loops (SADLs) -- Figure~\\ref{sads_sadls_example} (b)); however, the explanation for this can be derived simply from observational perspective. If the loops are viewed nearly edge-on as they retract through a bright current sheet, then SADs may represent the cross-sections of the SADLs (see Figure~\\ref{sads_sadls_diagram_eyes_ch4}). Since neither SADs nor SADLs can be observed 3-dimensionally by an independent imaging instrument, proving this hypothetical connection is not possible with a single image sequence. However, their general bulk properties, such as velocity, size, and magnetic flux, can be measured and should be comparable if this scenario is correct. Moreover, measuring these parameters for a large sample of SADs and SADLs yields constraints that are useful for development of numerical models/simulations of 3D magnetic reconnection in the coronae of active stars. \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=0.9\\textwidth]{sads_sadls_diagram_eyes.pdf} \\caption{(a) Cartoon depiction of supra-arcade downflows (SADs) resulting from 3-D patchy reconnection. Discrete flux tubes are created, which then individually shrink, dipolarizing to form the post-eruption arcade. (b) Cartoon depiction of supra-arcade downflowing loops (SADLs) also resulting from 3-D patchy reconnection. Note that the viewing angle, indicated by the eye position, is perpendicular to that of SADs observations.} \\label{sads_sadls_diagram_eyes_ch4} \\end{center} \\end{figure} \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=0.6\\textwidth]{cartoon_bright_dark.pdf} \\caption{Diagram depiction illustrating the possible reason for the lack of bright SADs despite observations of bright SADLs with a snapshot of a loop shrinking through a bright current sheet with a viewing angle \\textit{Top:} along the arcade axis and \\textit{Bottom:} perpendicular to the axis. (Different scales) The loop is bright on the left of both panels and dark on the right. Because the current sheet is thin and in a region of low signal, the bright loops are easier to observe along the arcade axis which is opposite for the dark loops.} \\label{cartoon_bright_dark} \\end{center} \\end{figure} While bright shrinking loops (SADLs) are often observed as well as dark ones when the background has been sufficiently illuminated, bright SADs are rarer. A possibility for the reason behind this lack of bright SADs observations is given in Figure~\\ref{cartoon_bright_dark}. In order to view SADs, the loops are viewed edge-on to their apex cross-sections as they travel through the bright current sheet which provides a bright background. Consequently, if the loops are nearly as bright as the current sheet, then they blend in with the background and are difficult to observe. Conversely, because the current sheet is thin, viewing it edge-on gives the appearance of a bright, thin line surrounded by a dark background against which the bright loops can be seen. The coronal background itself would need to be sufficiently illuminated in order to observe dark loops as they retract. This is often not the case, however, because during a flare, the footpoints are so bright that the exposure durations applied to the images are not long enough to result in any significant coronal background signal. A counter-example is the 2008 April 09 flare described in \\cite{savage_2010}. In this paper we provide a list of 62 flares observed by several instruments containing downflow signatures. We analyze flows from 35 of these flares and present comparative results of general bulk properties, including magnetic flux and shrinkage energy estimates, from SADs and SADLs in Section~\\ref{sadsiisec:analysis}. These comparisons provide compelling evidence linking SADs to SADLs and constraints on flare magnetic reconnection models. Possible trends in the data are speculated in Section~\\ref{sadsiisec:trends}. The effect of simple drag on the loops is investigated in Section~\\ref{sadsiisec:drag} as a possible reason for the slow downflow speeds. Some loop cooling observations are presented in Section~\\ref{sadsiisec:loop_cooling} relating the appearance of shrinking loops and the brightening of the arcade. Finally, in Section~\\ref{sadsiisec:discussion} we summarize our findings of quantities that are typical of the observed SADs and SADLs, and suggest goals for future models of 3D bursty reconnection. \\subsection{\\label{sadsiisec:observations}Observations} Because SADs and SADLs are located in regions of extremely low coronal emission near bright, dynamic sources, measurements for any one flow are naturally associated with a high degree of error. Therefore, in order to identify any trends in the data, analysis of several flows from each of many flares must necessarily be performed. Table~\\ref{sads_list_full} contains a list of flares which have been noted to display downflow signatures (i.e. observable SADs, SADLs, or swaying fan above the arcade as suggested by \\cite{khan-bain-fletcher_2007}). Several of these flares were selected from \\cite{khan-bain-fletcher_2007} (see Table~1 therein) and were supplemented from \\cite{mckenzie_2000} as well as by personal flare data investigation. Some flows from the \\cite{khan-bain-fletcher_2007} list were excluded from this study if the presence of flows was not confirmed by visual inspection. The majority of the flares are from SXT observations. Under the ``Filter\" heading in Table~\\ref{sads_list_full}, the ``Q-\", ``H-\" and the ``F-\" preceding the filter indicate whether the images examined are quarter-, half-, or full-resolution, respectively (where Q = 9.8, H = 4.9, \\& F = 2.5~arcsec/pix). The SXT resolution is about 2$-$10 times poorer than either TRACE (0.5~arcsec/pix) or XRT (1~arcsec/pix) (depending on SXT resolution); however, it was operational during two solar maxima (unlike XRT to date which has been operational primarily during an unusually quiet solar minimum) and observed hotter plasma than TRACE. Having the capacity to observe the hot plasma in the current sheet increases the height of flow observations. A few of the flares lack GOES assignments either because the soft X-ray (SXR) output was too low or the footpoints were too far beyond the limb to measure any significant signal. Note that the TRACE flares in Table~\\ref{sads_list_full} are extremely energetic (as indicated by their GOES X1.5, X4.8, \\& X28 classifications). In addition, the TRACE flares are observed with the 195~\\AA\\ filter which has temperature response peaks both in the .5 $-$ 2~MK and 11 $-$ 26~MK bandpasses. The high energies result in very high temperature plasma detectable above the underlying post-eruptive arcade. The increased intensity in the supra-arcade region, presumably within the current sheet \\citep{reeves_2010}, provides a bright background against which to observe the dark downflows. Not all of the flares listed are suitable for tracking flows due to various factors (e.g. cadence, flow visibility, flare position, image quality, etc.). The last column of Table~\\ref{sads_list_full} indicates whether analysis of a flare was performed using the semi-automated routines described in \\cite{mckenzie-savage_2009}~(Section~2 therein) or a supplementary manual-tracking routine. Flows for 35 out of the 62 flares were evaluated. Table~\\ref{sads_list_analyzed} includes a list of the analyzed flares from this study. Whether the flows were determined to be clearly shrinking loops (SADLs) is indicated in the table so that the SADs results can be compared with those of the SADLs. (Both SADs and SADLs are clearly observed in the 2002 April 21 TRACE event.) Also indicated is the position of the flare on the Sun. Flares beyond the limb are given a limb designation for the instrument field of view (FOV). Flares occurring on the disk (within $\\sim50^{\\circ}$ from disk center) yield unreliable trajectory information due to the inability to accurately measure heights above the surface; therefore, their results are treated as detections only and removed from the following statistical analysis. It should be noted that the number of flows being reported are those that were deemed to be the most reliable and complete although additional flows (\\textgreater~50) have been processed. Inevitably, only a portion of the flows could be tracked for most flares due to noise, image quality, cadence, etc. (particularly for SXT) -- hence the need to process many flares in order to build up a catalog of flow parameters. The high resolution (0.5~arcsec/pix), high energy (GOES X1.5+) selected TRACE flares yield by far the highest number of clearly defined, easily-trackable SADs and SADLs, but even in those flares, there is substantial untrackable downflowing motion whose shrinkage energy contribution cannot be included in the final total estimates. This untracked motion is primarily present during the impulsive phase of the flare. \\subsection{\\label{sadsiisec:uncertainties}Sources of Uncertainty} Several variables were measured for each flow including height, velocity, acceleration, area, magnetic flux, and shrinkage energy. A description of these measurements can be found in \\cite{mckenzie-savage_2009}~(Section~2 therein). Flows that were either too difficult for the automatic routine to follow or contained shrinking loops were tracked manually instead. Although conservatively determined by visually judging the cross-sectional diameters and extrapolating a circular area, areas assigned manually are typically smaller than those determined using the automatic threshold technique. All of the manually evaluated flows are assigned a single area per flare whereas the automatically determined areas vary in time. The manual trajectories are thus better determined though the flow sizes are not temporally flexible. This is especially true with SXT data because of the low spatial resolution. As noted in \\cite{mckenzie-savage_2009}~(Section~3 therein), degrading the resolution of TRACE images to that of SXT's half-resolution leads to flow areas comparable to that of SXT. The result is that several smaller flows become undetectable, some flows that are near one another spatially are combined, and several of the flow ``heads\" are merged with their trailing ``tails\" making them appear larger. The square root of the largest area extent is used as the error on the flow positions. An additional large source of error is the initial height location. This height is biased by instrument-dependent detection capabilities (e.g. dynamic range and FOV) and is limited by the low emission high above the arcade. The initial height detection limits the path length of the measured flow trajectory which is used in the shrinkage energy calculation. Of larger consequence, the initial height determines the initial magnetic field invoked from the PFSS model (\\citeauthor{schatten-wilcox-ness_1969}~\\citeyear{schatten-wilcox-ness_1969}; \\citeauthor{schrijver-derosa_2003}~\\citeyear{schrijver-derosa_2003}) to calculate the flux and shrinkage energy. The ``Cartwheel CME\" flare from 2008 April 9 (\\citeauthor{ko_2010}~\\citeyear{ko_2010}; \\citeauthor{landi_2010}~\\citeyear{landi_2010}; \\citeauthor{patsourakos-vourlidas_2010}~\\citeyear{patsourakos-vourlidas_2010}; \\citeauthor{savage_2010}~\\citeyear{savage_2010}) is a nice example of being able to observe near the actual flow initiation region due to long exposure durations enabled by limb-obscured footpoints. This is a rare example, however, because active regions are not often observed for long after crossing the western solar limb and it is difficult to anticipate flares prior to crossing the eastern limb. AIA, an EUV imager aboard the recently-launched Solar Dynamics Observatory (SDO), will improve the number of these necessary limb observations since it observes the full solar disk continuously. Also noted in \\cite{mckenzie-savage_2009}, the PFSS model itself is another source of uncertainty considering that the flows are associated with flaring active regions which are expected to be non-potential. It was estimated that the uncertainties from the model do not exceed about 30\\% (Priest, private communication). The larger contributions to the magnetic field uncertainty are therefore the height input into the model, the footpoints assigned to the flows, and the flare's position on the Sun. Limb flares provide the optimum viewing angle for flow detections; however, the magnetograms used to extrapolate the magnetic fields into the corona are more reliable on the disk. Additionally, modeling for the east limb is even less reliable because magnetograms prior to the flare are unavailable. This discrepancy is shown in Figure~\\ref{radial_field} as the footpoint used for an X-class east limb flare (2002 July 23) is circled in panel (a) while the footpoints for X-class west limb flares (2002 April 21 \\& 2003 November 4) are circled in the bottom two panels. The initial height of the first flow per flare was used as a basis for the magnetic field represented in the panels. There is some variance, however, in the initial magnetic field strength between flows because the field extrapolation depends on initial height. Note the strong magnetic field indicated by the magnetograms for the west limb flares compared to that of the east. Several of the flares in Table~\\ref{sads_list_analyzed} occurred on the east limb (coordinates are shown in Table~\\ref{sads_list_full}). Magnetic fields are still estimated for these flares. The effect of the underestimated magnetic field becomes apparent in Figure~\\ref{quart3}~(c). Velocity results may be biased by the inability to track flows that are so fast they only show up in one frame or so slow that they are unobservable in the difference images. Most flows appear to be moving well within these constraints. Flows with areas near the resolution of the instrument are also difficult to detect which contributes an additional bias to the statistics. The obvious instrumental effect on area measurements is shown in Figure~\\ref{quart2}~(a). \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=0.6\\textwidth]{radial_field.pdf} \\caption{Radial magnetic fields derived from PFSS magnetic field modeling for the active regions producing flares observed by TRACE on a) 2002 July 23, b) 2002 April 21, and c) 2003 November 4. The initial height of the first flow per flare was used to create the figures. The footpoint region is circled. The east limb event in (a) predicts much weaker fields than the other two west limb events.} \\label{radial_field} \\end{center} \\end{figure} \\clearpage ", "conclusions": "Discussion \\& Conclusions} The preceding sections have presented many downflow observations from several instruments. Measurements of any one flow have too many possible sources of uncertainty to contribute to our understanding of the reconnection process occurring during long duration solar flares. SXT observed the most flares during its lifetime because it was operational throughout two solar maxima and was the first instrument to observe SADs; however, the poor spatial resolution leaves much to be desired with respect to being able to observe many flows per flare and make reliable measurements. LASCO has observed a number of flows in the outer corona; however, they are not always associated with flaring events, and because of the imager's much higher observational regime above the solar surface and considerably lower resolution, comparisons with other SADs and SADLs observations are difficult to interpret. LASCO speeds and accelerations are the only parameters comparable to the other instruments. TRACE has the best resolution of all the instruments used for this study and has observed many hundreds of flares during its lifetime; however, flows are difficult to observe with TRACE unless the 195 \\AA\\ filter is being utilized for broader temperature coverage and the flare is atypically large (all three TRACE flares in this study were GOES X-class flares). These requirements are due to the need for the hot plasma above the flare arcade to be illuminated. Finally, XRT has high enough spatial resolution to observe the flows as well as the optimal temperature coverage; however, solar activity has been unusually low throughout most of its lifetime to date making the amount of available flare data small. For all of these reasons, combining flow observations from all instruments improves our understanding of the flows themselves and ultimately our understanding of the reconnection process. Also, comparing the measurements between the instruments allows us to determine if the appearance of flows is temperature or density dependent. Interpreting SADs as the cross-sections of retracting reconnected flux tubes also means that if they are viewed from an angle that is not near perpendicular to the arcade axis (i.e. the polarity inversion line), the downflows will instead appear as shrinking loops. These shrinking loops (SADLs) have indeed been clearly observed with all of the instruments under investigation. Therefore, comparing observations of SADs to those of SADLs can help to support or refute the hypothesis for SADs. Figures~\\ref{quart1} through \\ref{quart3} present a summary of the instrument and SAD/SADL comparisons. These figures show that the flow velocities and accelerations agree between the instruments quite well. Height measurements agree except for those measured with XRT due to the exceptional heights observed for the ``Cartwheel CME\" flare. Figure~\\ref{quart2} (a) shows that the area measurements are understandably resolution dependent, which indicates that we may not be able to observe the smallest loop sizes. The flux and energy measurements are area dependent and therefore instrument dependent. There is also a limb dependence with the magnetic measurements due to the use of modeling based on magnetograms. Even so, there is decent agreement between all of the instruments. (LASCO is only included with the velocity and acceleration comparisons as explained in Section~\\ref{sadsiisec:quartiles}.) Beyond the agreement between the SADs and SADLs measurements, the high-resolution TRACE observations clearly show both SADs and SADLs occurring during the same flare depending on the arcade viewing angle which curves within the active region. The SADs versus SADLs diagram shown in Figure~\\ref{sads_sadls_diagram_mag} (previously provided in \\cite{mckenzie-savage_2009} -- Figure~10 \\& \\cite{savage_2010} -- Figure~23) still remains applicable after this analysis and has been updated to include SADLs with magnetic estimates. The measured cross-sectional areas range from $\\sim$2--90~Mm$^{2}$, with at least 75\\% being smaller than 40~Mm$^{2}$ (Figure~\\ref{quart2}~(a)). The flows typically move at speeds on order of 10$^{2}$~km~s$^{-1}$ with accelerations that are near zero or slightly decelerating. The most complete flow paths show significant deceleration near the top of the arcade. There is a range of initial heights depending on the quality of the image set, but they are generally about 10$^{5}$~km above the solar surface with a path length of $\\sim$10$^{4}$~km. Each tube carries $\\sim$10$^{18}$~Mx of flux and releases on order of 10$^{27}$~ergs of energy as it retracts. A lower limit of 10$^{16}$~Mx~s$^{-1}$ can be put on the reconnection rate by considering the total flux released by the observed flows for 5 flares (Section~\\ref{sadsiisec:magnetic}). These observations and measurements support the conclusion that SADs are indeed post-reconnection loops relaxing to form the post-eruption arcade. Also, the lack of instrument dependency of the dark flow observations suggests that either the loops are filled with cold material or are depleted. The temperature coverage of the instruments used in this study goes up to about 100 MK; therefore, it is unlikely that the loops are filled with hotter material. Combining this with the SUMER analysis of the 2002 April 21 TRACE flare from \\cite{innes-mckenzie-wang_2003a} which showed the lack of continuum absorption or emission in the C II, Fe XII, and Fe XXI lines at the flow sites supports the hypothesis that the tubes are depleted. This is at least true for loops reconnecting following the flare impulsive phase, which may be due to the fact that, according to the standard model, subsequent loops reconnect higher in the corona where less plasma is available to fill the loops. Loops shrinking very early during the 2002 April 21 TRACE flare appear to be bright in the hot 195 \\AA\\ bandpass as noted in Section~\\ref{sadsiisec:loop_cooling}. Few bright flows are analyzed in this study because they are more difficult to observe as SADs due to the low contrast although bright SADLs have easily been observed (see Figure~\\ref{cartoon_bright_dark}). Density analysis has not been performed due to the small sample of bright, tracked flows and especially lack of spectral coverage. \\clearpage \\begin{figure}[!ht] \\begin{center} \\includegraphics[width=.7\\textwidth]{sads_sadls_diagram_labeled.pdf} \\caption{(a) Schematic diagram of supra-arcade downflows (SADs) resulting from 3-D patchy reconnection. Discrete flux tubes are created, which then individually shrink, dipolarizing to form the post-eruption arcade. The measured quantities shown are averages from the limb events listed in Table~\\ref{sads_list_analyzed} containing SADs. (b) Schematic diagram of supra-arcade downflowing loops (SADLs) also resulting from 3-D patchy reconnection. The measured quantities shown are averages from the limb events listed in Table~\\ref{sads_list_analyzed} containing SADLs.} \\label{sads_sadls_diagram_mag} \\end{center} \\end{figure} There are a few trends to note from Figure~\\ref{trend_limb_selected}. Panel (a) shows that SADs tend to have increased areas in regions of weaker magnetic field. This result may be due to the fact that the weaker magnetic fields are generally associated with higher coronal heights where the signal to noise is very low. The noise impedes precise area measurements. If the correlation is real though, it may indicate that a given reconnection episode is associated with a limited amount of flux transfer. There is also the possibility that the loop cross-sections shrink as they retract either physically or apparently (due to filling with heated plasma from chromospheric evaporation). In this situation, lower initial heights, where the magnetic field strength is larger, would be associated with smaller areas as well. Then in panel (e), the initial heights seem to increase with time which is an expected consequence from the CHSKP model. This trend, however, is likely the result of an observational bias due to background brightening as the flare progresses rather than evidence of an upwardly migrating X-point. The observational findings presented in this paper provide a more complete description of the SAD/SADL phenomenon than has previously been available. Assuming that SADs and SADLs are thin, post-reconnection loops based on this body of evidence, the measurements obtained through this analysis and summarized in Figure~\\ref{sads_sadls_diagram_mag} provide useful constraints for reconnection models. Area estimates can constrain the diffusion time per episode and reconnection rates can be derived to distinguish between fast and slow reconnection. Creation of outflowing flux tubes carrying on order of 10$^{18}$~Mx of flux, with net reconnection rates of at least 10$^{16}$~Mx~s$^{-1}$, should be an objective of realistic models of 3D reconnection. The lack of acceleration of the downflow speeds and their discrete nature tends to favor 3D patchy Petschek reconnection. Speeds almost an order of magnitude slower than traditionally assumed Alfv\\'{e}n speeds are an unexpected consequence of the flow measurements; therefore, analyzing the effect of some source of drag on the downflow trajectories using models (an effort begun by \\cite{linton-longcope_2006}) could provide valuable insight into this discrepancy. \\subsection*{Acknowledgements} This work was supported by NASA under contract NNM07AB07C with the Harvard-Smithsonian Astrophysical Observatory. The authors would like to thank Drs. D. Longcope, C. Kankelborg, J. Qiu, and A. Des Jardins for constructive conversations. Hinode is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in cooperation with ESA and NSC (Norway). Yohkoh data are provided courtesy of the NASA-supported Yohkoh Legacy Archive at Montana State University. \\clearpage" }, "1101/1101.1589_arXiv.txt": { "abstract": "The dense, cold gas of Infrared Dark Clouds (IRDCs) is thought to be representative of the initial conditions of massive star and star cluster formation. We analyze $\\thco$ $J=1-0$ line emission data from the Galactic Ring Survey of Jackson et al. for two filamentary IRDCs, comparing the mass surface densities derived from $\\thco$, $\\Sigco$, with those derived from mid-infrared small median filter extinction mapping, $\\Sigsmf$, by Butler \\& Tan. After accounting for molecular envelopes around the filaments, we find approximately linear relations between $\\Sigco$ and $\\Sigsmf$, i.e. an approximately constant ratio $\\Sigco/\\Sigsmf$ in the clouds. There is a variation of about a factor of two between the two clouds. We find evidence for a modest decrease of $\\Sigco/\\Sigsmf$ with increasing $\\Sigma$, which may be due to a systematic decrease in temperature, increase in importance of high $\\thco$ opacity cores, increase in dust opacity, or decrease in $\\thco$ abundance due to depletion in regions of higher column density. We perform ellipsoidal and filamentary virial analyses of the clouds, finding that the surface pressure terms are dynamically important and that globally the filaments may not yet have reached virial equilibrium. Some local regions along the filaments appear to be close to virial equilibrium, although still with dynamically important surface pressures, and these appear to be sites where star formation is most active. ", "introduction": "Massive, high column density Infrared Dark Clouds (IRDCs), typically identified as being opaque against the Galactic background at $\\sim 10 {\\rm \\rm{\\mu} m}$, are thought to contain the sites of future massive star and star cluster formation (e.g. Rathborne et al. 2006), since their densities ($n_{\\rm H}\\gsim 10^4\\:{\\rm cm^{-3}}$) and mass surface densities ($\\Sigma \\gsim 0.1\\:{\\rm g\\:cm^{-2}}$) are similar to regions known to be undergoing such formation activity (Tan 2007). Studies of molecular line emission from IRDCs can help determine their kinematics. In particular, we would like to know if they are gravitationally bound, if they are near virial equilibrium and if there is evidence for coherent gas motions that might indicate that IRDC formation involves converging atomic flows (Heitsch et al. 2008) or converging molecular flows from cloud collisions (Tan 2000). In this study we use $\\thco$ $J=1-0$ line emission data from the Galactic Ring Survey (GRS) (Jackson et al. 2006) for two filamentary IRDCs, clouds F ($l=34.437^\\circ$, $b=0.245^\\circ$, $d=3.7$~kpc) and H ($l=35.395^\\circ$, $b=-0.336^\\circ$, $d=2.9$~kpc) from the sample of 10 relatively nearby massive and dense IRDCs of Butler \\& Tan (2009, hereafter BT09), comparing $\\thco$-derived mass surface densities, $\\Sigco$, with small median filter (SMF) mid-infrared (MIR) ($\\rm 8\\mu m$) extinction mapping derived mass surface densities, $\\Sigsmf$, using the method of BT09 applied to the {\\it Spitzer} Infrared Array Camera (IRAC) band 4 images of the Galactic plane taken as part of the Galactic Legacy Mid-Plane Survey Extraordinaire (GLIMPSE) (Benjamin et al. 2003). We consider systematic errors in each of these methods, which is necessary before analyzing larger samples of clouds. We are also able to look for evidence of changing CO abundance with column density, e.g. due to possible depletion of CO at high densities. We then perform a virial analysis of the clouds to determine their dynamical state. There have been a number of other studies comparing $\\thco$ derived mass surface densities with those from other methods. For example, Goodman, Pineda \\& Schnee (2009) compared near infrared (NIR) dust extinction, far infrared (FIR) dust emission and $\\thco$ line emission in the Perseus giant molecular cloud (GMC), probing values of $\\Sigma$ up to $\\sim 0.02\\:{\\rm g\\:cm^{-2}}$ (i.e. up to $A_V\\simeq 8$~mag). Even after accounting for temperature and optical depth variations they concluded that $\\thco$ emission was a relatively unreliable tracer of mass surface density, perhaps due to threshold, depletion and opacity effects. Our study probes higher values of $\\Sigma$, from $\\sim 0.01$ to $\\sim 0.05\\:{\\rm g\\:cm^{-2}}$, and compares $\\thco$ emission with MIR extinction in order to investigate these processes. Battersby et al. (2010) used $\\thco$ emission, MIR extinction and FIR dust emission methods to measure $\\Sigma$ and mass of clumps in 8 IRDCs, one of which is IRDC F of our study. They did not present a specific comparison of $\\Sigco$ with other methods, although derived clump masses were in reasonable agreement. Their sample also included MIR-bright regions, associated with ultra-compact HII regions, for which the MIR extinction method cannot be applied. As we describe below, our approach differs in a number of ways, including by focusing on filamentary and mostly quiescent regions of IRDCs for which the MIR extinction method is most reliable and which are likely to be closer to the initial conditions of the massive star and star cluster formation process. We note that while IRDC F in particular does contain some regions of quite active star formation, including an ultra-compact \\ion{H}{2} region, here we have concentrated on its more quiescent portions. ", "conclusions": "We have compared measurements of mass surface density, $\\Sigma$, in two IRDC filaments based on $\\thco$ observations, $\\Sigco$, with those derived from MIR extinction mapping, $\\Sigsmf$, finding agreement at the factor of $\\sim 2$ level. A systematic decrease in $\\Sigco/\\Sigsmf$ with increasing $\\Sigma$ may be due to a systematic decrease in temperature, increase in the contribution of under-resolved high optical depth regions, increase in dust opacity, or decrease in $\\thco$ abundance due to depletion in regions of higher column density. Future studies that spatially resolve the temperature structure and MIR dust absorption properties can help to distinguish these possibilities. We have then used the kinematic information derived from $\\thco$ to study the dynamical state of the IRDCs. In particular we have evaluated the terms of the steady-state virial equation, including surface terms, under the assumption of ellipsoidal and filamentary geometries. In both cases we find evidence that the surface pressure terms are important and possibly dominant, which may indicate that the filaments, at least globally, have not yet reached virial equilibrium. These results would be consistent with models of compression of dense gas in colliding molecular flows, e.g. GMC collisions. Tan (2000) proposed that this mechanism may trigger the majority of star formation in shearing disk galaxies. The expected collision velocities are $\\sim 10\\:{\\rm km\\:s^{-1}}$. It is less clear whether colliding atomic flows, (e.g. Heitsch et al. 2008), which form the molecular gas after shock compression of atomic gas, would also produce such kinematic signatures: recall that we are inferring large surface pressures based on $\\thco$ emission from the envelopes around the IRDC filaments. Recent observations of extended, parsec-scale SiO emission, likely produced in shocks with velocities $\\gtrsim 12\\:{\\rm km\\:s^{-1}}$ in IRDC H by Jim\\'enez-Serra et al. (2010) may also support models of filament formation from converging flows. However, we caution that the observed extended SiO emission is very weak and may also be produced by multiple protostellar outflow sources forming in the IRDC (see Jim\\'enez-Serra et al. 2010 for further discussion). Our resolved filamentary virial analysis also indicates that the regions closest to virial equilibrium (strips F6, F7 and H2) are those which have initiated the most active star formation. This would be expected if models of slow, equilibrium star formation (Tan et al. 2006; Krumholz \\& Tan 2007) apply locally in these regions. In this case, these dense regions that have become gravitationally unstable, perhaps due to the action of external pressure and/or converging flows, then persist for more than one local dynamical time and so are able to reach approximate pressure and virial equilibrium with their surroundings. In this scenario, they are stabilized by the ram pressure generated by protostellar outflow feedback from the forming stars (Nakamura \\& Li 2007)." }, "1101/1101.4153_arXiv.txt": { "abstract": "*{LS~5039 is one of the few X-ray binaries detected at VHE, and potentially contains a young non-accreting pulsar. The outflow of accelerated particles emitting synchrotron emission can be directly mapped with high resolution radio observations. The morphology of the radio emission strongly depends on the properties of the compact object and on the orbital parameters of the binary system. We present VLBA observations of LS~5039 covering an orbital cycle, which show morphological and astrometric variability at mas scales. On the other hand, we discuss the possible association of LS~5039 with the supernova remnant SNR~G016.8$-$01.1.} \\abstract{LS~5039 is one of the few X-ray binaries detected at VHE, and potentially contains a young non-accreting pulsar. The outflow of accelerated particles emitting synchrotron emission can be directly mapped with high resolution radio observations. The morphology of the radio emission strongly depends on the properties of the compact object and on the orbital parameters of the binary system. We present VLBA observations of LS~5039 covering an orbital cycle, which show morphological and astrometric variability at mas scales. On the other hand, we discuss the possible association of LS~5039 with the supernova remnant SNR~G016.8$-$01.1.} ", "introduction": " ", "conclusions": "\\label{discussion} The emission at 5~GHz traces the short-lived electrons of the outflow up to 15~AU. The morphological changes along the orbit allow us to model the velocity, the energy, and the cooling time-scale of the flow of particles that originates the extended emission. The morphological and astrometric information constrains the inclination of the orbit ($i$), a key parameter of the system. The mass function of the system and $i$ yield the mass of the compact object. On the other hand, it should be possible to clearly trace the peak position along the orbit with observations at higher frequencies (where the phase calibrator is much more compact), yielding direct information on the absorption around the system. The VLBA images from 2007 at 5~GHz contain relevant information to test the models (Mold\\'on et al.\\ in prep.), but the lack of continous astrometric information following the peak of the emission at all orbital phases is a key point to better constrain the physical properties of the system. If the peak position is expected to be shifted between 1--2~mas, it is not possible to unambiguously distinguish the displacement of the peak from intrinsic variations of the extended emission. The flow velocity and cooling times strongly depend on this ambiguity, which can be disentangled with accurate astrometry, to be obtained in the future." }, "1101/1101.3545_arXiv.txt": { "abstract": "{The lower limit to the distribution of orbital periods $P$ for the current population of close-in exoplanets shows a distinctive discontinuity located at approximately one Jovian mass. Most smaller planets have orbital periods longer than $P \\sim 2.5$ days, while higher masses are found down to $P \\sim 1$ day.} {We analyze whether this observed mass-period distribution could be explained in terms of the combined effects of stellar tides and the interactions of planets with an inner cavity in the gaseous disk.} {We performed a series of hydrodynamical simulations of the evolution of single-planet systems in a gaseous disk with an inner cavity mimicking the inner boundary of the disk. The subsequent tidal evolution is analyzed assuming that orbital eccentricities are small and stellar tides are dominant.} {We find that most of the close-in exoplanet population is consistent with an inner edge of the protoplanetary disk being located at approximately $P \\gtrsim 2$ days for solar-type stars, in addition to orbital decay having been caused by stellar tides with a specific tidal parameter on the order of $Q'_* \\simeq 10^7$. The data is broadly consistent with planets more massive than one Jupiter mass undergoing type II migration, crossing the gap, and finally halting at the interior 2/1 mean-motion resonance with the disk edge. Smaller planets do not open a gap in the disk and remain trapped in the cavity edge. CoRoT-7b appears detached from the remaining exoplanet population, apparently requiring additional evolutionary effects to explain its current mass and semimajor axis.}{} ", "introduction": "Close-in planets (semimajor axis $a < 0.1$ AU) constitute a special subset of the exoplanetary population. Since it is unclear whether in-situ formation occurs, the current orbital and physical characteristics of these planets provide important constraints on their past evolution and formation process. Several mechanisms have been proposed to explain the pile-up of hot planets with a three~day orbital period, including a truncation of the gaseous disk by the star (Lin et al. 1996, Kuchner and Lecar 2002), planetary scattering combined with Kozai resonance and tidal circularization (Nagasawa et al. 2008), planetary evaporation (Davis and Wheatly 2009), and tidal interactions with the parent star (Jackson et al. 2009). In particular, Kuchner and Lecar (2002) suggested that a giant planet in circular orbit could halt its migration when its orbital period was half that of the inner edge of the disk. In this configuration, all the planet's circular Lindblad resonances would lie in the inner cavity (IC) and no further interchange of angular momentum would take place. Masset et al. (2006) performed a series of hydrodynamical simulations to follow the evolution of low-mass planets in disks including an IC. They found that all bodies migrated until reaching a point slightly exterior to the cavity edge, where they were effectively trapped in a stable configuration in almost circular orbits. Although this result appears different from that predicted by Kuchner and Lecar (2002), each is valid, as we shall see, for a different range of planetary masses. The first reference to a possible correlation between mass and orbital period for close-in planets was proposed by Mazeh et al. (2005) for only six transiting bodies. They found that both parameters seemed to follow a linear law, with more massive bodies being located at smaller semimajor axes. Southworth et al. (2007) and Davis and Wheatley (2009) extended the analysis to a larger transiting population, finding a similar result although with a much broader dispersion. They proposed that smaller planets closer to the star might have been lost because of evaporation, similar to that currently ongoing at least in HD209458b (Vidal-Madjar et al. 2003) and WASP-17 (Anderson et al. 2010). \\begin{figure} \\centerline{\\includegraphics*[width=20pc]{15774fg1.eps}} \\caption{Distribution of orbital periods and planetary masses for close-in exoplanets with orbital period $P < 12$ days. Data from http://exoplanets.eu. Black circles denote planets with both Doppler and transit data, while gray circles mark bodies without detected transits. Empty squares correspond to planets with retrograde orbits with respect to the stellar spin.} \\label{fig1} \\end{figure} Jackson et al. (2009) also analyzed the distribution of close-in planets, this time focusing on the correlation between the semimajor axis of the planet and the age of its star. They found that the lower limit of the semimajor axis was lower for younger stars, which implies that tidal effects could be responsible. Exoplanets with very short orbital periods in older stars would have had enough time to be tidally disrupted, thus they would only be presently observable in relatively young systems. In this paper, we revisit the mass-period distribution, taking advantage of the recent increase in the exoplanet population. Figure \\ref{fig1} shows the orbital periods~$P$, as a function of the mass~$m$, for the known population of exoplanets with $P \\le 12$~days (137 planets). Black circles correspond to cases for which both transits and Doppler data are available; bodies without detected transits are shown in gray. Exoplanets in apparent retrograde motion with respect to stellar rotation are identified by an empty square. These are WASP-8b (Queloz et al. 2010), WASP-17b (Anderson et al. 2010), WASP-33b (Collier Cameron 2010), Hat-P-7b (Winn et al. 2009), WASP-2b (Triaud et al. 2010), and WASP-15b and WASP-17b (Triaud et al. 2010). Although it may be argued that these bodies are not consistent with planetary migration (Triaud et al. 2010), they may also point towards primordial spin-orbit misalignment and not be related to subsequent orbital evolution of the planets (Lai et al. 2010). The distribution exhibits a noticeable ``step'', exoplanets larger than one Jupiter mass ($M_{\\rm Jup}$) appear to have a lower inner boundary (down to $\\sim 1$ day) while for $m < M_{\\rm Jup}$ the distribution seems restricted to larger values of $P$. The only exceptions are three bodies in the Super-Earth range, CoRoT-7b, GJ1214b, and GJ876d, which are all marked in the plot. Of these, the latter two planets belong to low-mass stars ($M_*=0.17$ and $M_*=0.32$ solar masses, respectively), thus constitute special cases. CoRoT-7b, however, belongs to a solar-type star (Rouan et al. 2009). This planet has a very short orbital period ($\\sim 0.85$ days) but also a very low mass ($m \\sim 0.015 M_{\\rm Jup}$), and does not seem to comply with the rest of the exoplanet distribution. In particular, Jackson et al. (2009) pointed out that CoRoT-7b could reach the Roche radius on timescales of $10^7 - 10^9$ years, depending on the value of the specific tidal parameter $Q'_*$. Regardless of these isolated cases, there seems to be a very clear discontinuity (or bump) in the mass-period distribution, located at approximately $m=M_{\\rm Jup}$. Moreover, for high masses the lower limit in orbital periods appears very close to a $2/1$ mean-motion resonance with the disk edge for small planetary bodies. This appears consistent with a scenario in which the planetary traps proposed by Masset et al. (2006) would dominate the low-mass region, while the mechanism of Kuchner and Lecar (2002) would be mainly responsible for the upper end of the mass spectrum. The main objective of this study is to test whether the combined action of planetary traps in the gaseous disk plus subsequent tidal interactions with the parent star could explain the observed distribution of close-in exoplanets. Since exoplanets in retrograde orbits should have exotic disk-planet and tidal evolutions, the study of these exosystems is beyond the scope of the present work, and we focus mainly on bodies believed to have orbital motion in the same direction as the stellar spin. Even in this case, we assume a zero inclination with respect to the stellar equator. In Section 2, we present a series of hydrodynamical simulations adopting different planetary masses and analyzing the relative halting distance from the central star. Not only are we interested in seeing whether such a hybrid and mass-selective process is possible, but also whether the boundary between both mechanisms is consistent with the observed distribution of close-in planets. Section 3 is devoted to the subsequent evolution of exoplanets under the stellar tide and their effects on any initial disk-driven distribution in the mass-period diagram. In Section 4, we analyze the case of the CoRoT-7 planetary system and present possible explanations of the present location of CoRoT-7b. Finally, conclusions close the paper in Section 5. \\begin{figure} \\centerline{\\includegraphics*[width=20pc]{15774fg2.eps}} \\caption{Snapshot of the surface density profile of one of our simulations. The inner cavity is centered around $r=1.8$. In the top frame, the low-density regions are shown in black, while the high-density regions are shown in white. The planet ($m = 0.1 M_{\\rm Jup}$) is located on the $x$-axis.} \\label{fig2} \\end{figure} ", "conclusions": "We have attempted to understand the dynamical origin and evolution of the mass-period distribution of close-in exoplanets. The present-day population shows a distinctive discontinuity located at approximately one Jovian mass. Smaller planets have orbital periods longer than $P \\sim 2.5$ days, while higher masses are found to have periods as short as $P \\sim 1$ day. We have found that the combined effects of tidal evolution and disk-planet interactions with an inner cavity (IC) in the gaseous disk can explain most of the observed characteristics. The current distribution appears to be compatible with an inner disk edge located approximately at distances of ${\\bar a} \\simeq 0.035$, which for solar-type stars corresponds roughly to orbital periods of $P \\simeq 2.5$ days. This value is consistent with the inner gas radii for T Tauri stars as estimated from CO spectroscopy (Najita et al. 2007, Carr 2007). Planets below a certain critical mass $m_c \\sim M_{\\rm Jup}$ are trapped just outside the IC as found by Masset et al. (2006). The location of the stationary solution with respect to the IC is practically mass-independent. In contrast, bodies with $m > m_c$ enter a regime characterized by a Type II migration that causes significant perturbations to the density profile of the disk; consequently the IC edge cannot generate a significant positive corotational torque and does not stop the orbital decay. As predicted by Kuchner and Lecar (2002), migration only brakes inside the inner edge at a $2/1$ mean-motion resonance with the cavity edge. For reasonable values of the disk viscosity, we expect a gap opening to occur when the height of the disk is approximately equal to the Hill radius of the planet. For a scale height equal to $H/r \\simeq 0.05$, this implies a minimum mass of $m \\simeq 0.4 M_{\\rm Jup}$, a value similar to our critical mass $m_c$. Adopting a value of one Jovian mass for the critical mass leads to a slightly larger value of $H/r \\simeq 0.07$. However, given the uncertainties involved in both the gap opening criteria and its dependence on the viscosity, the values may be considered to be virtually equivalent. Thus, the location of the bump in the observed distribution of close-in planets appears to be consistent with a mass threshold for gap opening (Crida et al. 2006 and references therein). Although the present data are sparse and plagued by the additional effects described in this work, we may expect that the aforementioned results will be confirmed by further detections. In contrast, when the statistics become sufficiently robust, the location of this bump may be used to place constraints on the physical properties of the inner disk, in terms of both temperature and effective viscosity, as these quantities feature in the gap opening criterion. Since the edge of the IC created in our simulations is placed at an arbitrary distance from the star, we have a certain degree of freedom when fitting the numerical mass-period distribution to the real planets (as seen in Figure \\ref{fig4}). We have adopted a vertical displacement that minimizes the number of exoplanet inside the cavity edge, but this is not the only option. It may be argued that it would be better to fit the synthetic curve with the location of sub-giants (i.e. $m \\sim 0.5 M_{\\rm Jup}$) in the $m$-$P$ diagram. We note that these bodies exhibit a smaller dispersion in orbital period than observed for any other mass range, and can be seen as a compact group in both plots of Figure \\ref{fig4}. However, this displacement would lead to our accepting a larger number of small planets inside the cavity edge, bodies whose subsequent tidal evolution should have been very small. Whatever the choice, the qualitative results are not affected. Moreover, the ratio of the stopping values of $P$ (for small and large bodies) is scale-independent, and found to be slightly larger than $2$. This is because although higher masses are stopped in a $2/1$ MMR with the IC, the small bodies are trapped outside the disk edge. Once again, the distribution of real planets seems to yield a similar ratio. The subsequent tidal evolution of the close-in planets is consistent with a stellar tidal parameter of $Q'_* = 10^7$, a value similar to that predicted by Schlaufman et al. (2010) from synthetic population models. Smaller parameters, leading to higher rates of orbital decay, do not lead to distributions similar to the observed population. This is also consistent with the analysis of Ogilvie \\& Lin (2007). A consequence of the tidal evolution is the removal of most of the original gas giants with short orbital periods and their substitution by exoplanets that were originally farther away. Thus, we expect that many of the primordial planets with $P < 2$ days and $m > 1 M_{\\rm Jup}$ might have been tidally disrupted and absorbed by their parent stars. Although this scenario is consistent with the properties of most of the exoplanetary population, it appears difficult to explain the present-day mass and orbit of CoRoT-7b. A possible explanation is to assume that the planet started its life as a gas giant whose gas envelope was completely evaporated (Valencia et al. 2010, Jackson et al. 2010). Last of all, in the present scenario we have neglected the role of the orbital inclinations; however, the same results should be expected as long as the inclinations are not very large. Three-dimensional studies of disk-planet interactions in the linear approximation (e.g. Tanaka and Ward 2004), as well as numerical simulations (e.g. Cresswell et al. 2007), show little effect on a finite inclination on the migration timescale. Similar results are also found for orbital evolution due to tidal effects (e.g. Ferraz-Mello et al. 2008, Barker and Ogilvie 2009). For planets in retrograde orbits results are, however, is more difficult to evaluate. It is unclear whether gas disks in retrograde motion (with respect to the stellar spin) would have inner cavities or how such a structure would interact with planets in its vicinity. Similarly, the tidal evolution of bodies in retrograde motion is poorly understood, thus it\tis impossible at present to ascertain how the results of this work could be extended to these systems." }, "1101/1101.4399.txt": { "abstract": "We present a survey for low-ionization metal absorption line systems towards 17 QSOs at redshifts $z_{\\rm em} = 5.8-6.4$. Nine of our objects were observed at high resolution with either Keck/HIRES or Magellan/MIKE, and the remainder at moderate resolution with Keck/ESI. The survey spans $5.3 < z_{\\rm abs} < 6.4$ and has a pathlength interval $\\Delta X = 39.5$, or $\\Delta z = 8.0$. In total we detect ten systems, five of which are new discoveries. The line-of-sight number density, $\\ell(X) = 0.25^{+0.21}_{-0.13}$ (95\\% confidence), is consistent with the combined number density at $z \\sim 3$ of DLAs and sub-DLAs, which comprise the main population of low-ionization systems at lower redshifts. This apparent lack of evolution may occur because low ionization systems are hosted by lower-mass halos at higher redshifts, or because the mean cross section of low-ionization gas at a given halo mass increases with redshift due to the higher densities and lower ionizing background. The roughly constant number density notably contrasts with the sharp decline at $z > 5.3$ in the number density of highly-ionized systems traced by \\civ. The low-ionization systems at $z \\sim 6$ span a similar range of velocity widths as lower-redshift sub-DLAs but have significantly weaker lines at a given width. This implies that the mass-metallicity relation of the host galaxies evolves towards lower metallicities at higher redshifts. These systems lack strong \\siiv\\ and \\civ, which are common among lower-redshift DLAs and sub-DLAs. This is consistent, however, with a similar decrease in the metallicity of the low- and high-ionization phases, and does not necessarily indicate a lack of nearby, highly-ionized gas. The high number density of low-ionization systems at $z \\sim 6$ suggests that we may be detecting galaxies below the current limits of $i$-dropout and \\lya\\ emission galaxy surveys. These systems may therefore be the first direct probes of the `typical' galaxies responsible for hydrogen reionization. ", "introduction": "\\label{sec:intro} The universe at redshift six remains one of the most challenging observational frontiers. The high opacity of the \\lya\\ forest inhibits detailed measurements of the intergalactic medium (IGM), which hosts the vast majority of baryons in the early universe. At the same time, extreme luminosity distances make it difficult to assemble large samples of galaxies, or to study individual objects in detail. Despite these challenges, however, significant observational progress is being made at $z \\sim 6$. Statistics of the scarce transmitted flux in the \\lya\\ forest are being combined with insights from quasar proximity zones to learn about the very high-redshift IGM \\citep[e.g.,][]{fan2006b,lidz2006,becker2007,gallerani2008,maselli2009,bolton2010,calverley2010,carilli2010,mesinger2010}. Deep galaxy surveys from the ground and from space are also beginning to yield considerable information on the properties of galaxies at $z \\sim 6$ and beyond \\citep[e.g.,][]{bouwens2007,bouwens2010,stanway2007,martin2008,ouchi2008,ouchi2009,bunker2010,hu2010,labbe2010,mclure2010,oesch2010,stark2010}. A more complete understanding of the universe at such early times, however, will require insight from a wide variety of observations. Metal absorption lines offer a unique probe into the high-redshift universe. Transitions with rest wavelengths longer than \\hi\\ \\lya\\ will remain visible in QSO and GRB spectra even when the forest becomes saturated. While absorption lines studies at $z \\sim 6$ are still in their early days, metal lines are expected to provide insight into past and ongoing star formation, the action of galactic winds, as well as conditions in the interstellar media of galaxies, just as they do at lower redshifts \\citep[e.g.,][]{songaila2001,simcoe2002,schaye2003,adelberger2005,wolfe2005,scannapieco2006,fox2007a}. At high redshifts, the role of metal absorption lines in studying galaxy evolution becomes increasingly significant as galaxies become more difficult to study in emission. Metal lines play another unique role at $z \\gtrsim 6$ as potential probes of hydrogen reionization. In a scenario where galaxies provide most of the ionizing photons, dense regions of the IGM may become chemically enriched early on, yet remain significantly neutral until the end of reionization due to the short recombination times at high densities. High-density peaks are commonly predicted to host the last pockets of neutral gas following the growth and overlap of large \\hii\\ regions \\citep{me2000, ciardi2003, furoh2005, gnedinfan2006}. If these regions are metal-enriched, then \u00d2forests\u00d3 of low-ionization absorption lines such as \\oi, \\siii, and \\cii\\ may be observable during this last phase of reionization \\citep{oh2002, furloeb2003}. We have conducted a two-part survey for metals out to $z \\sim 6$. In the first paper of this series \\citep[][hereafter Paper I]{becker2009} we used high-resolution, near-infrared spectra to search for \\civ\\ over $z = 5.3-6.0$. \\civ\\ is a widely-used tracer of highly-ionized metals in the IGM at $z < 5$ \\citep[e.g.,][]{songaila2001}, and initial studies had indicated that \\civ\\ may remain abundant out to $z \\sim 6$ \\citep{simcoe2006b, rw2006}. In contrast, the longer path length and greater sensitivity of our data allowed us to demonstrate that the number density of \\civ\\ systems declines by at least a factor of four from $z \\sim 4.5$ to $z > 5.3$. This downturn was also seen by \\citet{rw2009} using a larger sample of low-resolution spectra. Since \\civ\\ is expected to be an effective tracer of metals in the IGM at $z \\sim 5-6$ \\citep{oppenheimer2006, oppenheimer2009}, the decline suggests that there are fewer metals in the IGM towards higher redshifts, consistent with gradual enrichment of the IGM by star-forming galaxies. The apparent rapid evolution of \\civ\\ near $z \\sim 5$ may also point to changes in the ionization state of the metal-enriched gas, perhaps due to a hardening of the UV background associated with the beginning of \\heii\\ reionization \\citep{madauhaardt2009}. In this paper we present a complementary search for low-ionization metal absorbers over $5.3 < z < 6.4$. This is an extension of an earlier work \\citep{becker2006}, in which we used high-resolutions spectra of nine QSOs spanning $4.9 < z_{\\rm em} < 6.4$, including five at $z_{\\rm em} > 5.8$, to search for `\\oi\\ absorbers,' so named due to the presence of \\oi~$\\lambda1302$. Although we identified a possible excess of these systems at $z \\gtrsim 6$, the significance of that result remained unclear due to the relatively short survey pathlength, and the fact that the $z \\sim 6$ systems occurred almost entirely along a single sightline. Here we increase our sample of $z_{\\rm em} \\sim 6$ QSOs to 17. The expanded survey offers us the opportunity to gather a more representative sample of low-ionization systems, providing insight into early galaxies and possibly the tail end of hydrogen reionization. The remainder of the paper is organized as follows. The data are described in Section~\\ref{sec:data}. We present our detections and the measurements for individual systems in Section~\\ref{sec:systems}. In Section~\\ref{sec:sample} we analyze a number of ensemble properties, including the number density, ionic mass densities, relative abundance patterns, lack of strong high-ionization lines, and velocity width distribution. We also present evidence that the mass-metallicity relation of the host galaxies of low-ionization absorbers evolves out to $z \\sim 6$. Possible causes for the apparent lack of evolution in the number density of low-ionization absorbers are discussed in Section~\\ref{sec:discussion}, where we also consider the likely hosts of these systems. Finally, we summarize our conclusions in Section~\\ref{sec:summary}. Throughout this paper we assume $(\\Omega_{\\rm m}, \\Omega_{\\Lambda}, h) = (0.274, 0.726, 0.705)$ \\citep{komatsu2009}. ", "conclusions": "\\label{sec:discussion} We now consider whether the number density of low-ionization systems is expected to remain roughly constant with redshift, as observed. Previous works have shown that the number density (per unit X) of both DLAs and sub-DLAs shows little evolution over $3 \\le z \\le 5$ \\citep{peroux2005,prochaska2005,prochaska2009,noterdaeme2009,guimaraes2009}. Our observations suggests that the number density remains roughly constant out to $z \\sim 6$, but why should this happen? The number density of a population of absorbers will depend on their comoving number density, $\\phi_{\\rm c}$, and their mean physical absorption cross section, $\\langle \\sigma_{\\rm abs} \\rangle$, as \\begin{equation} \\ell(X) = \\frac{c}{H_0} \\, \\phi_{\\rm c} \\, \\langle \\sigma_{\\rm abs} \\rangle \\, . \\label{eq:number_density} \\end{equation} Observational \\citep[e.g.,][]{cooke2006a,cooke2006b} and theoretical \\citep[e.g.,][]{gardner1997a,haehnelt1998,nagamine2004a} arguments generally suggest that DLAs are hosted by dark matter haloes of mass $M \\gtrsim 10^{8-9}$~\\Msun, and may be largely associated with halos in the range $10^9 < M/M_{\\odot} < 10^{11}$ \\citep{pontzen2008}. For a Sheth-Tormen halo mass function, $\\phi_{\\rm c}$ for $10^{10}$~\\Msun\\ halos, for example, will be lower at $z=6$ than at $z=3$ by a factor of $\\sim$4 \\citep{sheth2001,reed2007}. It is possible that the host halos of low-ionization systems are somewhat less massive at higher redshifts, which would increase their space density. The similarity of the velocity width distributions of the $z \\sim 6$ systems and low-redshift sub-DLAs suggests that the halo masses are similar, although the velocity width may only be moderately sensitive to the mass. For an NFW dark matter profile \\citep{navarro1997}, the virial velocity scales with the halo mass as $v_{\\rm vir} \\propto M_{\\rm vir}^{1/3}H(z)^{1/3}$, where $H(z)$ is the Hubble parameter at redshift $z$. If \\dv\\ scales with the virial velocity \\citep[e.g.,][]{haehnelt1998}, then a decrease in the typical halo mass by a factor of ten from $z=3$ to 6 would only produce a decline in the typical \\dv\\ of $\\sim$0.2~dex. Such a decline is consistent with the present data, but as noted in Section~\\ref{sec:vel_distrib}, the distribution of \\dv\\ values for the $z \\sim 6$ sample is formally consistent with that of lower-redshift sub-DLAs. A larger sample is needed to determine whether there is a genuine decrease in the mean velocity width. Multiple factors may increase the absorption cross-section with redshift. At a given mass, an NFW profile will have a higher density, particularly in the inner regions. The external UV background also decreases from $z \\sim 3$ to 6 \\citep{fan2006b,bolton2007b,wyithe2010,calverley2010}. Both of these factors should increase the radius out to which a halo of a given mass can become self-shielded. We explored simple models using NFW profiles to predict how the cross-section of low-ionization gas should evolve with redshift. Assuming that the gas traces the dark matter and is in photoionization equilibrium, it is straightforward to predict the radius at which the gas becomes self-shielded and to estimate the cross-section over which the projected \\hi\\ column density exceeds a given value. For this we assumed a gas temperature of 20,000~K and a metagalactic \\hi\\ ionization rate $\\Gamma = 10^{-11.9}~{\\rm s^{-1}}$ ($10^{-12.8}~{\\rm s^{-1}}$) at $z = 3$ (6). This simple model suggests that the cross section over which a halo of a given mass would be a DLA should scale roughly linearly with the halo mass, which agrees broadly with recent theoretical models \\citep{nagamine2004a,pontzen2008,tescari2009}. It also predicts that the DLA cross section for a given mass halo should increase by a factor of $\\sim$4 from $z = 3$ to 6. If DLAs are mainly hosted by dark matter halos with masses $M \\ge 10^9$~M$_\\sun$, then this increase in the cross section will offset the decline in the comoving spatial density to produce a roughly constant number density of DLAs with redshift. The same is true for sub-DLAs in this model. We stress that this is an extremely simple model, and it is unlikely to reproduce detailed observations such as the column density or velocity width distributions of DLAs. Nevertheless, it gives some assurance that the apparent lack of evolution in the number density of low-ionization absorbers can be plausibly explained by a decline in the spatial density of dark matter halos combined with an increase in the mean absorption cross section. More detailed analysis must be left for future work. The high number density of low-ionization systems makes it likely that they arise from galaxies below the detection limits of current galaxy surveys. Integrating the $z \\sim 6$ luminosity function of $i$-dropout galaxies from \\citet{bouwens2007} down to the limit of the Hubble Ultra Deep Field ($\\sim$0.04$L^{*}_{z=3}$) gives a space density $\\phi(M_{\\rm AB} \\le -18) = 5 \\times 10^{-3}$ per comoving Mpc$^3$. For $\\ell(X) \\simeq 0.25$, this implies a mean absorption cross-section $\\langle \\sigma_{\\rm abs} \\rangle \\simeq 1 \\times 10^4~{\\rm kpc}^{-2}$, or $\\langle R_{\\rm abs}\\rangle \\simeq 60$~kpc for a circular projected geometry. The \\lya-emitting galaxies (LAEs) at $z \\sim 6$ included in current surveys would need an even larger cross-section. \\citet{ouchi2008} found a density of LAEs at $z=5.7$ of $\\sim$$7 \\times 10^{-4}$ per Mpc$^{-3}$ down to a \\lya\\ luminosity of $2.5 \\times 10^{42}~{\\rm erg~s^{-1}}$. These LAEs would require $\\langle R_{\\rm abs}\\rangle \\simeq 160$~kpc to account for the observed number density of low-ionization absorption systems. In contrast, constraints from paired QSO lines of sight suggest a typical DLA size at $z = 1$-2 of $R_{\\rm DLA} \\simeq 5-10$~kpc \\citep{monier2009,cooke2010}, although in some cases they may be larger \\citep[e.g.,][]{briggs1989}. The sizes of sub-DLAs are less well constrained, but they may plausibly be a factor of two larger than DLAs (see, for example, the scaling relation in \\citet{monier2009}, or Figure~2 in \\citet{pontzen2008}). In that case, they would still need to be significantly larger at $z = 6$ to arise solely from galaxies similar to those in current surveys. It seems more plausible, therefore, that these low-ionization systems trace fainter, more numerous sources. This would be consistent with the results of \\citet{rauch2008}, who find a likely connection between DLAs and a population of ultra-faint LAEs at $z \\sim 3$. We note that these absorption systems are potentially our first direct probe of the `typical' galaxies that are responsible for hydrogen reionization. It is widely believed that in order for reionization to complete by $z = 6$, the majority of ionizing photons must come from galaxies below the limits of current imaging surveys \\citep[e.g.,][]{bouwens2007,bolton2007b,ouchi2009,bunker2010,mclure2010,oesch2010}. Metal absorption lines may therefore be the most efficient means of detecting these sources, as directly detecting them in emission must await facilities such as ALMA and JWST. This study already delivers insights into the properties of reionization galaxies, including the suggestion that they are likely to be more metal poor than their lower-redshift counterparts, at least in the gas phase. If their stars also have low metallicities, then they may be highly efficient in producing ionizing photons \\citep{schaerer2003,venkatesan2003}. As noted by others \\citep[e.g.,][]{ouchi2009}, this would potentially help to explain how reionization is completed by $z \\sim 6$. We add that, in addition to their role in reionization, these low-mass galaxies may contribute significantly to the chemical enrichment of the IGM \\citep[e.g.,][]{booth2010}. We have conducted a search for low-ionization metal absorption systems spanning $5.3 < z < 6.4$. Our survey includes high- and moderate-resolution spectra of 17 QSOs with emission redshifts $z_{\\rm em} = 5.8-6.4$. The total survey pathlength is $\\Delta X = 39.5$ ($\\Delta z = 8.0$, or $\\Delta l = 3.6$~comoving Gpc), of which roughly half is covered by high-resolution data. We searched for low-ionization systems by looking for coincidences in redshift between \\siii, \\cii, and \\oi. In total we detect ten systems, five of which were previously reported by \\citet{becker2006}. Each contains \\siii\\ and \\cii, and all but one contains \\oi. The majority are detected in our high-resolution data, which is consistent with the fact that these data are more sensitive to narrow absorption lines. The line-of-sight number density of absorption systems, uncorrected for completeness, is $\\ell(X) = 0.25^{+0.21}_{-0.13}$ (95\\% confidence). This is similar to the number density over $3 < z < 5$ of DLAs and sub-DLAs ($\\log{N({\\mbox \\hi})} \\ge 19.0$), which constitute the main population of low-ionization absorbers at those redshifts. The fact that the number density of low-ionization absorbers is roughly constant out to $z \\sim 6$ is in sharp contrast with the evolution of high-ionization systems traced by \\civ, which show a marked decline at $z > 5.3$ \\citep[Paper I;][]{rw2009}. At $z \\sim 6$, low-ionization systems with $\\log{N({\\mbox \\cii})} \\gtrsim 13$ are more common than high-ionization systems with $\\log{N({\\mbox \\civ})} \\gtrsim 13$, a reversal from lower redshifts. The roughly constant number density of low-ionization systems over $3 \\lesssim z \\lesssim 6$ may be explained if they are hosted by lower-mass dark matter halos at higher redshifts. Alternatively, if the systems at $z \\sim 6$ are hosted by halos with masses similar to those that host DLAs and sub-DLAs at lower redshifts, then the apparent lack of evolution may occur if dark matter halos of a given mass have a larger mean cross-section of low-ionization gas at higher redshifts due to the higher gas densities and weaker UV background. Although the \\hi\\ column densities cannot be measured for these systems, we are able to infer some of their properties by comparing the velocity widths and metal line strengths to samples at lower redshifts. For the seven systems with high resolution data, the velocity widths span a similar range as sub-DLAs at $2 \\lesssim z \\lesssim 4$, although there is some indication that the $z \\sim 6$ systems tend to be narrower. The lines in the $z \\sim 6$ systems are also weaker in the sense that, at a given velocity width, the inferred equivalent width of \\siii~$\\lambda 1526$ is lower than for $z \\lesssim 4$ sub-DLAs by a factor of two, and by a factor of six compared to $z \\lesssim 4$ DLAs. This implies that the mass-metallicity relation of the host galaxies evolves towards lower metallicities at higher redshifts, a trend that has also been noted over $2 \\lesssim z \\lesssim 4$ \\citep{ledoux2006}. Assuming the $z \\sim 6$ systems span a similar range in $\\log{N({\\mbox \\hi})}$ as the low-ionization systems at $z < 4$, the strength of the metal lines implies a decline in the gas-phase metallicity from $2 \\lesssim z \\lesssim 4$ to $z \\sim 6$ of at least $\\sim$0.4 dex. The mean relative abundances are $[{\\rm Si/O}] = -0.03 \\pm 0.03$ and $[{\\rm C/O}] = -0.18 \\pm 0.03$ (1$\\sigma$ uncertainties), assuming no dust depletion and no ionization corrections. These are consistent with the values found for metal-poor DLAs by \\citet{pettini2008} and for low-metallicity halo stars by \\citet{akerman2004}. As noted by \\citet{becker2006}, the relative abundances are broadly consistent with enrichment from Type II supernovae of low-metallicity massive stars. Our $z \\sim 6$ systems are also notable in that they lack strong high-ionization lines (\\siiv\\ and \\civ), which are ubiquitous among lower-redshift DLAs and sub-DLAs \\citep{fox2007a}. The absence of these lines is consistent with a similar fractional decline in the metallicity of the low- and high-ionization phases, however, and does not necessarily indicate that these systems lack halos of highly-ionized gas. Deeper spectra will help to determine whether the $z \\sim 6$ absorbers have significantly lower $N({\\mbox \\siiv})/N({\\mbox \\siii})$ and $N({\\mbox \\civ})/N({\\mbox \\cii})$ ratios than lower-redshift systems. The overall consistency between the properties of the $z \\sim 6$ systems and those of lower-redshift DLAs and sub-DLAs suggests that the $z \\sim 6$ absorbers arise from galaxy halos, rather than from remnants of neutral IGM at the tail end of hydrogen reionization. One notable aspect of our survey, however, is that all ten of our systems occur at $z > 5.75$, while roughly 40\\% of our pathlength is at lower redshifts. Expanded surveys over $4.5 < z < 5.7$ will help to determine whether the number density of low-ionization systems truly remains constant out to $z \\sim 6$, or whether there is a decline up to $z \\sim 5.7$ followed by an increase at higher redshifts. The latter scenario would potentially indicate a strong evolution in the UV background near $z \\sim 6$. Finally, the high-number density of low-ionization systems at $z \\sim 6$ suggests that we are probing galaxies below the detection limits of current $i$-dropout and \\lya-emission galaxy surveys. As such, these absorption systems are potentially the first observations of `typical' galaxies responsible for hydrogen reionization. The low metallicities we infer suggest that these galaxies may be highly efficient at producing ionizing radiation, a fact which would help to explain how the IGM becomes fully ionized by $z \\sim 6$. As more and higher-redshift QSOs are discovered, absorption lines will continue to provide a unique probe of the reionization era that will complement studies with ALMA, JWST, and other next-generation facilities." }, "1101/1101.5822_arXiv.txt": { "abstract": "We present results from continued {\\em Chandra} X-ray imaging and spectroscopy of a flux-limited sample of flat spectrum radio-emitting quasars with jet-like extended structure. X-rays are detected from 24 of the 39 jets observed so far. We compute the distribution of $\\alpha_{rx}$, the spectral index between the X-ray and radio bands, showing that it is broad, extending at least from $0.8$ to $1.2$. While there is a general trend that the radio brightest jets are detected most often, it is clear that predicting the X-ray flux from the radio knot flux densities is risky so a shallow X-ray survey is the most effective means for finding jets that are X-ray bright. We test the model in which the X-rays result from inverse Compton (IC) scattering of cosmic microwave background (CMB) photons by relativistic electrons in the jet moving with high bulk Lorentz factor nearly along the line of sight. Depending on how the jet magnetic fields vary with $z$, the observed X-ray to radio flux ratios do not follow the redshift dependence expected from the IC-CMB model. For a subset of our sample with known superluminal motion based on VLBI observations, we estimate the angle of the kpc-scale jet to the line of sight by considering the additional information in the bends observed between pc- and kpc-scale jets. These angles are sometimes much smaller than estimates based on the IC-CMB model with a Lorentz factor of 15, indicating that these jets may decelerate significantly from pc scales to kpc scales. ", "introduction": "Many fundamental physical properties of quasar jets remain uncertain, such as the nature of the energy-carrying particles, whether the particle energy densities are in equipartition with the local magnetic field energy densities, and how much entrainment there is. From the observation of superluminal motion with the VLBI technique, it is generally agreed that the pc-scale jets of high power quasars are highly relativistic, with bulk Lorentz factors ($\\Gamma$) of 10-30. However, it is not certain whether most jets at the kpc-scale also have high Lorentz factors in bulk motion and whether the jets are oriented close to our line of sight, as inferred for \\pks\\ \\citep{celotti,tavecchio} because of its X-ray bright knots \\citep{schwartz}. The model posited by Celotti et al.\\ and Tavecchio et al.\\ involved inverse Compton scattering of photons from the Cosmic Microwave Background (IC-CMB), in contrast to earlier synchrotron and synchrotron self-Compton models. For a review of relativistic jet physics and the role of X-ray observations, see \\citet{2009A&ARv..17....1W} and references therein. It is now becoming evident that the simplest, single-zone IC-CMB model is inadequate in many cases \\cite[e.g.][]{ks05,hardcastle06,2006ApJ...648..900J,aneta2}. One concern with this model is that the lifetimes of the electrons responsible for the X-ray emission are orders of magnitude longer than those producing the radio emission so the observed correspondence of radio and X-ray structures would not be expected \\citep{tavecchio03,schwartz06}. Extra synchrotron components are proposed by others \\citep{2006ApJ...648..900J,hardcastle06}. In some cases it gets difficult to generate an adequate physical model \\citep{aneta2}. Currently, the field is in a fruitful phase of mutually driven theoretical and observational advances. Solutions seem to be as varied as the sources themselves, bolstering the need for more detailed case studies. This need provides the primary motivation for our X-ray imaging survey with {\\em Chandra} \\citep[][hereafter, Paper I]{marshall05}. Our survey is similar to that undertaken by \\citet{sambruna02,sambruna04} but the sample is somewhat larger and the exposures correspondingly shorter. This being a shallow survey, we leave detailed modeling of individual sources to later, follow-up analyses of deeper observations This paper is a continuation of Paper I and presents observations of another 19 quasars from the original sample of 56. We describe the sample properties in section 2. In section 3, we describe the {\\em Chandra} observations and compare the X-ray maps to newly obtained radio images. In section 4, we examine the sample properties in the context of beaming emission models and test the IC-CMB model in a limited context, as examined previously by \\citet{2004ApJ...600L..23C} and \\citet{ks05}. We use a cosmology in which $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_\\Lambda = 0.7$. ", "conclusions": "A hypothesis that bears testing with these data is that the X-ray emission results from IC-CMB photons off relativistic electrons and that the bulk motion of the jet is highly relativistic and aligned close to the line of sight. We have several lines of evidence that suggest that the jets in our sample are consistent with this interpretation. \\subsection{Detection Statistics} We detected 12 X-ray emitting jets among the 19 targets observed, half of which were previously reported. Of these detections, 9 were in the A subsample of 10 sources, while only 3 were in just the B subsample: 0234$+$285, 2007$+$777, and 2123$-$463. If detections were equally likely in both B and A samples, then the {\\it a priori} probability that there would be $<4$ B detections would be 7.3\\%, so the hypothesis that the morphology selection is just as good as a flux selection is marginally acceptable. Of the aggregate of 39 sources from Paper I and this paper, 22 were in the A subsample. Jets were detected in 16 of the 22 images, for a 73\\% detection rate. This detection rate is similar to that obtained by \\citet{sambruna04} and \\citet{marshall05}. The jet detection rate for the B-only subsample is not as high. Of all the B-only quasars, 7 of 17 jets are detected (41\\%). These rates could be biased, however, because those targets observed in other programs were generally the brightest A targets. The typical X-ray flux densities of detected jets are greater than 2 nJy. Flux densities in the radio band were generally lower in B targets than in A targets while the X-ray flux limits are all about the same; consequently, the lower limits on $\\alpha_{rx}$ are higher for the A targets (see Fig.~\\ref{fig:alpharx-z}). However, Fig.~\\ref{fig:arxdistribution} shows that the distribution of $\\alpha_{rx}$ is slightly shifted toward lower values of $\\alpha_{rx}$ for B targets compared to A targets, indicating slightly larger X-ray flux densities relative to the jets' radio flux densities as a group. Thus, it appears that morphological selection may yield jets brighter in the X-ray band. However, the distribution differences are not statistically significant, due to small number of detected sources in the B subsample. Furthermore, due to the systematically higher redshifts of the B subsample, jet detection rate differences between the two subsamples may result from redshift dependences. \\subsection{Modeling the X-ray Emission} \\subsubsection{Distribution of $\\alpha_{rx}$ and Redshift Dependence} As in Paper I, there are bright X-ray jets even in sources with weak extended radio flux, confirming that the ratio of the X-ray to radio flux densities has a wide range (see Fig.~\\ref{fig:arxdistribution}). The $\\pm 1\\sigma$ width of the $\\alpha_{rx}$ distribution is about 0.15 -- a factor of 15 in $R$, the ratio of the X-ray and radio flux densities (as extrapolated to a common frequency, see appendix~\\ref{app:d95}). The jets' radio flux densities extend over a factor of almost 100 for the {\\em detected} jets in our sample. While there is a general trend that the brightest jets are detected most often, it is clear that predicting the X-ray flux from the radio knot flux densities is risky, so a shallow survey is practically the only efficient means for finding jets that are X-ray bright. We note that the two jets detected out of four exposures longer than 10 ks would have been detected with just the first 10 ks. Whereas detailed individual analyses of the brighter quasar jets can test physical models \\citep[e.g.,][]{schwartz06,perlman08}, we explore here how even relatively short exposures can prove useful for statistical tests of the model in which the X-rays result from inverse Compton (IC) scattering of cosmic microwave background (CMB) photons by relativistic electrons in a jet moving with high bulk Lorentz factor nearly along the line of sight (the IC-CMB model). Particular support for this model arises in individual cases where the optical flux lies below the radio to X-ray interpolation, indicating that synchrotron radiation from a single population of relativistic electrons cannot fit the spectral energy distribution. We note that our objective is comparable to that examined by \\citet{2004ApJ...600L..23C} and \\citet{ks05} but with a larger, more homogeneous sample of FR II quasars containing a much larger fraction of sources with $z > 1$. We limit our analysis to the 34 quasars in our sample with known redshifts greater than 0.1 so as to avoid the slight contamination by flat spectrum, core-dominated FR I radio galaxies. Following \\citet[][HK02]{hk02} and Paper I (see also Appendix~\\ref{app:d95}) in the context of the IC-CMB model, $R$ can be related to the equipartition magnetic field in the absence of beaming, $B_1$, derived from the radio flux and emitting volume, and beaming parameters $\\Gamma = (1-\\beta^2)^{-1/2}$ and $\\mu = \\cos \\theta$, where $\\theta$ is the angle to the line of sight, as \\begin{equation} \\label{eq:r} R = A \\bigg( \\frac{b}{B_1} \\frac{ (1 - \\beta) (1+ \\mu) }{(1-\\mu \\beta)^{2}} \\bigg)^{(1+\\alpha)} (1+z)^{3+\\alpha} \\end{equation} \\noindent where $A = 6.9 \\times 10^{-21}$ and $b = 38080$~G are constants and $B_1$ has units of G. The spectral index, $\\alpha$, defined by $S_{\\nu} \\propto {\\nu}^{-\\alpha}$, is assumed to be 0.8 for both the X-ray and radio bands. Equation~\\ref{eq:r} can be solved to give $\\mu$ for an assumed value of $\\beta$ (Paper I) or for $\\beta$ for a given value of $\\mu$ \\citep{2006xru..conf..643M, mc07}. In Appendix~\\ref{app:d95}, we show that the HK02 approach is equivalent to the inverse Compton model developed by \\cite{dermer95} which was later written in a form independent of the system of units by \\citet{2009A&ARv..17....1W}. The quantity $R$ depends on $\\alpha_{rx}$ via the relationship $R = (\\nu_{\\rm x}/\\nu_{\\rm r})^{\\alpha-\\alpha_{rx}}$, so $\\alpha_{rx}$ depends on quasar redshift in the IC-CMB model. No significant correlation of $\\alpha_{rx}$ with $z$ is apparent in Fig.~\\ref{fig:arxdistribution}. However, with such a broad distribution of $\\alpha_{rx}$ it would be difficult to discern such a trend. We tested the possibility that $\\alpha_{rx}$ depends on $z$ by splitting the sample into two redshift ranges. For $0.55 < z < 0.95$, the average $\\alpha_{rx}$ is 1.001 $\\pm$ 0.020, compared to a value of 0.954 $\\pm$ 0.019 for $0.95 < z < 2$. The difference is insignificant. A more sensitive test is to explore the dependence of $R$ upon $z$. We use the method developed by \\cite{marshall92} to fit $R$ to the form $(1+z)^a$; details are given in Appendix~\\ref{app:regress}. However, $B_1$ is calculated from observations and depends on redshift according to model assumptions (as discussed below). Generally, we expect $B_1 \\propto f(z)$, giving \\begin{equation} \\label{eq:rz} R \\propto (1+z)^{3+\\alpha} [f(z)]^{-(1+\\alpha)} \\end{equation} \\noindent In the simple case where the distribution of intrinsic magnetic fields is independent of redshift, then we may set $f(z) = 1$. The log likelihood for this case is shown in Fig.~\\ref{fig:arxtest}, for which we find $a = 0.7 \\pm 1.6$ at 90\\% confidence ($\\Delta \\chi^2 = 2.71$). The likelihood ratio test then rejects $a > 3.5$ at $>99$\\% confidence, whereas we expect $a = 3.8$ for the IC-CMB model. It is common to estimate the magnetic field in individual sources based on observations and assume minimum energy. We note that a simple dependence of $B_1$ calculated this way (Tables~\\ref{tab:knotresults} and \\ref{tab:beaming}) with $z$ is not readily apparent in our data (see Fig.~\\ref{fig:b1vz}), but other factors entering the calculation of $B_1$ (particularly, the jet's radio flux density and angular length) have a broad scatter and probably serve to mask any relationship. A simple case to consider is one described by \\citet{2009A&ARv..17....1W}. If source volume is estimated via angular sizes in two dimensions assuming that the path through the jet is independent of redshift, then the volume $V \\propto d_{\\rm A}^2 \\propto d_{\\rm L}^2/(1+z)^4$, where $d_{\\rm A}$ and $d_{\\rm L}$ are angular and luminosity distances, respectively.\\footnote{The values of the volume reported in Paper I were computed incorrectly, so we provide the correct values of $V$, $B_1$, $K$ in Table~\\ref{tab:beaming}. The sense of the error is that the volumes in Paper I were too large, causing $B_1$ and $K$ to be too small by about a factor of 10 in some cases, and $\\theta$ to be about a factor of 2 larger than we now find.} For minimum energy (or equipartition), \\begin{equation} \\label{eq:fdmw} f(z) \\propto \\bigg [ \\frac {L_{\\rm s}(z)}{V(z)} \\bigg ]^{1/(\\alpha+3)} \\propto \\bigg [ \\frac {(1+z)^{(\\alpha -1)} d_{\\rm L}^2}{V(z)} \\bigg ]^{1/(\\alpha+3)} \\propto (1+z) \\end{equation} \\noindent Here we have assumed that the minimum-energy field is measured over fixed electron energies in the rest frame of the source. In the case of calculations over fixed frequencies in the observer's frame ($10^7$ to $10^{15}$ Hz are actually adopted for $B_1$ in Tables 5 and 6) the result is similar, with exponent $2/7$ rather than $1/(\\alpha+3)$, where (as in Paper I) we assume that $\\alpha=0.8$. Combining Equations \\ref{eq:rz} and \\ref{eq:fdmw} gives \\begin{equation} \\label{eq:Rdmw} R \\propto (1 + z)^2 \\end{equation} \\noindent which agrees with equation 13 of \\citet{2009A&ARv..17....1W}. Under this assumption for the jet volume, the fit value of $a$ is consistent with the prediction of the IC-CMB model. Alternatively, the volume can be estimated assuming that the jet is a cylinder of constant angular radius matched to {\\em Chandra}'s resolution (as adopted in Paper I and used for the estimates of $B_1$ in this paper). Here, $V(z) \\propto d_{\\rm A}^3$, so $d_{\\rm A}$ does not cancel in the equations, giving \\begin{equation} \\label{eq:fmh} f(z) \\propto (1+z) d_A^{-{1/(\\alpha + 3)}} \\end{equation} \\noindent In this case, $f(z)$ does not have a simple dependence on $(1+ z)$ over the redshift distribution of our sources. Instead, we define a new quantity that is derived from the observed data for each source, $Q \\equiv R B_1^{1+\\alpha}$. In the IC-CMB model, $Q \\propto (1+z)^{3+\\alpha}$, while our fit to $Q \\propto (1+z)^{a}$ gives $a = 1.35 \\pm 1.36$ (at 90\\% confidence). Here $a = 3+\\alpha$ is rejected at better than 99\\% confidence for $\\alpha > 0.5$. The best fit resulted in a smaller index, $a = -0.37 \\pm 1.35$, and $a = 3.5$ is still rejected at 99\\% confidence. Thus, we have two circumstances where the IC-CMB model can be ruled out and one in which it is still viable, where the jet volume is computing using the assumption described above Equation \\ref{eq:Rdmw}. The circumstances involve different but plausible conditions dictating the dependence of the intrinsic magnetic field with redshift, so it is difficult to provide a definitive test using these data alone. The factors that go into estimating the magnetic field bear further investigation as source details are obtained in follow-up radio, X-ray, and optical observations in order to develop a refined test of the model. One source of uncertainty in our method of using the X-ray and radio emission for the entire jet rather than for individual knots is that the jet geometries are often complex. Furthermore, the termination knots may also be included in some cases, where it is unlikely that both the radio and X-ray emission regions are moving relativistically relative to the nucleus. This paper is concerned primarily with shallow observations and deeper individual analyses would be best suited to examine these more subtle issues. \\subsubsection{Angles to the Line of Sight} As in Paper I, we computed the distribution of angles to the line of sight for these kpc scale jets, under the assumptions that 1) X-rays arise from the IC-CMB mechanism, and 2) all jets have a common Lorentz factor, $\\Gamma$. \\cite{vlba} estimated the intrinsic Lorentz factor distribution for a flux-selected set of core-dominated quasars, finding that it appears broad, with most values of $\\Gamma$ between 5 and 25 (see their Fig.~9). For now, we assume $\\Gamma = 15$ and find that $\\theta$ ranges from 6\\arcdeg\\ to 13\\arcdeg\\ for the quasars in our sample (see Table~\\ref{tab:beaming}). For these sources, the Doppler factor, $\\delta$, is in the range 3-8, compared to the assumed Lorentz factor of 15. Because the jet surface brightness is not constant and the spatial variations between the radio and X-ray bands can differ, it is possible that systematic errors result from considering the entire jet. To estimate the effect of restricting attention to knots within the jets, we have computed X-ray and radio flux densities for a selection of 3\\arcsec$\\times$3\\arcsec\\ regions from the jets. Measurements are given in Table~\\ref{tab:knotresults} and angles to the line of sight are given in Table~\\ref{tab:knotbeaming}. The angles usually decrease by a degree or less from the full-jet estimates. For the remainder of this section, we will only consider results for the entire jet, leaving analysis of individual knots to follow-up work which will require deeper X-ray observations with higher knot counts. See section~\\ref{sec:sources} for comments about individual sources and references to more detailed analyses, where available. Many of these sources are in the MOJAVE program, which consists of VLBI observations of several hundred compact active galaxies and quasars used to measure pc-scale proper motions. Of the 22 sources in common with our sample, we have X-ray data for 14, as listed in Table~\\ref{tab:bends}. For all but one quasar of the 22, there is apparent superluminal motion. Values of $\\beta_{\\rm app}$, the apparent velocity of the most rapidly moving pc-scale component relative to $c$, are given in the table. See the discussion of individual sources for references. The population modeling by \\citet{2007ApJ...658..232C} based on the MOJAVE sample provides a basis for testing the IC-CMB model for our sample. As a first step, it is important to determine that our sample is a representative subset of the MOJAVE sample. For the flux-limited MOJAVE sample, \\citet{2007ApJ...658..232C} showed that $\\sin \\theta$ of the pc-scale jets are generally within 50\\% of $1/\\beta_{\\rm app}$. Fig.~\\ref{fig:sldistribution} shows that the distribution of $1/\\beta_{\\rm app}$ for our sample is as concentrated below about 5\\arcdeg\\ as the MOJAVE sample. Also shown in this figure is the distribution of the values of $\\theta$ for the large scale jets, as derived from the IC-CMB model. These angles are generally below 11\\arcdeg\\ but systematically larger than the angles estimated for the pc-scale jets. This difference is not surprising because the pc-scale and kpc-scale jets are not aligned in projection on the sky but suggests that most misalignments are small. Position angle differences are given in Table~\\ref{tab:bends}. We now attempt to quantify the comparison of the angles to the line of sight for the kpc-scale jets with information in the pc-scale jets. In appendix~\\ref{app:angle}, we show how one may estimate the range of kpc-scale angles to the line of sight by using only the values of $\\beta_{\\rm app}$ for the pc-scale jets and the position angle differences. At the same time, intrinsic bend angles, $\\zeta$, are estimated and a probable range for these angles are computed. The results are given in Table~\\ref{tab:bends} and shown in Fig.~\\ref{fig:angles}, where it can be seen that these independent estimates are generally consistent. However, there are some notable exceptions, particularly where the angles from the IC-CMB calculation are of order a factor of two larger than those based on geometry and superluminal motion of the pc-scale jet. For these exceptions, one may infer that the jets decelerate substantially from pc to kpc scales." }, "1101/1101.5155_arXiv.txt": { "abstract": "We model the evolution of galaxy clustering through cosmic time to investigate the nature of the power-law shape of $\\xir$, the galaxy two-point correlation function. While $\\xir$ on large scales is set by primordial fluctuations, departures from a power law are governed by galaxy pair counts on small scales, subject to non-linear dynamics. We assume that galaxies reside within dark matter halos and subhalos. Therefore, the shape of the correlation function on small scales depends on the amount of halo substructure. We use a semi-analytic substructure evolution model to study subhalo populations within host halos. We find that tidal mass loss and, to a lesser extent, dynamical friction dramatically deplete the number of subhalos within larger host halos over time, resulting in a $\\sim 90\\%$ reduction by $z=0$ compared to the number of distinct mergers that occur during the assembly of a host halo. We show that these non-linear processes resulting in this depletion are essential for achieving a power-law $\\xir$. We investigate how the shape of $\\xir$ depends on subhalo mass (or luminosity) and redshift. We find that $\\xir$ breaks from a power law at high masses, implying that only galaxies of luminosities $ \\lesssim \\Lstar$ should exhibit power-law clustering. Moreover, we demonstrate that $\\xir$ evolves from being far from a power law at high redshift, toward a near power-law shape at $z = 0$. We argue that $\\xir$ will once again evolve away from a power law in the future. This is in large part caused by the evolving competition between the accretion and destruction rates of subhalos over time, which happen to strike just the right balance at $z \\approx 0$. We then investigate the conditions required for $\\xir$ to be a power law in a general context. We use the halo model along with simple parametrizations of the halo occupation distribution (HOD) to probe galaxy occupation at various masses and redshifts. We show that key ingredients determining the shape of $\\xir$ are the fraction of galaxies that are satellites, the relative difference in mass between the halos of isolated galaxies and halos that contain a single satellite on average, and the rareness of halos that host galaxies. These pieces are intertwined and we find no simple, universal rule for which a power-law $\\xir$ will occur. However, we do show that the physics responsible for setting the galaxy content of halos do not care about the conditions needed to achieve a power law $\\xir$ and these conditions are met only in a narrow mass and redshift range. We conclude that the power-law nature of $\\xir$ for $\\Lstar$ and fainter galaxy samples at low redshift is a cosmic coincidence. ", "introduction": "\\label{intro} The two-point correlation function of galaxies was measured four decades ago and found to be consistent with a $\\xir \\propto r^{-2}$ power law \\citep{totsuji69,peebles73,hauserpeebles73,peebleshauser74,peebles74}. Since that time, successively larger galaxy redshift surveys \\citep[e.g.,][]{huchra83,daCosta88,santiago95, shectman96,saunders00,colless01,york00a} have mapped the distribution of galaxies with ever increasing precision and confirmed correlation functions consistent with power laws over a large range of scales \\citep[e.g.,][]{delapparent88,marzke95,hermit96,tucker97,jing98,jing02,norberg02,zehavi02}. The scales on which a single power-law description is valid span a range from large regions exhibiting mild density fluctuations ($r \\gtrsim 10$~Mpc), to smaller regions with large density fluctuations experiencing rapid non-linear evolution ($r \\sim 1-10$~Mpc), to collapsed and virialized galaxy groups and clusters ($r \\lesssim 1$~Mpc). It has long been noted that the lack of any feature delineating the transitions among these scales is surprising \\citep[e.g.,][]{peebles74,gott_turner79,hamilton_tegmark02,masjedi06a,liwhite09}. This is especially true given that the matter correlation function in the now well-established concordance cosmological model differs significantly from a power law. In this paper, we return to this long-standing problem and address the origin of a power-law galaxy correlation function in the context of our modern paradigm for the growth of cosmic structure. This conundrum can be refined within the contemporary framework in which galaxies live within virialized halos of dark matter \\citep{whiterees78,blumenthal_etal84}. In such a model, galaxy clustering statistics can be modeled as a combination of dark matter halo properties and a halo occupation distribution (HOD) that specifies how galaxies occupy their host halos \\citep[e.g.,][]{peacock00a,scoccimarro01a,berlind02,cooray02}. In this {\\em halo model} approach, the galaxy correlation function is a sum of two terms: On small scales, pairs of galaxies reside in the same host dark matter halo (the ``one-halo'' term), whereas on large scales, the individual galaxies of a pair reside in distinct halos (the ``two-halo'' term). These two terms depend on the HOD in different ways, requiring delicate tuning in order to spawn an unbroken power law \\citep[e.g.,][]{berlind02}. Consequently, a feature in $\\xir$ at scales corresponding to the radii of the typical, viralized halos that host luminous galaxies is expected. In a dramatic success for the halo model, \\citet{zehavi04a} first detected a statistically-significant departure from a power law due to the high precision measurements of the Sloan Digital Sky Survey, and demonstrated that the halo model provides an acceptable fit to the data. \\citet{zehavi05a} confirmed this result, adding that power-law departures grow stronger with galaxy luminosity \\citep[see also][]{blake08,ross10}. $\\xir$ has since been shown to deviate from a power law at high redshifts \\citep{ouchi05,lee06,coil06,wake_etal11}. Nevertheless, it remains a fact that deviations from a power law at low redshifts are small and the galaxy correlation function is roughly a power law over an enormous range of galaxy-galaxy separations. Deviations have been revealed only through ambitious observational efforts. Halos are known to be replete with self-bound structures, dubbed ``subhalos'' \\citep{Ghigna98,Klypin99a,moore_etal99}, and both halos and subhalos are thought to be the natural sites of galaxy formation. Subhalos were isolated halos in their own right, hosting distinct galaxies before merging into a larger group or cluster halo\\footnote[1]{\\emph{Satellites} or \\emph{subhalos} are used throughout the paper to refer to self-bound entities lying within the virial radius of a larger halo. Those that do not lie within a larger system are designated as \\emph{centrals}, \\emph{host halos} or simply \\emph{hosts}.}. Remarkably, the clustering of host halos along with their associated subhalos is very similar to that of observed galaxies \\citep{kravtsovklypin99,colin99,kravtsov04a}, suggesting a simple correspondence of galaxies with host halos and subhalos. This was clearly demonstrated by \\citet{conroy06} who compared the correlation functions of hosts and subhalos to that of galaxies over a broad range of luminosities and redshifts ($z \\sim 0-4$), finding excellent agreement. These results indicate that an understanding of the physics governing the subhalo populations within host halos may provide insight into the physics of galaxy clustering and the near power-law form of the galaxy two-point correlation function. In this paper, we examine the causes of the observed power-law correlation function by studying the mergers, survival, and/or destruction of dark matter subhalos. Our focus in this paper is on the gross features of the galaxy two-point function and {\\em not} on detailed comparisons to specific data sets. We explore more sophisticated galaxy-halo assignments and statistical comparisons with data in a forthcoming follow-up study (Watson et al. in prep.). We argue that the nearly power-law, low-redshift galaxy correlation function is a coincidence. The correlation function of common $L \\lesssim \\Lstar$ galaxies evolves from relatively strong small-scale clustering at early times, through a power-law at the present epoch, and most likely toward relatively weak small-scale clustering in the future. The origin of the present-day power law, in turn, relies on the tuning of several disconnected ingredients, at least three of which are: the normalization of primordial density fluctuations determined by early Universe physics; a halo mass scale for efficient galaxy formation determined largely by atomic physics, stellar physics, and the physics of compact objects; and relative abundances of baryonic matter, dark matter, and dark energy in the Universe. Our paper is organized as follows. In \\S~\\ref{halomodel} we review the halo model and restate the problem in terms of this framework. In \\S~\\ref{model} we give an overview of our primary modeling technique. In \\S~\\ref{dynamics} we investigate the individual roles of merging, dynamical friction, and mass loss in shaping the halo occupation statistics of subhalos, as well as the resulting halo correlation function. In \\S~\\ref{dependence} we show how $\\xir$ depends on host halo mass and redshift. In \\S~\\ref{make_powerlaw} we explore a standard parametrization of the HOD to see what is required to get a power-law $\\xir$, and we predict the masses and redshifts at which a power-law $\\xir$ can be constructed. In \\S~\\ref{summary} we give a summary of our results and our primary conclusions. Throughout this paper, we work within the standard, vacuum-dominated, cold dark matter ($\\Lambda$CDM) cosmological model with $\\Omega_{\\mathrm{m}}=0.3$, $\\Omega_{\\Lambda}=0.7$, $\\Omega_{\\mathrm{b}}=0.04$, $h_{0}=0.7$, $\\sigma_{8}=0.9$, and $n_{\\mathrm{s}}=1.0$. These values differ slightly form the WMAP best-fit values, however this has little effect on our general results and was chosen in order to compare to previous work that used similar cosmological models. ", "conclusions": "\\label{summary} It has been recognized for decades that the two-point correlation function has a simple, power-law form with $\\xir \\sim r^{-2}$. Observational determinations of galaxy two-point clustering spanning more than thirty years all yielded results consistent with a single power law extending from linear and quasi-linear length scales ($r \\gtrsim 30\\, \\hmpc$) to deeply non-linear scales ($r \\lesssim 0.1\\, \\hmpc$). In this paper, we cast the problem in the contemporary setting in which galaxies form in halos and subhalos of dark matter and set out to understand the physical processes that drive this surprisingly simple result. Our primary conclusion is that the nearly power-law correlation function of relatively common, $\\Lstar$ and sub-$\\Lstar$ Galaxies at $z \\sim 0$ is a coincidence and does not reflect any general principle of structure formation or galaxy evolution. So how did we arrive at this conclusion? First, the efficiency of galaxy formation is dependent upon halo mass and it has been determined both theoretically and empirically that there is a halo mass scale below which galaxy formation is inefficient, roughly $M_{\\mathrm{gal}} \\sim 10^{10.5}\\, \\hMsun$ \\citep{conroy_wechsler09,behroozi10,guo10}. A number of things can set this scale including atomic and molecular physics and feedback from supernovae and active galactic nuclei \\citep[for a recent review article see][]{BensonReview10}. This mass scale is $M_{\\mathrm{gal}} < \\Mstar$, so $\\Lstar$ and sub-$\\Lstar$ galaxies are common. Had $M_{\\mathrm{gal}}$ been greater than or similar to $\\Mstar$, most bright galaxies would lie in comparably rare halos and be rare themselves. In such a case, one-halo clustering would be too strong to be compatible with a power law. $\\Mstar$ is {\\em not} determined by galaxy formation physics but is set by the completely unrelated processes that establish the amplitude of cosmological density fluctuations, presumably primordial inflation. Second, power-law clustering requires that some of the galaxies formed within relatively large subhalos are destroyed. Destruction is due primarily to mass loss, and, to a lesser extent, merging with the central galaxy as a result of dynamical friction. Without this destruction, satellite fractions would be too high and small-scale clustering too strong compared with large-scale clustering. In a forthcoming paper, we perform more sophisticated modeling to make the connection between subhalo mass loss and stellar mass loss in order to make predictions for the amount of intracluster light. Large-scale clustering is principally set by large-scale matter density fluctuations and is insensitive to the details of galaxy formation within halos, while the strength of small scale clustering grows in proportion to the fraction of galaxies that are satellites and in inverse proportion to the number density of the galaxies of interest. As it turns out, precisely the right amount of subhalo destruction has occurred by redshift $z \\sim 0$ in a concordance cosmology to produce a single, unbroken, power-law $\\xir$. Evolution of the satellite fraction is set by a competition between halo mergers, which increase $\\fsat$, and destruction by dynamical processes, which occur on a dynamical timescale and reduce $\\fsat$. At high redshifts, mergers occur more rapidly than destruction for halos with masses $\\gtrsim M_{\\mathrm{gal}}$. The low-redshift merger rate declines in part due to the fact that $M_{\\mathrm{gal}} < \\Mstar$ at $z \\lesssim 1$. Halos with masses below $\\Mstar$ become relatively more likely to merge with a larger object than to acquire new substructure compared to counterparts with masses greater than $\\Mstar$ \\citep[see][]{zentner07}. More importantly, the rate of halo mergers is quenched at $z \\lesssim 1$ as dark energy begins to suppress further cosmological structure growth. As merger rates decline, satellites are depleted with time. Therefore, at $z \\sim 0$, the correlation function is nearly a power law because the competition between the accretion and destruction rates has struck just the right balance to yield the appropriate value of $\\fsat$. The merger and destruction rates will once again become unbalanced in the future as halo merging is stifled by dark energy and existing satellite galaxies are slowly destroyed over many dynamical times through complex interactions in their host environments. We show that this will result in small-scale clustering that will be significantly {\\em too weak} to be consistent with a power law. Largely as a consequence of the merger/destruction competition, $\\xir$ evolves through cosmic time, achieving a power law only near $z\\sim 0$ for $L \\sim \\Lstar$ and dimmer galaxies. The processes of galaxy formation, the amplitude of cosmological density fluctuations, the abundance of dark matter, and the nature of the dark energy are thought to be completely distinct and determined by {\\em unrelated physics}. So the power-law $\\xir$ at $z \\sim 0$ is a coincidental conspiracy. In establishing these broad conclusions, we have performed an exhaustive investigation of the ingredients of the galaxy correlation function, which has revealed many interesting, more detailed conclusions. These can be summarized as follows. \\begin{enumerate} \\label{main_results} \\item{We find that satellite halo mass loss is the principle dynamical process responsible for depleting sufficient substructure so as to nearly align the one- and two-halo terms to yield a power-law correlation function at low redshift. Dynamical friction plays a smaller supporting role, accounting for an additional $\\sim 15\\%$ of subhalo destruction.} \\item{The shape of the correlation function is strongly mass dependent. For instance, at low redshift deviations from a power law $\\xir$ grow with increasing host halo mass. This drives stronger deviations from a power law for higher luminosity galaxy samples. The best power-law fits derived from our model are for galaxies residing in halos that are common enough to correspond to $\\sim \\Lstar$ and dimmer galaxies, in agreement with observations.} \\item{The correlation function is highly redshift-dependent. The sensitivity of the one-halo term to the HOD, coupled with the relative insensitivity of the two-halo term, implies that achieving a power-law requires fine-tuning the number of satellite galaxies per halo. The satellite galaxy abundance evolves with redshift, driven by the evolving balance between accretion and destruction, with an enhanced amount of substructure at high redshift. Therefore, the correlation function can only achieve a power law during those epochs when substructure has evolved to align the one- and two-halo terms. The correlation function is boosted on small scales at high $z$, the one- and two-halo terms join at $z = 0$ to form a power-law, then the power law is once again broken in future epochs.} \\item{For three chosen number densities corresponding to low-redshift, $\\sim \\Lstar$ and dimmer galaxies, we probed the most likely power-law space as a function of redshift for a parametrized HOD. We find that there is a relatively narrow range of satellite fractions for $\\xir$ to be consistent with a single power law (assuming $\\sim 10\\%$ measurement errors) at any given redshift. At all redshifts and masses, power-law correlation functions have satellite fractions in the range $\\fsat \\sim 0.1 - 0.25$. It is difficult to achieve a power-law correlation function at $z \\gtrsim 3$ for any number density.} \\item{We find that to achieve a power law $\\xir$ at high mass or redshift, the slope $\\alpha$ of the satellite galaxy occupation function must be significantly steeper than unity (for instance, greater than 2 at $z=3$). This would imply that the mapping of galaxies to halos is much more complicated than we think, since the number of galaxies would have to be very different than the number of subhalos of a particular size. Instead, it appears that the processes that govern galaxy formation do not care about the conditions needed to achieve a power law $\\xir$. } \\item{The ratio $\\Mone/\\Mmin$ (the ``plateau'' of the HOD) is a key ingredient for predicting the shape of $\\xir$. The prominence of the plateau is a measure of substructure abundance. Along with $\\Mone/\\Mmin$, it is also necessary to characterize the ratio $\\Mmin/\\Mstar$, which specifies what halo masses galaxies occupy relative to the halo mass function. By maintaining the combination of $\\Mone/\\Mmin \\sim 30$ and $\\Mmin/\\Mstar \\sim 0.05$ we can achieve a near power law for redshifts in the range $0-1.5$ and the appropriate mass threshold at each redshift (the mass threshold is $\\Mmin \\sim \\Mstar/20$, with $\\Mstar$ set by the redshift). At higher redshifts this criterion is met for galaxies that are most likely too dim to be observed. For example, achieving the requisite $\\Mmin \\sim \\Mstar/20$ at $z = 2$ corresponds to a halo mass of $\\Mmin \\sim 10^{9}\\, \\hMsun$ in which star formation is inefficient. } \\end{enumerate} This work has allowed us to formulate a general picture of the nature of the galaxy two-point correlation function. Halo abundances and subhalo populations evolve with time. At high redshifts, halos large enough to harbor galaxies are rare and subhalos are abundant within these hosts. With time, host halos that harbor galaxies generally become more common (though the specifics of this evolution can be subtle) and subhalos within these hosts become relatively less abundant. All the while, large-scale matter correlations grow, but the clustering bias of large halos evolves to largely compensate for this large-scale growth of structure. These effects, considered either individually or in tandem, change the HOD and the shape of $\\xir$. As a result, the correlation function evolves through an epoch where it is close to a power law and this epoch happens to be near $z\\sim 0$. From our broad discussion and detailed conclusions, it is clear that a nearly power-law correlation function requires a conspiracy between otherwise unrelated processes such as the early Universe physics that established the initial conditions for low redshift structure, the detailed physical processes that determine galaxy and star formation efficiency, and the growth rate of cosmic structure set largely by the abundances of dark matter and dark energy. The low-redshift power-law galaxy two-point function is thus a mere cosmic coincidence." }, "1101/1101.2543_arXiv.txt": { "abstract": "We explore the evolution of the emissions by accelerated electrons in shocked shells driven by jets in active galactic nuclei (AGNs). Focusing on powerful sources which host luminous quasars, we evaluated the broadband emission spectra by properly taking into account adiabatic and radiative cooling effects on the electron distribution. The synchrotron radiation and inverse Compton (IC) scattering of various photons that are mainly produced in the accretion disc and dusty torus are considered as radiation processes. We show that the resultant radiation is dominated by the IC emission for compact sources ($\\lesssim 10{\\rm kpc}$), whereas the synchrotron radiation is more important for larger sources. We also compare the shell emissions with those expected from the lobe under the assumption that a fractions of the energy deposited in the shell and lobe carried by the non-thermal electrons are $\\epsilon_e \\sim 0.01$ and $\\epsilon_{e, {\\rm lobe}} \\sim 1$, respectively. Then, we find that the shell emissions are brighter than the lobe ones at infra-red and optical bands when the source size is $\\gtrsim 10{\\rm kpc}$, and the IC emissions from the shell at $\\gtrsim 10~{\\rm GeV}$ can be observed with the absence of contamination from the lobe irrespective of the source size. In particular, it is predicted that, for most powerful nearby sources ($L_{\\rm j} \\sim 10^{47}~{\\rm ergs~s^{-1}}$), $\\sim {\\rm TeV}$ gamma-rays produced via the IC emissions can be detected by the modern Cherenkov telescopes such as MAGIC, HESS and VERITAS. ", "introduction": "It is well established that radio-loud active galactic nuclei (AGNs) are accompanied by relativistic jets \\citep[e.g.,][for review]{BBR84}. These jets dissipate their kinetic energy via interactions with surrounding interstellar medium (ISM) or intracluster medium (ICM), and inflate a bubble composed of decelerated jet matter, which is often referred to as cocoon. Initially, the cocoon is highly overpressured against the ambient ISM/ICM \\citep{BC89} and a strong shock is driven into the ambient matter. Then a thin shell is formed around the cocoon by the compressed ambient medium. The thin shell structure persists until the cocoon pressure decreases and the pressure equilibrium is eventually achieved \\citep{RHB01}. While large number of radio observations identified cocoons with the extended radio lobe, no clear evidence of radio emissions is found for the shocked shells \\citep[e.g.,][]{CPD88}. Due to the lack of detections, in the previous studies on the extragalactic radio sources, it is usually assumed that non-thermal emissions are dominated by or originated only in the cocoon \\citep[e.g.,][]{SBM08}. However, since strong shocks are driven into tenuous ambient gas with a high Mach number, the shocked shells are expected to offer site of particle acceleration as in the shocks of supernova remnants (SNRs) % and therefore give rise non-thermal emissions \\citep{FKY07, B08}. Hence, although the observations at radio seems to be unsuccessful, non-thermal emissions from the shell may be accessible at higher frequencies. In fact, while not detected in radio, recent deep X-ray observation have reported % the presence of non-thermal emissions from the shell associated with the radio galaxy Centaurus A \\citep{CKH09}. Theoretically, emissions by the accelerated particles residing in shocked shells have been studied by \\citet{FKY07}. However, they paid attention only to the the extended sources of $\\sim 100~{\\rm kpc}$ and the inverse Compton (IC) scattering of external photons was not included in the radiative processes. The cooling effects on the energy distribution of non-thermal electrons were not considered, either. Motivated by these backgrounds, we explore in this paper the temporal evolution of the non-thermal emissions by the accelerated electrons in the shocked shells, properly taking into account the Comptonization of photons of various origins as well as the cooling effects on the electron distribution. Focusing on the powerful sources which host luminous quasar in its core, we show that the shell can produce prominent emissions ranging from radio up to $\\sim 10~{\\rm TeV}$ gamma-ray and discuss the possibility for the detection. The paper is organized as follows. In \\S\\ref{model}, we introduce the dynamical model, which describes the evolution of the shell expansion, and explain how the energy distribution of electrons residing in the shell and the spectra of the radiations they produce are evaluated based on the dynamical model. The obtained results are presented in \\S\\ref{results}. In \\S\\ref{comp}, we compare the emissions from the shell with those from the lobe and discuss its detectability. We close the paper with the summary in \\S\\ref{summ}. ", "conclusions": "\\label{summ} We have explored the temporal evolution of the emissions by accelerated electrons in the shocked shell produced by AGN jets. Focusing on the powerful sources which host luminous quasars, we have calculated the spectra of the synchrotron emission as well as the IC scatterings of various photons that will be relevant in this context. We have used a simple analytic model that describes the dynamics of the expanding shell and estimated the energy distribution of non-thermal electrons based on this model, taking properly into account both the adiabatic and radiative coolings. Below we summarize our main findings in this study. \\vspace{2mm} 1. When the source is small ($R \\lesssim R_{\\rm IC/syn} \\sim 27 L_{\\rm IR, 46}^{1/2} B_{-5}^{-2}{\\rm kpc}$), the dominant radiative process is the IC scattering of IR photons emitted from the dust torus. For larger sources, on the other hand, the synchrotron emissions dominate over the IC emissions, since the energy density of photons becomes smaller than that of magnetic fields ($U_{B} > U_{\\rm ph} \\propto R^{-2}$). Through the entire evolution, the spectrum is rather broad and flat, and the peak luminosity is approximately given by $(\\nu L_{\\nu})_{\\rm peak} \\sim 4.0 \\times 10^{40} \\epsilon_{-2}L_{45}~{\\rm ergs~s^{-1}}$, since it is roughly equal to the energy injection rate, which is in turn determined by the jet power $L_{\\rm j}$ and acceleration efficiency $\\epsilon_e$. \\vspace{2mm} 2. The broadband spectra extend from radio up to $\\sim 10~{\\rm TeV}$ gamma-ray energies for a wide range of source size ($R\\sim 1-100~{\\rm kpc}$) and jet power ($L_{\\rm j}\\sim 10^{45}-10^{47}~{\\rm ergs~s^{-1}}$). By comparing the emissions with those from the lobe, we find that the synchrotron emissions at IR and optical frequencies can be observed without being hampered by the lobe emissions for extended sources ($R \\gtrsim R_{\\rm IC/syn}$), while the IC emissions at $h \\nu \\gtrsim 10~{\\rm GeV}$ can be observed with the absence of contamination from the lobe irrespective of the source size. In particular, it is predicted that, for most powerful nearby sources ($L_{\\rm j} \\sim 10^{47}~{\\rm ergs~s^{-1}}$, $D \\lesssim 100~{\\rm Mpc}$), $\\sim {\\rm TeV}$ gamma-rays produced via the IC emissions can be detected by the modern Cherenkov telescopes such as MAGIC, HESS and VERITAS. %" }, "1101/1101.5612_arXiv.txt": { "abstract": "{ GRB 050509b, detected by the \\emph{Swift} satellite, is the first case where an X-ray afterglow has been observed associated with a short gamma-ray burst (GRB). Within the fireshell model, the canonical GRB light curve presents two different components: the proper-GRB (P-GRB) and the extended afterglow. Their relative intensity is a function of the fireshell baryon loading parameter $B$ and of the CircumBurst Medium (CBM) density ($n_{CBM}$). In particular, the traditionally called short GRBs can be either ``genuine'' short GRBs (with $B \\lesssim 10^{-5}$, where the P-GRB is energetically predominant) or ``disguised'' short GRBs (with $B \\gtrsim 3.0 \\times 10^{-4}$ and $n_{CBM}\\ll1$, where the extended afterglow is energetically predominant). } { We verify whether GRB 050509b can be classified as a ``genuine'' short or a ``disguised'' short GRB, in the fireshell model. } { We investigate two alternative scenarios. In the first, we start from the assumption that this GRB is a ``genuine'' short burst. In the second attempt, we assume that this GRB is a ``disguised'' burst. } { If GRB 050509b were a genuine short GRB, there should initially be very hard emission which is ruled out by the observations. The analysis that assumes that this is a disguised short GRB is compatible with the observations. The theoretical model predicts a value of the extended afterglow energy peak that is consistent with the Amati relation. } { GRB 050509b cannot be classified as a ``genuine'' short GRB. The observational data are consistent with a ``disguised'' short GRB classification, i.e., a long burst with a weak extended afterglow ``deflated'' by the low density of the CBM. We expect that all short GRBs with measured redshifts are disguised short GRBs because of a selection effect: if there is enough energy in the afterglow to measure the redshift, then the proper GRB must be less energetic than the afterglow. The Amati relation is found to be fulfilled only by the extended afterglow excluding the P-GRB. } ", "introduction": "The traditional classification of gamma ray bursts (GRBs) is based on the observed time duration of the prompt emission measured with the criterion of ``$T_{90}$'', which is the time duration in which the cumulative counts increase from $5\\%$ to $95\\%$ above the background, encompassing $90\\%$ of the total GRB counts. This parameter shows that there are two groups of GRBs, the short ones with $T_{90}<2$ s, and the long ones with $T_{90}>2$ s. This analysis motivated the standard classification in the literature of short and long GRBs \\citep{1992grbo.book..161K,1992AIPC..265..304D,1993ApJ...413L.101K}. The observations of GRB 050509b by BAT and XRT on board the \\emph{Swift} satellite \\citep[see][]{2004ApJ...611.1005G,2005SSRv..120..165B} represent a new challenge to the classification of GRBs as long and short, since it is the first short GRB associated with an afterglow \\citep{2005Natur.437..851G}. Its prompt emission observed by BAT lasts 40 milliseconds, but it also has an afterglow in the X-ray band observed by XRT, which begins 100 seconds after the BAT trigger (time needed to point XRT to the position of the burst) and lasts until $\\approx$ 1000 seconds. It is located 40 kpc away from the center of its host galaxy \\citep[][, see Fig. \\ref{fbloom}]{2006ApJ...638..354B}, which is a luminous, non-star-forming elliptical galaxy with redshift $z=0.225$ \\citep{2005Natur.437..851G}. Although an extensive observational campaign has been performed using many different instruments, no convincing optical-IR candidate afterglow nor any trace of any supernova has been found associated with GRB 050509b \\citep[see][]{2005GCN..3401....1C,2005GCN..3521....1B,2005ApJ...630L.117H,2005A&A...439L..15C,2005GCN..3386....1B,2005GCN..3417....1B,2006ApJ...638..354B}. An upper limit in the $R$-band $18.5$ days after the event onset imply that the peak flux of any underlying supernova should have been $\\sim 3$ mag fainter than the one observed for the type Ib/c supernova SN 1998bw associated with GRB 980425, and $2.3$ mag fainter than a typical type Ia supernova (\\citealp{2005A&A...439L..15C}, see also \\citealp{2005ApJ...630L.117H}). An upper limit to the brightening caused by a supernova or supernova-like emission has also been established at $8.17$ days after the GRB: $R_c \\sim 25.0$ mag \\citep{2006ApJ...638..354B}. While some core-collapse supernovae might be as faint as (or fainter than) this limit \\citep{2007Natur.449E...1P}, the presence of this supernova in the outskirts of an elliptical galaxy would be truly extraordinary \\citep{2005A&A...433..807M,2005PASP..117..773V}. Unfortunately, we cannot obtain exhaustive observational constraints for this GRB because XRT data are missing in-between the first 40 milliseconds and 100 seconds. However, this makes the theoretical work particularly interesting, because we can infer from first principles some characteristics of the missing data, which are inferred by our model, and consequently reach a definite understanding of the source. This is indeed the case, specifically, for the verification of the Amati relation \\citep{2002A&A...390...81A,2006MNRAS.372..233A,2009A&A...508..173A} for these sources as we see in section \\ref{amati}. \\begin{figure} \\centering \\includegraphics[width=\\hsize]{gband} \\caption{Keck LRIS G-band image, zoomed to show the XRT error circle. The larger, blue circle is the revised XRT position from \\citet{2005GCN..3395....1R}; the smaller, green circle to the west and north of that is the $2\\sigma$ confidence region of the XRT position computed in \\citet{2006ApJ...638..354B}. The 11 sources consistent with the \\citet{2005GCN..3395....1R} X-ray afterglow localization are labeled in the image. North is up and east is to the left. G1 is the large galaxy to the west and south of the XRT. Bad pixel locations are denoted with ``BP''. Figure reproduced from \\citet{2006ApJ...638..354B} with the kind permission of J. Bloom and of the AAS.} \\label{fbloom} \\end{figure} GRB 050509b is an example that the usual classification is at least incomplete. Within the fireshell model, we propose three classes of GRBs: long, genuine short and disguised short \\citep[][and references therein]{2009AIPC.1132..199R}. We have a well-defined way of differentiating between the classes, which is based on two parameters, the baryon loading parameter $B$ and the CircumBurst Medium (CBM) number density $n_{CBM}$ (see next section), that help to make the classification clearer. In this paper, we analyze GRB 050509b within the fireshell model. We proceed with the identification of the two basic parameters, $B$ and $n_{CBM}$, within two different scenarios. We first investigate the ``ansatz'' that this GRB is the first example of a ``genuine'' short bursts. After disproving this possibility, we show that this GRB is indeed another example of a disguised short burst. In the next section, we briefly introduce the fireshell model and explain the classification, in section \\ref{analysis} we show the analysis of the data, in section \\ref{amati} we present the theoretical spectrum and the study of the fulfillment of the Amati relation, in section \\ref{sec:discussions} we comment on the results, and in section \\ref{sec:conclusions} we finally present our conclusions. ", "conclusions": "\\label{sec:conclusions} It has been shown that GRB 050509b originates from the gravitational collapse to a black hole of a merging binary system consisting of two degenerate stars according to three different and complementary considerations: \\begin{enumerate} \\item Very stringent upper limits on an associated supernova event have been established \\citep[see][]{2005GCN..3401....1C,2005GCN..3521....1B,2005ApJ...630L.117H,2005A&A...439L..15C,2005GCN..3386....1B,2005GCN..3417....1B,2006ApJ...638..354B}; \\item The host galaxy has been identified with a luminous, non-star-forming elliptical galaxy \\citep{2005GCN..3390....1P,2005Natur.437..851G,2005GCN..3386....1B,2006ApJ...638..354B}; \\item The GRB exploded in the halo of the host galaxy \\citep{2006ApJ...638..354B}, because the binary system spiraled out before merging. \\end{enumerate} From an astrophysical point of view, there are three possible cases of merging binary systems that must be considered: \\begin{enumerate} \\item Neutron star / neutron star: unlike the case of GRB 970228 \\citep{2007A&A...474L..13B}, the low energetics of GRB 050509b disfavor this hypothesis; \\item Neutron star / white dwarf: this appears to be the most likely case for GRB 050509b, as in GRB 060614 \\citep{2009A&A...498..501C} and in GRB 071227 \\citep{2010A&A...521A..80C}; \\item White dwarf / white dwarf: this case is viable only for two very massive white dwarfs, allowing the critical mass of neutron stars against gravitational collapse to a black hole to be overcome in the merging process; that low massive white dwarf / white dwarf merging binary systems may lead to low energetics events has been widely expressed in the literature \\citep[see e.g.][]{1984ApJS...54..335I,1985ASSL..113....1P,2010Natur.463...61P}. \\end{enumerate} From the point of view of GRB classification, we conclude that: \\begin{enumerate} \\item GRB 050509b is a disguised short GRB occurring in a low CBM density environment ($n_{CBM} < 10^{-3}$ particles/cm$^3$), typical of a galactic halo; \\item The baryon loading of GRB 050509b, and consequently the ratio of the P-GRB to the extended afterglow energetics, is typical of canonical long-duration GRBs; \\item The possible origin of a genuine short GRB from a merging binary system, as often purported in the literature (see e.g. \\citealp{2006RPPh...69.2259M} but also \\citealp{2009ARA&A..47..567G}), still remains an open issue both from an observational and a theoretical point of view; in theory, this will crucially depend on the amount of baryonic matter left over in the process of gravitational collapse originating the fireshell baryon loading, which must be $B \\lesssim 10^{-5}$. \\end{enumerate} From all the above considerations, it also follows that a binary system merging in a higher density region (i.e. $n_{CBM} \\sim 1$ particles/cm$^3$) would give rise to a canonical long-duration GRB without an associated supernova \\citep[see also][]{2007A&A...474L..13B,2009A&A...498..501C}." }, "1101/1101.2204_arXiv.txt": { "abstract": "Using a sample of 92 UV continuum-selected, spectroscopically identified galaxies with $\\langle z \\rangle = 2.65$, all of which have been imaged in the \\lya\\ line with extremely deep narrow-band imaging, we examine galaxy \\lya\\ emission profiles to very faint surface brightness limits. The galaxy sample is representative of spectroscopic samples of LBGs at similar redshifts in terms of apparent magnitude, UV luminosity, inferred extinction, and star formation rate and was assembled without regard to \\lya\\ emission properties. Approximately 45\\% (55\\%) of the galaxy spectra have \\lya\\ appearing in net absorption (emission), with $\\simeq 20$\\% satisfying commonly used criteria for the identification of ``Lyman Alpha Emitters'' (LAEs) [$W_0(\\lya) \\ge 20$ \\AA]. We use extremely deep stacks of rest-UV continuum and continuum-subtracted \\lya\\ images to show that all sub-samples exhibit diffuse \\lya\\ emission to radii of at least 10\\arcs\\ ($\\sim 80$ physical kpc). The characteristic exponential scale lengths for \\lya\\ line emission exceed that of the $\\lambda_0=1220$ \\AA\\ UV continuum light by factors of $\\sim 5-10$. The surface brightness profiles of \\lya\\ emission are strongly suppressed relative to the UV continuum light in the inner few kpc, by amounts that are tightly correlated with the galaxies' observed spectral morphology; however, all galaxy sub-subsamples, including that of galaxies for which \\lya\\ appears in net absorption in the spectra, exhibit qualitatively similar diffuse \\lya\\ emission halos. Accounting for the extended \\lya\\ emission halos, which generally would not be detected in the slit spectra of individual objects or with typical narrow-band \\lya\\ imaging, increases the total \\lya\\ flux [and rest equivalent width $W_0(\\lya)$] by an average factor of $\\sim 5$, and by a much larger factor for the 80\\% of LBGs not classified as LAEs. We argue that most, if not all, of the observed \\lya\\ emission in the diffuse halos originates in the galaxy \\ion{H}{2} regions but is scattered in our direction by \\ion{H}{1} gas in the galaxy's circum-galactic medium (CGM). The overall intensity of \\lya\\ halos, but not the surface brightness distribution, is strongly correlated with the emission observed in the central $\\sim 1$\\arcs\\-- more luminous halos are observed for galaxies with stronger central \\lya\\ emission. We show that whether or not a galaxy is classified as a giant ``Lyman $\\alpha$ Blob'' (LAB) depends sensitively on the \\lya\\ surface brightness threshold reached by an observation. Accounting for diffuse \\lya\\ halos, all LBGs would be LABs if surveys were sensitive to 10 times lower \\lya\\ surface brightness thresholds; similarly, essentially all LBGs would qualify as LAEs. ", "introduction": "Although the Lyman $\\alpha$ (\\lya) emission line of neutral H is expected to be produced in prodigious amounts by star-forming galaxies (e.g. \\citealt{partridge67,meier76a}), it has long been appreciated that the astrophysics affecting observations of \\lya\\ are far more complex than for other lines of abundant species due to resonant scattering (\\citealt{spitzer78,meier81,charlot93}). The very large cross-section in the \\lya\\ transition means that emission from a gas cloud or nebula may have been strongly altered in intensity, kinematics, and apparent spatial distribution by the time it reaches an observer. Similarly, information about the initial source of observed \\lya\\ emission may be lost or obscured, with the apparent source simply being \\ion{H}{1} gas responsible for scattering in the observer's direction. Consequently, the dominant process producing \\lya\\ emission may often be ambiguous; possibilities include photoionization by young stars or AGN, line emission following collisional excitation of H atoms, or simply scattering from intervening \\ion{H}{1} gas that happens to favor the observer's direction. In the absence of dust, the standard expectation for \\lya\\ emission produced in \\ion{H}{2} regions for ``Case B'' (i.e., ionization-bounded) recombination \\citep{brockle71} and a \\cite{chabrier03} stellar initial mass function (IMF) for high mass stars \\footnote{Note that this value is a factor 1.8 higher than would be obtained assuming a Salpeter (1955) IMF because a given number of ionizing photons is associated with a smaller total SFR for the Chabrier IMF.} is that each solar mass of star formation produces a \\lya\\ luminosity $L(\\lya) \\simeq 2.0 \\times 10^{42}$ ergs s$^{-1}$. For the same IMF, the far-UV continuum light produced per solar mass of SFR near the wavelength of \\lya\\ has an expected monochromatic luminosity in the range $40.0 \\simlt{\\rm log}~L_{\\lambda,cont} \\simlt 40.3$ ergs s$^{-1}$ \\AA$^{-1}$ (\\citealt{leitherer99}). The predicted rest equivalent width of \\lya\\ emission is then given by $W_0(\\lya) \\simeq L(\\lya)/L_{\\lambda,cont} \\simeq 100 - 200$ \\AA\\ (see also \\citealt{charlot93}), with values near the lower end of this range expected for continuous star formation lasting more than $\\sim 3 \\times 10^7$ yrs, roughly the minimum dynamical timescale for L* LBGs at $z \\sim 2-3$ (e.g., \\citealt{erb+06b}). Under the above assumptions, the period of time over which \\lya\\ emission has $W_0(\\lya) > 100$ \\AA\\ would be very brief, after which the line-to-continuum ratio reaches an asymptotic value of $W_0(\\lya) \\simeq 100$ \\AA. Thus, for a UV continuum-selected sample, one would expect only a small fraction of galaxies to be caught during a time when their intrinsic $W_0(\\lya)$ exceeds 100 \\AA \\footnote{For a sample selected by \\lya\\ (as opposed to continuum) emission, this may not be the case.}. When dust is mixed throughout the scattering medium, one expects selective extinction of \\lya\\ photons compared to those in the nearby UV continuum due to the much larger effective path length traversed by a line photon before escaping into the intergalactic medium (e.g., \\citealt{meier81,hartmann84,neufeld90}). This effect is often cited when observed \\lya\\ emission lines are much weaker than the Case B expectations discussed above (e.g., \\citealt{charlot93,shapley03,hayes10,kornei10}). Since most continuum-selected high redshift galaxies in current spectroscopic surveys appear to have at least some dust, and the vast majority have \\lya\\ equivalent widths $W_0(\\lya) < 100$ \\AA\\ (e.g., \\citealt{shapley03,kornei10}), this conclusion would seem reasonable. On the other hand, it is also possible, at least in principle, for \\lya\\ photons to experience {\\it less} attenuation by dust than continuum photons, in the case of a clumpy ISM in which dust is located only within the clumps which are rarely penetrated by \\lya\\ photons (\\citealt{neufeld91, finkelstein08}). There is no reason to believe that the two competing effects could not {\\it both} be at work within different regions of the same galaxy. Even without dust, however, resonant scattering produces spatial and/or spectral diffusion of \\lya\\ photons leading to emergent line emission whose properties depend on the geometry, kinematics, and \\ion{H}{1} optical depth distributions within the gaseous circumgalactic medium (CGM) surrounding a galaxy (\\citealt{steidel2010} [S2010]). In the zero-dust case, the total \\lya\\ luminosity would be unaltered by resonant scattering, but, as we detail below, the {\\it detectability} of \\lya\\ could be very strongly affected. \\begin{figure}[thb] \\centerline{\\includegraphics[height=9cm]{f1.eps}} \\caption{ Comparison of \\lya\\ line equivalent widths measured from spectra compared to those inferred from Cont-NB colors in deep\\lya\\ imaging. The imaging measurements use isophotal apertures defined by the extent of \\lya\\ flux to a surface brightness limit of $\\simeq 1-2 \\times 10^{-18}$ ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$, which is typical of the deepest \\lya\\ narrow-band imaging surveys. \\label{fig:ew_vs_ew} } \\end{figure} \\begin{figure}[thb] \\centerline{\\includegraphics[height=9cm]{f2.eps}} \\caption{Comparison of the \\lya\\ line equivalent width distribution from spectroscopic measurements versus that inferred from CB-NB colors in \\lya\\ imaging. The imaging measurements use colors within isophotal apertures defined by the extent of \\lya\\ flux to a surface brightness limit of $\\simeq 1-2 \\times 10^{-18}$ ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$, which is typical of the deepest \\lya\\ narrow-band imaging surveys. The statistics are for the mean and standard deviation (left) of individual values (left), and the median and inter-quartile range (right) for each set of measurements. \\label{fig:ew_hist} } \\end{figure} In S2010, we characterized the distribution of cool gas in the CGM of star-forming galaxies with redshifts $2 \\simlt z \\simlt 3$ and attempted to understand the kinematics and line strength of the ISM absorption and \\lya\\ emission in the context of galaxy-scale gaseous outflows. In brief, we found that UV-selected galaxies within a factor of a few of L* in the far-UV continuum luminosity function (corresponding at $z \\sim 2.5$ to apparent magnitudes ${\\cal R} \\simeq 24-24.5$-- see \\citealt{reddy09}) have a CGM that can be traced by \\ion{H}{1} (\\lya\\ and \\lyb\\ absorption) and several strong absorption lines of metallic species (e.g., \\ion{C}{2}, \\ion{C}{4}, \\ion{Si}{2}, \\ion{Si}{4}) to galactocentric distances of $\\simlt 120$ kpc using the spectra of faint background galaxies. The measurement used more than 500 galaxy pairs on angular scales $1-15$\\arcs\\ to map out the absorption line strength as a function of galaxy impact parameter $b$ (i.e., the physical separation of the two lines of sight at the redshift of the foreground galaxy) for each observed species. In slit spectra of the CGM ``host galaxies'', the bulk of observed \\lya\\ emission, when present, is almost always strongly redshifted, while the strong interstellar (IS) absorption lines are strongly blue-shifted. S2010 presented a geometric and kinematic model that reproduces many of the observed trends. In the context of the model, \\lya\\ photons escape the galaxy in an observer's direction mainly by scattering from optically thick \\ion{H}{1} gas located on the far side of the galaxy's stars, but having the same overall (outflowing) kinematics as the IS gas seen in blue-shifted absorption. We used the transverse information from the galaxy pairs combined with line-of-sight information available from the galaxies' own far-UV spectra to construct a consistent geometric and kinematic model of galaxy scale outflows in the context of a very well-studied population of high redshift star-forming galaxies. That is, we combined the line profiles of IS absorption lines and \\lya\\ emission in the galaxy spectra themselves (sampling the kinematics and line strength for galactocentric impact parameter $b \\sim 0$) with IS line strength measurements at $b>>0$ (using close angular pairs of galaxies) to infer the 3-dimensional distribution of CGM gas surrounding an average galaxy in the spectroscopic sample. We suggested that the CGM gas seen in absorption would also constitute a scattering medium through which \\lya\\ photons must traverse in order to be observed. High velocities and large velocity gradients together with gas covering fraction $f_c \\le 1$ through much of the CGM allow \\lya\\ photons to diffuse spatially outward, favoring escape of \\lya\\ photons last scattered (in the observer's direction) from atoms with velocities well off resonance with respect to any \\ion{H}{1} that remains between the location of the last scattering and the observer. If true, one might then expect to observe scattered \\lya\\ emission over the same spatial scales for which strong HI and low-ion metallic {\\it absorption} is seen, i.e., $\\simeq 80-90$ kpc, even if all \\lya\\ photons originated in the galaxy's \\ion{H}{2} regions. Clearly, scattering will substantially modify both the spatial and spectral distribution of \\lya\\ photons emergent in a particular direction, and at the very least may cause \\lya\\ emitting regions to appear distinct from the UV continuum emission even if both share a common origin. Slit spectra commonly optimized for the compact size of the continuum emitting regions of typical star-forming high redshift galaxies may encompass only a fraction of emergent \\lya\\ emission. The relevant angular scale for the optically-thick CGM \\ion{H}{1} gas is $\\simeq 10$\\arcs\\ ($\\simeq 80$ physical kpc at $z\\sim 2.5$), whereas a typical extraction aperture for a slit spectrum is $\\sim 1\\secpoint2 \\times 1\\secpoint4$ -- a difference of a factor of more than 180 in solid angle. Thus, even if the Case B-expected production rate of \\lya\\ photons were to escape the CGM of a galaxy, it is likely that the emission would be distributed over such a large region that a narrow slit would miss most of the \\lya\\ flux; even very deep narrow-band images might leave much of the flux unaccounted-for due to limited surface brightness sensitivity. In this paper, we present direct observational evidence showing that extended \\lya\\ scattering ``halos'' are a generic property of high redshift star-forming galaxies, including those that have no apparent \\lya\\ emission lines in their far-UV spectra. In \\S2 we describe a sample of 92 UV-continuum-selected galaxies for which both rest-far-UV spectra and deep narrow-band \\lya\\ images are available, and discuss the relationship between \\lya\\ properties measured using both techniques. In \\S3 we use composite UV spectra, as well as \\lya\\ and continuum image stacking, to measure \\lya\\ emission extending to very low surface brightness thresholds for various galaxy sub--samples. The results and their implications for the nature of \\lya\\ emission in star-forming galaxies are described in \\S4, discussed in \\S5, and summarized in \\S6. Throughout the paper we assume a Lambda-CDM cosmology with $\\Omega_m = 0.3$, $\\Omega_{\\Lambda}=0.7$, and $h=0.7$. ", "conclusions": "We have shown above that, on average, LBGs with far-UV luminosities $0.3\\simlt (L/L_{UV}^*) \\simlt 3$ at $\\langle z \\rangle = 2.65$ exhibit spatially extended \\lya\\ emission to physical radii of at least 80 kpc (10\\arcs), even when \\lya\\ appears only in absorption for regions coincident with the UV continuum starlight. Figures~\\ref{fig:sb_profile}, \\ref{fig:emabs_plot}, and \\ref{fig:blob_plot_mod} show that the {\\it profiles} of the \\lya\\ emission are quite similar in shape independent of the spectral morphology, with the main difference being the overall intensity normalization and the presence or absence of emission spatially coincident with the continuum light (i.e., the inner $\\pm 5$ kpc). The observations suggest that the \\lya-scattering CGM may be statistically universal, with the main variable being the fraction of \\lya\\ photons able to emerge from the inner few kpc region without being destroyed. For example, the difference between the \\lya\\ Em and \\lya\\ Abs (see Table~\\ref{table:tab2}) spectrally classified subsets is an overall factor of $\\sim 5$ in the \\lya\\ surface brightness at the full continuum extent (Figure~\\ref{fig:emabs_plot}), beyond which the ratio of $S(b)$ for the two sub-samples remains essentially constant. The scale lengths for \\lya\\ emission ($b_l \\simeq 25\\pm3$ kpc) are consistent among the statistically distinct galaxy sub-samples in spite of the fact that the integrated line-to-continuum ratio varies by large factors among the same sub-samples. \\subsection{Previous Results on Statistical \\lya\\ Detections} \\lya\\ emission with physical extent larger than that of a galaxy's continuum starlight is not a surprising result from a theoretical perspective (e.g., \\citealt{barnes09,barnes10,laursen09a,laursen09b}), and has been observed and noted in many individual cases both in the nearby (e.g. \\citealt{mas-hesse03,hayes07,ostlin09}) and high redshift (e.g.,\\citealt{franx97,moller98,steidel00,fynbo03,matsuda04,adelberger06,ouchi08}) universe. However, relatively few surveys at high redshift have reached adequate \\lya\\ surface brightness limits to allow the detection of the very low surface brightness levels discussed above. An exception is the extremely deep spectroscopic survey for \\lya\\ emission conducted by \\cite{rauch08} [R08]. Using a \\lya\\ --selected sample distributed over the redshift range $2.7 \\le z \\le 3.8$, these authors noted that extended \\lya\\ emission was a common feature of the LAEs discovered in their survey. A spatial stack of all of the \\lya\\ emitting sources exhibited significant emission (with threshold $\\simeq 1.5\\times 10^{-19}$ ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$) to an angular scale of $\\sim 4$\\arcs, or $\\sim 30$ kpc projected physical radius. The R08 sample, as the authors themselves point out, covers a different range of UV luminosity compared to most continuum-selected LBG spectroscopic surveys-- only one of 27 objects has $V < 25.5$, while 80\\% our sample (which has a median $V\\simeq 25.0$) has $V < 25.5$, although there is is a tendency for the faintest objects to be among those with the strongest \\lya\\ emission lines (see Tables~\\ref{table:tab2} and ~\\ref{table:tab3}). \\footnote{Moving our continuum-selected sample to the somewhat higher median redshift of R08 would result in $\\simeq 50$\\% of our sample having $V > 25.5$. } Nevertheless, the average surface brightness profile for the R08 \\lya-selected sample is remarkably similar to that of our continuum-selected sample (e.g., compare Figure~\\ref{fig:emabs_plot} to Figure 20 of R08). For objects in our ``\\lya\\ Em'' sub-sample (Table~\\ref{table:tab2}), the peak \\lya\\ SBs are somewhat higher than for the R08 sample, while the angular extent (at the same limiting SB of $\\sim 1\\times10^{-19}$ ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$) is $\\simeq 2.5-3$ times larger in the present LBG sample. Within our sample there is a significant dependence of $W_0(\\lya)$ on apparent UV continuum luminosity, but the average \\lya\\ profiles are similar, as shown in Figure~\\ref{fig:Rs_compare}. \\begin{figure}[thb] \\centerline{\\includegraphics[width=9cm]{f14.eps}} \\caption{A comparison of the continuum and \\lya\\ surface brightness profiles of the full sample divided into two at the median continuum apparent magnitude. The ``UV bright'' sample is a factor of $\\simeq 2.0$ times brighter in the continuum than that of the ``UV faint'' sample (CB(Bright)$=24.22$ versus CB(Faint)$=24.95$), but the average \\lya\\ flux for the UV bright sub-sample is 10\\% {\\it smaller} than that of the UV-faint sub-sample, i.e. $W_0(\\lya,bright)=22.0$ \\AA, while $W_0(\\lya,faint)=48.5$ \\AA). \\label{fig:Rs_compare} } \\end{figure} In any case, it is worth pointing out that, under the hypothesis that \\lya\\ scattering, and not fluorescence, is the dominant process producing the observed \\lya\\ halos, the scattering medium need not be optically thick in the \\ion{H}{1} Lyman continuum. This means that it is not necessarily correct to associate the observed physical extent of \\lya\\ emission with regions having ${\\rm N(HI) > 3 \\times 10^{18}}$ cm$^{-2}$ as R08 have suggested -- in principle, ${\\rm N(HI)}$ could be 1000 times lower and still remain optically thick to \\lya\\ photons. Perhaps more directly analogous to the results of the present sample is the narrow-band \\lya\\ survey of \\cite{hayashino04}. These authors used deep NB \\lya\\ images in the SSA22 field, and stacked the \\lya\\ images of 22 $z =3.09$ continuum-selected LBGs from the survey of \\cite{steidel03}, of which 19 are in common with our current SSA22 sample\\footnote{ The new NB image used in the present sample includes both archival Subaru data as well as an additional 10 hours' integration using LRIS on the Keck 1 telescope, and so is substantially deeper ($\\sim$ factor of 2-3), but covers a much smaller area, than that of \\cite{hayashino04}. }. Indeed, Hayashino et al showed that significant emission extends to angular scales of at least 4\\arcs\\, and that the ``ring'' in the range 2-4\\arcs\\ often contains as much or more \\lya\\ flux than the inner $\\theta \\le 2$ \\arcs\\ region. They also stated (but did not show) that a stack of the 13 galaxies which did not individually exhibit extended \\lya\\ emission results in a significant detection on the same 2-4\\arcs\\ scales. Although the authors did not discuss what physical mechanism might have been responsible for their observation, these results clearly provided an early indication of the nature of \\lya\\ emission in L* galaxies, borne out by our larger and more sensitive sample. \\subsection{Has the Whole Iceberg Been Detected?} The level of sensitivity to low-SB \\lya\\ emission at high redshifts is unlikely to improve by large factors using the current generation of ground-based telescopes, and so a natural question would be: How much more is there at still lower SB? Many \\lya\\ surveys (e.g., \\citealt{rauch08,bunker98}) have been designed to detect \\lya\\ fluorescence induced by the metagalactic radiation field at redshifts $2 \\simlt z \\simlt 3$. The radiation field intensity is usually expressed as $J_{\\nu} \\simeq 2-10 \\times 10^{-22}$ ergs s$^{-1}$ cm$^{-2}$ Hz$^{-1}$ sr$^{-1}$, where the quoted range indicates the dispersion among published observational or theoretical estimates (e.g., \\citealt{shapley06, bolton05,scott00,fauch08}). The expected maximum fluorescent signal at $z \\simeq 2.5-3.0$ is in the range $0.2-2 \\times 10^{-19}$ ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$ if the only source of ionizing photons is the general UV background (see e.g. \\citealt{cantalupo05,kollmeier10,fauch10}). These expectations clearly lie at or below the current SB thresholds of any survey completed to date. The difficulty of detecting the fluorescent signal from the metagalactic UV field has instead inspired several searches for fluorescence near bright sources of ionizing photons, such as QSOs (\\citealt{francis04,cantalupo05, adelberger06,hennawi09}). The results from such studies have been mixed. A different argument can be used to suggest that fluorescence from the UV background will always be overwhelmed by \\lya\\ scattering from the CGM of star-forming galaxies, at least at $z \\sim 2-3$. This assertion follows from the fact that S2010 found that the total absorption cross-section contributed by the CGM of LBGs (using ${\\rm R_{eff}} = 80-90$ kpc for the detection of low-ionization absorption species) can account for a large fraction of all gas with N(\\ion{H}{1})$\\simgt 2\\times10^{17}$ cm$^{-2}$ (i.e., $\\tau \\ge 1$ in the Lyman continuum, also known as ``Lyman Limit Systems''). In other words, any gas of sufficiently high $N(HI)$ to produce a detectable signal from fluorescence also lies within $\\sim 90$ kpc of a star-forming galaxy with properties similar to those in our sample. We have shown that these galaxies generically exhibit diffuse \\lya\\ emission on the same physical scales when a surface brightness threshold of $S(\\lya) \\sim 1 \\times 10^{-19}$ ergs s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$ is reached. Unless the fluorescent \\lya\\ signal lies at the very top of the allowed range, it will have much lower SB than the signal we have attributed to scattering from the inside of the galaxy out. It is more difficult to assess what fraction of observed \\lya\\ emission may be due to cooling processes such as those described by a number of recent authors (e.g., \\citealt{dijkstra09,kollmeier10,fauch10,goerdt10}.) In particular, the predictions of the emergent \\lya\\ emission from cooling gas accreting onto galaxies are extremely sensitive to gas temperature (\\citealt{kollmeier10,fauch10}) and to the small-scale structure in the gas. As a result, the range in \\lya\\ flux and SB, as well as the galaxy mass dependence and spatial distribution of cooling emission, must be regarded as uncertain by a factor of $\\simgt 10$, with an upper bound (based on energetic arguments) that can be as large as $L(\\lya) \\simgt 10^{44}$ ergs s$^{-1}$, but which under different assumptions could be as small as $\\sim 5 \\times 10^{41}$ ergs s$^{-1}$ for a galaxy with $M_{halo} \\simeq 9 \\times 10^{11}$ M$_{\\sun}$ (\\citealt{fauch10}), approximately the mean halo mass of the galaxies in the present sample (see \\citealt{asp+05,conroy08,steidel2010}). The observations appear to argue against a significant contribution of cooling radiation to the detected \\lya\\ halos, at least on average. We have shown that the shape of the observed radial surface brightness distribution among the LBGs in the sample is remarkably consistent beyond the inner $\\sim 10$ kpc, within which the \\lya\\ intensity for a given continuum luminosity varies by orders of magnitude. Moreover, the overall intensity scaling for the \\lya\\ emission at large radii is strongly correlated with the behavior of \\lya\\ emission in the inner 5-10 kpc region --- at the same continuum luminosity, \\lya\\ absorption-dominated galaxies (on average) exhibit diffuse \\lya\\ emission with a factor of 3-4 lower normalization than galaxies with spectroscopically detected \\lya\\ emission. In the context of \\lya\\ cooling radiation, one might expect the extended \\lya\\ emission to be strongly correlated with galaxy mass and/or SFR since it is believed by some (e.g., \\citealt{goerdt10}) that the baryonic accretion rate ultimately controls the SFR. In this scenario, the central region of \\lya\\ emission might be suppressed by higher \\ion{H}{1} column densities mixed with dust, but the outer regions would have no obvious way to ``know about'' the number of \\lya\\ photons being produced at smaller radii. Instead, one might expect that the brightest \\lya\\ halos would be associated with the ``\\lya\\ Abs'' sub-sample, since these have a median SFR nearly 3 (4.5) times larger than the ``\\lya\\ Em'' (LAE) sub-samples. Clearly, the observations are inconsistent with this expectation. If on the other hand most or all of the \\lya\\ emission at all radii originates in the central regions and is subsequently scattered by the CGM gas, the density of photons available for scattering at (for example) $r=50$ kpc will be very tightly linked to the number of \\lya\\ photons that successfully diffuse past $r \\sim 5$ kpc, beyond which the \\lya\\ halos appear ``self-similar''. The emergent \\lya\\ luminosities are entirely consistent with the observed level of star formation in the galaxies, and are more attenuated than the UV continuum, for all sub-samples except the LAEs. It is not necessary to invoke sources of \\lya\\ emission other than scattering (from the inside outward) to account for both the \\lya\\ luminosity and its spatial distribution. Under the scattering hypothesis, and further assuming that the scattering medium is self-similar for all galaxies, then sub-samples with more luminous \\lya\\ halos should provide information on the degree to which even the current SB threshold might lead to an underestimate of the total \\lya\\ flux emergent from a galaxy. To increase the dynamic range for detecting diffuse \\lya\\ emission, one might use the observed properties of giant LABs (which are well-detected in the stack to $b\\simeq 15$\\arcs) to estimate how much additional \\lya\\ flux may lie beyond the SB detection threshold near $b \\sim 8$\\arcs\\ for more typical galaxies. Under the assumption that diffuse emission from LABs and LBGs has a similar origin and differs only in total \\lya\\ luminosity, the curve-of-growth for LABs (Figure~\\ref{fig:lya_cum_plot}) suggests that an aperture of radius $\\simeq 8$ \\arcs\\ would underestimate the total \\lya\\ flux by only $\\sim 10$\\%. Thus, further aperture corrections to the integrated \\lya\\ would probably leave the values of $W_0(\\lya)$ (Table~\\ref{table:tab2}) and $f_{esc,rel}(\\lya)$ (Table~\\ref{table:tab3}) more or less unchanged. At least at $z\\simeq 2.65$, the current SB limit appears to be sufficient to detect most of the ``iceberg''. Finally, we note that the differences in the intensity of the large-scale diffuse emission among sub-samples divided according to their spectral morphology suggest that galaxy viewing angle is relatively unimportant (on average) for \\lya\\ emission; that is, most galaxies are not LAEs in some directions but strong \\lya\\ Abs systems in others, consistent with the inference of generally axisymmetric CGM gas distributions inferred from the absorption line studies (S2010). \\subsection{IS Absorption, \\lya\\ Emission, and the CGM} Perhaps the strongest correlation (first explored in detail by \\citealt{shapley03} for galaxies at $z \\sim 3$) among the observed spectral properties of LBGs is between the strength of low-ionization IS absorption lines and the spectral morphology and equivalent width of \\lya. Galaxies with the strongest \\lya\\ emission (among the continuum-selected samples) invariably have much {\\it weaker} than average low-ionization IS absorption lines (see \\citealt{erb2010} for a well-observed example), while those with the \\lya\\ appearing strongly in absorption have correspondingly strong IS absorption features, often reaching zero intensity over some or most of the line profile (see e.g. \\citealt{prs+02}) indicating unity covering fraction. These trends are easy to understand in the context of the CGM model discussed by S2010 and extended in this paper to cover the expectations for \\lya\\ scattering and its effects on the observability of \\lya\\ emission: both the IS absorption lines and \\lya\\ line strengths and morphologies are controlled by the kinematics and geometry of the same interstellar and circum-galactic gas. Dust certainly plays a role in determining the fraction of both \\lya\\ and continuum photons that will end up reaching an observer. However, the gas-phase geometry and kinematics are more directly responsible for the observed line strength (and line-to-continuum ratios) in the spectra. If a galaxy has strong \\lya\\ emission emerging from the same region as the UV continuum, it {\\it must} have shallow IS absorption lines; if it did not, then at least the spatial distribution of \\lya\\ (if not also its integrated flux) would be substantially modified-- it would become more spatially diffuse. When a slit spectrum (generally a small-aperture measurement) shows very strong and deep low-ionization IS absorption lines, including \\lya, it {\\it must} be the case that any \\lya\\ seen in emission will have escaped either from a region spatially distinct from the continuum (the subject of this paper), or by way of scattering from very high velocity material (see S2010). \\lya\\ emission seen in spectra which also show strong IS absorption will be primarily in the latter category, hence the nearly universal systemic redshift of \\lya\\ emission in LBG spectra. We have emphasized above that any \\lya\\ photons that are not destroyed by dust will eventually find their way out of their host galaxy-- but will be much harder to detect by the time they do. The point is that IS absorption and \\lya\\ emission are causally intertwined through their mutual dependence on the structure and kinematics of the CGM on scales from a few kpc to $\\simeq 100$ kpc." }, "1101/1101.2032_arXiv.txt": { "abstract": "{ \\Planck\\ has observed the entire sky from 30 GHz to \\getsymbol{HFI:center:frequency:857GHz:units}. The observed foreground emission contains contributions from different phases of the interstellar medium (ISM). We have separated the observed Galactic emission into the different gaseous components (atomic, molecular and ionised) in each of a number of Galactocentric rings. This technique provides the necessary information to study dust properties (emissivity, temperature, etc.), as well as other emission mechanisms as a function of Galactic radius. Templates are created for various Galactocentric radii using velocity information from atomic (neutral hydrogen) and molecular ($^{12}$CO) observations. The ionised template is assumed to be traced by free-free emission as observed by {\\it WMAP}, while 408 MHz emission is used to trace the synchrotron component. Gas emission not traced by the above templates, namely ``dark gas'', as evidenced using \\Planck\\ data, is included as an additional template, the first time such a component has been used in this way. These templates are then correlated with each of the \\Planck\\ frequency bands, as well as with higher frequency data from {\\it IRAS} and DIRBE along with radio data at 1.4\\,GHz. The emission per column density of the gas templates allows us to create distinct spectral energy distributions (SEDs) per Galactocentric ring and in each of the gaseous tracers from 1.4\\,GHz to 25 THz ($12\\micron$). The resulting SEDs allow us to explore the contribution of various emission mechanisms to the \\Planck\\ signal. Apart from the thermal dust and free-free emission, we have probed the Galaxy for anomalous (e.g., spinning) dust as well as synchrotron emission. We find the dust opacity in the solar neighbourhood, $\\tau/N_{\\rm H} = 0.92\\pm0.05\\times10^{-25} {\\rm cm}^2$ {\\changes at 250 $\\mu$m}, with no significant variation with Galactic radius, even though the dust temperature is seen to vary from over 25 K to under 14 K. Furthermore, we show that anomalous dust emission is present in the atomic, molecular and dark gas phases throughout the Galactic disk. Anomalous emission is not clearly detected in the ionised phase, as free-free emission is seen to dominate. {\\changes The derived dust propeties associated with the dark gas phase are derived but do not allow us to reveal the nature of this phase.} For all environments, the anomalous emission is consistent with rotation from polycyclic aromatic hydrocarbons (PAHs) and, according to our simple model, accounts for $(25\\pm5)\\%$ (statistical) of the total emission at 30 GHz. } ", "introduction": "\\label{sec:intro} In order to understand our own Galaxy, it is necessary to explore Galactic Plane ($|b|\\la10\\degr$), where we are able to observe the emission coming from a large range of distances. However the observed emission is the sum of a large number of line-of-sight components, often probing very different environments. Several previous studies \\citep{Bloemen1986, Bloemen1990, Giard1994, Sodroski1997, Paladini2007} have separated observed integrated emission into a number of Galactocentric radii in order to study its properties as a function of Galactic position and in different phases of the interstellar gas (e.g., atomic, molecular and ionised). The radial velocity of the gas is used to separate the Galactic gas emission into a number of Galactocentric rings and then the spectral energy distribution of each ring/gas phase is fitted with a physical model of dust and gas emissions. These methods have been used successfully in previous studies to map out basic properties of the ISM throughout the Galaxy. For instance, \\cite{Giard1994} demonstrated that polycyclic aromatic hydrocarbons (PAHs) are a ubiquitous component of the interstellar medium, and \\cite{Bloemen1990} showed that the dust temperature decreases with distance to the Galactic Centre in a way which is fully consistent with an exponential decrease of the interstellar radiation field (ISRF) in the stellar disk. \\cite{Sodroski1997} suggested that the abundance of large dust grains within each gas phase exhibits a gradient that is equivalent, within the uncertainties, to the metallicity gradient in the Galactic disk. \\cite{Paladini2007} showed that the dust in molecular clouds appears to be heated in a significant way by young massive stars still embedded in their parent clouds. With the advent of the \\Planck\\ satellite\\footnote{\\Planck\\ (http://www.esa.int/Planck) is a project of the European Space Agency (ESA) with instruments provided by two scientific consortia funded by ESA member states (in particular the lead countries: France and Italy) with contributions from NASA (USA), and telescope reflectors provided in a collaboration between ESA and a scientific consortium led and funded by Denmark.}, it is now possible to perform an inversion on the emission arising from the entire infrared, millimetre and centimetre range, making it possible to determine the dust and gas properties in many different environments in the Milky Way. The High Frequency Instrument (HFI) channels (100--857 GHz) allow us to properly constrain the big grain temperature and emissivity, while using the Low Frequency Instrument (LFI) channels (30--70 GHz), it is possible to constrain non-thermal emission mechanisms such as free-free, synchrotron or anomalous dust. We have used these data along with other ancillary data to perform a large-scale, low-resolution analysis of the Milky Way ISM emission. Knowledge of the dust emission and physical conditions throughout the Milky Way provides a self-consistent context for other \\Planck\\ studies: for example, the environmental effect on the properties of cold cores \\citep{planck2011-7.7a,planck2011-7.7b} can be evaluated, the modelling of individual anomalous dust regions can be placed in a Galaxy-wide context \\citep{planck2011-7.2}, or the Milky Way values can simply be compared to other galaxies \\citep{planck2011-6.4a,planck2011-6.4b}. In Sect.~\\ref{sec:data} we describe the data used in the study, including both those to be inverted and those used as templates to represent a particular phase of the ISM. In Sect.~\\ref{sec:method}, we describe how we optimise the separation of the gas observation into a number of Galactocentric rings, as well as how we report uncertainties on our best solution. We present the results in Sect.~\\ref{sec:results}, and discuss potential sources of uncertainty and bias in Sect.~\\ref{sec:discussion}. Finally we conclude in Sect.~\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} The analysis performed in this study provides a realistic description of the dust and gas properties as a function of Galactic radius in the Milky Way. The dust temperature in the \\ion{H}{i} is seen to decrease as a function of Galactocentric distance from 24.0 to 13.9\\,K, and the temperature in the molecular phase is heated by star formation associated to the Galactic spiral arms. The opacity found for grains in the solar circle is $\\tau/N_{\\rm H}=0.92\\pm0.05\\times10^{-25}\\, {\\rm cm}^2$, and with no significant variation with Galactic radius, even though the dust temperature is seen to drop by over 10\\,K from the centre to the outer Galaxy. The dust temperature in the \\ion{H}{i} gas in the solar circle, $17.6\\pm0.1$\\,K is also compatible with the recent \\Planck\\ value of high-latitude cirrus ($17.8\\pm0.9$\\,K). The extension of our analysis to lower frequencies than previously used in inversion techniques has allowed us to place constraints on the free-free, synchrotron and anomalous microwave emission: \\begin{itemize} \\item Anomalous dust emission is clearly seen in the atomic, molecular and dark gas phases. It is well fit by a very simple model consisting of spinning PAH molecules embedded in each of the gas phases. We highlight regions where spinning dust emission may be strong with respect to other emission mechanisms. According to our simple model, in the Galactic plane spinning dust accounts for $25\\pm5\\%$(statistical) of the total emission at 30 GHz. However, systematic uncertainties linked to our model may make this value more uncertain. \\item The dark gas phase has been explored spectrally across the Galactic plane for the first time. Its SED is similar to the molecular phase and therefore this phase seems to be tracing diffuse H$_2$ where CO is disassociated by energetic photons, as well as tracing optically thick CO and \\ion{H}{i}. \\item The free-free emission is completely traced by our ionised component and indicates that the free-free estimate from {\\it WMAP} component separation methods may be too high. In the ionised phase, spinning dust may be present but it is dominated by free-free emission. \\item Synchrotron emission is well characterised by a power law with a spectral index of $-1.0$. \\end{itemize} This first analysis of the large-scale Galactic emission as seen by \\Planck\\ has constrained the large scale ISM properties in the Galaxy. It could be improved in the future via several modifications. Individual regions could be studied using higher angular resolution data, which would probe smaller scale fluctuations that have been missed by the present study. The radial separation could be done differently to explore the difference between arm and inter-arm ISM, in order to focus more on the role of ISM conditions on the star formation process. It should also be possible to obtain an independent estimate of the dark gas in the plane, and ideally some radial information about its behaviour. In the near future, \\Planck\\ component separation products could be used in the place of {\\it WMAP} ones. Use of radio recombination line surveys currently underway would be very useful to constrain the free-free emission in the plane, and as a function of Galactic radius." }, "1101/1101.0037_arXiv.txt": { "abstract": "The {\\it Millimeter Sky Transparency Imager (MiSTI)} is a small millimeter-wave scanning telescope with a 25-cm diameter dish operating at 183~GHz. MiSTI is installed at Atacama, Chile, and it measures emission from atmospheric water vapor and its fluctuations to estimate atmospheric absorption in the millimeter to submillimeter. MiSTI observes the water vapor distribution at a spatial resolution of $0.5 \\arcdeg$, and it is sensitive enough to detect an excess path length of $\\lesssim 0.05$~mm for an integration time of 1~s. By comparing the MiSTI measurements with those by a 220~GHz tipper, we validate that the 183~GHz measurements of MiSTI are correct, down to the level of any residual systematic errors in the 220~GHz measurements. Since 2008, MiSTI has provided real-time (every 1~hr) monitoring of the all-sky opacity distribution and atmospheric transmission curves in the (sub)millimeter through the internet, allowing to know the (sub)millimeter sky conditions at Atacama. ", "introduction": "\\label{sect:intro} The major obstacle that leaves the millimeter (mm) to sub-millimeter (submm) windows still unexploited in ground-based astronomy is severe atmospheric absorption mainly caused by tropospheric water vapor. Moreover, small scale non-uniformity in its spatial distribution distorts wavefronts of astronomical signal, limiting the natural seeing at these wavelengths (e.g., \\cite{Carilli99}). At the ALMA \\citep{Wootten09} site, precipitable water vapor (PWV) contents of the atmosphere have been the subject of several site testing campaigns using devices such as tipping radiometers, Fourier-transform spectrometers (FTSs) and water vapor monitors (WVMs) operating at mm to submm wavelengths. These experiments succeeded in modeling atmospheric transparency at these wavelengths as a function of precipitable water vapor \\citep{Matsuo98, Matsushita99, Paine00, Paine04, Pardo01}. Another important fact from the viewpoint of telescope operation at mm and submm wavelengths is that observations above $\\sim 450$ GHz usually require very good weather conditions that are even at an excellent site like the Chajnantor plain in Chile. Dynamic scheduling is, therefore, necessary to make best use of a telescope under all conditions, and real-time monitoring of the atmospheric condition is indispensable to execute it. Tippers operating at 220--225~GHz are widely used at submm observatories (e.g., CSO, ASTE, APEX, etc.) for sounding the sky opacity and assessing astronomical data quality. An all-sky mid-infrared (IR) imager to monitor clouds on the sky, similar to those commonly used at optical/near-IR observatories \\citep{Takato03, Shamir05, Suganuma07, Sebag08, Miyata08}, is also working at the ASTE site (4860~m in elevation; Pampa la Bola, Atacama, Chile; \\cite{Ezawa04, Ezawa08} for the ASTE facility). There has been, however, no instrument to map the mm/submm emission of water vapor across the sky with high sensitivity. It will be important to monitor not only the amount of absorption by water vapor but also the amplitude of small-scale ($\\sim$10--100~m) fluctuations because they govern the astronomical seeing at mm and submm wavelengths (see \\citet{Altenhoff87} for single dish observations and \\citet{Thompson01} for interferometry). The turbulent layer of water vapor is known to be $\\approx$1~km above the ground level of the Atacama plateau \\citep{Delgado01, Robson02, Beaupuits05}. The {\\it Millimeter Sky Transparency Imager (MiSTI)} is a 25-cm diameter telescope operating at 183~GHz installed at the ASTE site. It is designed to measure the emission from the atmospheric water vapor and its time variation in any arbitrary azimuth ($Az$) and elevation ($El$) direction. A picture and schematic drawing of MiSTI are shown in figure~\\ref{oshima}. and monitoring, and using it to verify the quality of ALMA data (i.e., data quality assurance). We have developed MiSTI at the National Astronomical Observatory of Japan (NAOJ), and have installed and tested the instrument at the high-altitude ASTE site. The `first light' at the ASTE site was received in December 2007 (figure~\\ref{FirstLight}), and stand-alone automatic/remote operations were started in March 2008. Since May 2008, a dedicated web site for MiSTI\\footnote{\\texttt{http://aste-www.mtk.nao.ac.jp/\\symbol{\"7E}misti/opacity.html}} has been opened to provide the real-time 183~GHz sky images to radio astronomers who already have (and will have) telescopes in Atacama. In this paper, we present the instrumental overview of MiSTI and its initial results; more detailed analyses of the sky measurements will be presented in a subsequent paper. Requirements on the instrument and methods for measuring water vapor are introduced in Section~2. The system overview of MiSTI is described in Section~3. Application of the MiSTI as an all-sky opacity monitor is presented in Section~4. Sections~5 and 6 devote discussions and summary, respectively. ", "conclusions": "\\label{sect:discussions} The sky data obtained with the IR cloud monitor at the ASTE site can be used to examine the effect of the liquid component of water on the correlation between $\\tau_0^{\\mathrm{183}}$ and $\\tau_0^{\\mathrm{220}}$. The cloud monitor consists of a gold-plated Cassegrain-type mirror, a commercial mid-IR camera which has $320\\times 240$ image pixels, and a Linux based computer for control and data gathering. The cloud monitor can automatically take and store a fish-eye image of the sky every 5 minutes (figure~\\ref{irmon}). The camera is sensitive to the wavelength range of 8--12~$\\mu$m. Before installing the camera to the ASTE site, we measured its characteristics in the laboratory by taking images of an objects with various temperature ($T_{\\mathrm{obj}}$), and we confirmed that the output of each pixel is quite proportional to \\begin{equation} \\int _{8\\mu\\mathrm{m}}^{12\\mu\\mathrm{m}}B_{\\lambda}(T_{\\mathrm{obj}}) d\\lambda \\end{equation} in the range of $T_{\\mathrm{obj}} = 255$--330~K, where $B_{\\lambda}(T)$ is the Planck function. In each frame of the sky image, the IR emission from a black-painted aluminum block placed at the edge of the main mirror is also taken simultaneously (see figure~\\ref{irmon}). The temperature of the block is always monitored, and its emission is used as the brightness standard when the emissivity distribution of the sky is calculated. The obtained sky emissivity is always above 0; Even when there are no clouds, the emissivity is uniform and its typical value is 0.23--0.33, or 0.25--0.4 in optical depth at these IR wavelengths ($\\tau _{\\mathrm{IR}})$. This is mainly because there is an emission band of atmospheric ozone in this wavelength range. The measured accuracy is also limited by the accumulation of dirt on the mirror surface. More detailed description about the system can be found in the papers reported by \\citet{Takato03}, \\citet{Suganuma07}, and \\citet{Miyata08}. Figures~\\ref{compirmm}a and \\ref{compirmm}b show the comparison of optical depths at mm wavelengths and that at mid-IR during the period from April to August in 2008. These figures contain data that were taken only when all three environmental monitors at the ASTE site (i.e., MiSTI, tipper, and IR cloud monitor) were working. The time resolution of the data in figure~\\ref{compirmm} is smoothed to be 1 hour, which corresponds to 1.5 times the time resolution of MiSTI. When $\\tau_0^{\\mathrm{220}} < 0.05$ or $\\tau_0^{\\mathrm{183}} < 0.4$, $\\tau_{\\mathrm{IR}} = 0.25$--0.4, indicating that the sky was clear, while when the optical depths at mm wavelengths are large ($\\tau_0^{\\mathrm{220}} > 0.05$ or $\\tau_0^{\\mathrm{183}} > 0.4$), $\\tau_{\\mathrm{IR}}$ is large but also shows large scatter (figures~\\ref{compirmm}a and \\ref{compirmm}b). $\\tau_0^{\\mathrm{183}}$ changes by an order of magnitude even under clear sky conditions, suggesting that the observed frequency used by MiSTI can give us critical information to select the best conditions for submm observations. Figure~\\ref{compirmm}c presents comparisons between $\\tau_0^{\\mathrm{220}}$ and $\\tau_0^{\\mathrm{183}}$ under clear sky (i.e. $\\tau_{\\mathrm{IR}} \\leq 0.4$, dots) and cloudy conditions ($\\tau_{\\mathrm{IR}} > 0.4$, open circles). These comparisons reveal the followings: (i) $\\tau_0^{\\mathrm{220}}$ is always greater than 0.05 when the sky is cloudy, (ii) the correlation is much tighter under clear sky conditions, and (iii) $\\tau_0^{\\mathrm{220}}$ under cloudy sky shows large scatter and is even larger than the value expected from the correlation under clear sky. These results imply that the liquid component of water affects $\\tau_0^{\\mathrm{220}}$ more than $\\tau_0^{\\mathrm{183}}$, as discussed by \\citet{Matsushita00}. The correlation under clear sky shown in figure~\\ref{compirmm}c shows a bend at ($\\tau_0^{\\mathrm{183}}, \\tau_0^{\\mathrm{220}}$) = (0.4, 0.06) and the slope gets steeper in the regions of larger $\\tau_0^{\\mathrm{183}}$, which is also seen in figure~\\ref{tauplot}a. This may also be explained by more significant contribution of liquid water to $\\tau_0^{\\mathrm{220}}$. We have developed a small mm-wave telescope equipped with a 25-cm diameter dish, named the Millimeter Sky Transparency Imager (MiSTI). It measures emission from atmospheric water vapor at frequencies of $183.3 \\pm 4.5$~GHz, which maximize the sensitivity to tiny changes in the PWV content of the atmosphere. MiSTI observes the water vapor distribution at a spatial resolution of $0.5 \\arcdeg$, and it is sensitive enough to detect an excess path length of $\\lesssim 0.05$~mm for an integration time of 1~s. By comparing the measurements of MiSTI at 183~GHz with those from the 220~GHz tipper, we validate that the 183~GHz measurements of MiSTI are correct, down to the level of any residual systematic errors in the measurements at 220~GHz. When the mm opacities are large ($\\tau_0^\\mathrm{183} > 0.4$), optical depths measured at 220~GHz start to exceed those at 183~GHz. Given the comparison between the mm and IR opacities, this excess is likely attributed to water droplets in the sky, which contribute more largely to 220~GHz opacity than 183~GHz. Currently, MiSTI provides real-time (every 1~hr) monitoring of the all-sky opacity distribution and atmospheric transmission curves in the (sub)mm through the internet, allowing us to know the (sub)mm sky conditions at Atacama. It is important to investigate $\\sim$10--100~m scale structure of water vapor distribution because the small-scale fluctuations of the water vapor affect the natural seeing of submm single-dish telescopes and interferometers. We plan to update the functionality of MiSTI to further increase its sensitivity and measurement efficiency of such fluctuations. \\bigskip We are grateful to the referee for fruitful comments. We thank M.\\ Uehara, T.\\ Okuda, N.\\ Mizuno, and the ASTE team for the support in installation and operation of MiSTI. We also thank W.\\ Kimura, T.\\ Miyata and H.\\ Matsuo for contributions to the development of the IR cloud monitor. YT thank A.\\ Endo and M.\\ Kamikura for fruitful discussions. We acknowledge to the Advanced Technology Center (ATC) at NAOJ for allowing the use of the facilities. This work was financially supported by an ATC/NAOJ program for joint research and development and Grant-in-Aid for Scientific Research (A) (No.\\ 18204017). YT was financially supported by the Japan Society for the Promotion of Science (JSPS) for Young Scientists." }, "1101/1101.0347_arXiv.txt": { "abstract": "We study the active region NOAA 10960, which produces two flare events (B5.0, M8.9) on 04 June 2007. We find the observational signature of right handed helical twists in the loop system associated with this active region. The first B5.0 flare starts with the activation of helical twist showing $\\sim$3 turns. However, after $\\sim$20 minutes another helical twist (with $\\sim$2 turns) appears, which triggers M8.9 flare. Both helical structures were closely associated with a small positive polarity sunspot in the AR. We interpret these observations as evidence of kink instability, which triggers the recurrent solar flares. ", "introduction": "Solar flares are transient explosions in the solar atmosphere when the energy of stressed and twisted magnetic fields is released into heating and radiation. The flares associated with coronal mass ejections (CME) are known as ``eruptive flares\", while flares without CMEs are known as ``confined flares\". The large flares accompanied with energetic CMEs may be triggered by flux-rope eruption with significant changes in the photospheric fields [\\refcite{gary2004,liu2003}]. The flares can also be initiated due to the filament interactions followed by halo CMEs [\\refcite{kumar2010a} and references cited there] and filament eruptions [\\refcite{liu2008} and references cited there]. Observations show that the moderate flares without CME may be triggered by some instabilities (e.g., kink instability) [\\refcite{sri2010,kumar2010b}]. The instabilities may cause the destabilization of large-scale magnetic field, and can result CMEs [\\refcite{cho2009} and references cited there]. Instability of twisted magnetic flux tubes are well studied in theory and intense numerical simulations [\\refcite{torok2004,torok2005,hay2007}]. Kumar et al. [\\refcite{kumar2010b}] have recently presented a detailed multi-wavelength observations of the M8.9/3B class solar flare in the active region NOAA 10960 on 04 June 2007. They concluded that the positive flux emergence, the penumbral filament loss of the associated sunspot and the activation of the several twisted flux ropes in and around the flare site can be key candidates for the occurrence of this flare during 05:06 UT and 05:13 UT. The ``activation\" implies motion/brightening in the flux rope which is generated by some instability. A small B5.0 class flare has also been observed in the same active region during 04:40-04:51 UT, which seems to be a precursor for the M8.9 flare [\\refcite{sri2010}]. These two recurrent flares have been occurred during the activation of successive helical twist and kink unstable flux tubes from a positive polarity sunspot of AR 10960 [\\refcite{sri2010,kumar2010b}]. In this paper, we review the occurrence of recurrent solar flares and associated multi-wavelength phenomena. Multi-wavelength observations are described in section 2. The observational results are presented in section 3. The discussion and conclusions are given in the last section. \\begin{figure} \\centering \\psfig{file=goes_pk1.eps,width=10cm} $\\color{red} \\put(-235,140){\\vector(0,-1){15}}\\color{red} \\put(-244,142){B5.0}$ $\\color{red} \\put(-187,188){\\vector(-3,1){15}}\\color{red} \\put(-185,184){M8.9}$ \\caption{GOES Soft X-ray flux profiles in two different wavelength bands for both flares (indicated by arrows) on 4 June 2007.} \\label{fig1} \\end{figure} \\begin{figure} \\centering \\psfig{file=mdi_gband.eps,width=9cm} \\caption{SOHO/MDI magnetogram of the active region NOAA 10960 on 4 June 2007. The small positive polarity sunspot is shown in the box, indicated by an arrow. The enlarged view of the same sunspot is shown in Hinode SOT/G-band image in the bottom-right corner.} \\label{fig2} \\end{figure} ", "conclusions": "We study the recurrent flare activities in the active region NOAA AR 10960, which are accompanied by activations of helical twists over the small positive polarity $\\delta$ sunspot. Srivastava et al. [\\refcite{sri2010}] have observed the first activation of a highly (right-handed) twisted flux tube in AR 10960 during the period 04:43--04:52 UT. They have estimated the length and the radius of the loop as L$\\sim$80 Mm and a$\\sim$4.0 Mm respectively, and also estimated the total maximum twist angle as $\\sim$12$\\pi$ , by assuming quasi-symmetric distribution of the twist along the magnetic loop, which is much larger than the Kruskal--Shafranov instability criterion. They suggested that the right-handed twist is symmetrically distributed along the observed loop as a possible asymmetry can be smoothed over the short Alfv\\'en time of $\\sim$80 s. The detection of a clear double structure of the loop top during 04:47--04:51 UT in TRACE 171 \\AA \\ images are found to be consistent with simulated kink instability in curved coronal loops (T{\\\"o}r{\\\"o}k et al. [\\refcite{torok2004}]. They have suggested that the kink instability of this twisted magnetic loop triggered the B5.0 class solar flare, which occurred during 04:40 --04:51 UT in this active region. The co-spatial brightening in soft X-rays as observed by Hinode/XRT and the co-temporal occurrence of the right-handed twisting in the flux tube confirm the occurrence of the B5.0 flare during 04:40--04:51 UT probably due to the generation of the kink instability. Kumar et al. [\\refcite{kumar2010b}] have found multi-wavelength evidence of the successive activation of helical twists that may help in the energy build-up process at the flaring region in AR10960. The energy is released in the form of M-class flare after secondary activation of helical twist in the flux tube when it reconnects with neighboring opposite field. The activation of two helical structures/ropes played an important role in destabilizing and consecutive reconnection of magnetic field. The twist in the secondary magnetic structure crosses the threshold (2.5--3.5$\\pi$), which probably produces the kink instability in this structure. The energy release region in the M-class flare coincides with the twisted magnetic structures. The M-class flare showed agreement with the quadrupolar (closed--closed) reconnection model (breakout) between two closed field lines [\\refcite{anti1998}], which is evident in the decay phase of this flare (see bottom-right panel of Figure 5 also). The asymmetric evolution is driven by foot-point shearing of one side of an arcade, where reconnection between the sheared arcade and the neighboring (unsheared) flux system most probably triggers the flare. The kink instable twisted magnetic structure may undergo in a weak reconnection with the surrounding closed field lines in quadrupolar field configuration. Therefore, it triggers M-class flare in the active region without any eruption/CME [\\refcite{anti1999,aul2000,asc2004}]. Disappearance of Penumbrae during the decay phase and after the flare suggests that the magnetic field changes from inclined to almost vertical configuration [\\refcite{kumar2010b}]. This means that the part of penumbral magnetic field is converted into umbral fields. These results are in agreement with previous studies [\\refcite{wang2004,liu2005}]. The rotation of sunspot was the most plausible cause of the helical twist and thus energy release in B5.0 class flare [\\refcite{sri2010}]. On the other hand, the penumbral loss and umbral area enhancement were clearly evident during M-class flare [\\refcite{kumar2010b}]. Nobeyama 17 GHz radio contours overlaid the SoHO/MDI image during the M-class flare show the two radio sources, corresponding to the footpoints of magnetic loop system, which are generated during the impulsive phase of the M-class flare due to particle acceleration from the reconnection site. The bottom-right panel of Figure 5 shows the TRACE 171 \\AA\\ image in which the same bright loop system is clearly evident in the southward direction. The part of rising helical structure most likely reconnects with the southward loop-system and produce two radio sources due to particle acceleration from the reconnection site. These radio sources are the footpoints of the flaring loop system. The existence of co-spatial radio sources with this loop-system suggests the reconnection of twisted flux rope with the ambient field as most possible scenario for the flare triggering [\\refcite{kumar2010b}]. In conclusions, this paper reviews the rare observational signature of kink instability associated with failed eruptions and solar flares. Earlier, several researchers have been reported kink instability associated with CME eruptions [\\refcite{liu2008,cho2009} and references cited there]. The very interesting active region AR10960 was poor CME generator, but triggered many solar flares during its journey over the solar-disk. The multi-wavelength signature of helically twisted structures has been found as the cause of the recurrent solar flares on 04 June 2007. Therefore, we suggest that such flares may occur due to some instability/activation of twisted magnetic fields. The detailed multi-wavelength and statistical studies should be performed in future with the high- resolution space borne and ground-based observations." }, "1101/1101.0171_arXiv.txt": { "abstract": "Young solar-type stars rotate rapidly and are very magnetically active. The magnetic fields at their surfaces likely originate in their convective envelopes where convection and rotation can drive strong dynamo action. Here we explore simulations of global-scale stellar convection in rapidly rotating suns using the 3-D MHD anelastic spherical harmonic (ASH) code. The magnetic fields built in these dynamos are organized on global-scales into wreath-like structures that span the convection zone. We explore one case rotates five times faster than the Sun in detail. This dynamo simulation, called case~D5, has repeated quasi-cyclic reversals of global-scale polarity. We compare this case D5 to the broader family of simulations we have been able to explore and discuss how future simulations and observations can advance our understanding of stellar dynamos and magnetism. ", "introduction": "Magnetism is a ubiquitous feature of stars like our Sun. The magnetism we see at the surface probably has its origin in stellar dynamo action arising in the convective envelopes beneath the photosphere. There, turbulent plasma motions couple with rotation to build organized fields on global-scales. These processes occur in the Sun as well and are probably the source of the 11-year activity cycle. Despite intense study, solar and stellar dynamos are poorly understood, and at present we are unable to reliably predict even large-scale features of the solar cycle. Observations of young, rapidly rotating stars indicate that they have strong magnetic fields at their surfaces. There are clearly observed correlations between rotation and activity which appear to hold generally for stars on the lower main sequence \\citep[e.g.,][]{Pizzolato_et_al_2003}. Many of these stars show cycles of activity as well, though here the dependence on rotation rate, stellar mass and other fundamental parameters is less clear \\citep[e.g.,][]{Saar&Brandenburg_1999, Olah_et_al_2009}. At present even from a theoretical perspective we do not understand how the stellar dynamo process depends in detail on rotation. Motivated by this rich observational landscape, we have explored the effects of more rapid rotation on 3-D convection and dynamo action in simulations of stellar convection zones. These simulations have been conducted using the anelastic spherical harmonic (ASH) code to study global-scale magnetohydrodynamic convection and dynamo action in stellar convection zones \\citep[e.g.,][]{Clune_et_al_1999, Miesch_et_al_2000, Brun_et_al_2004}. In the past, global-scale convective dynamo simulations have focused primarily on the Sun, but now explorations are beginning for a variety of stars, ranging from A-type \\citep[e.g.,][]{Brun_et_al_2005, Featherstone_et_al_2009} to the M-type dwarfs \\citep{Browning_2008}. Here we will discuss simulations of G-type stars that rotate more rapidly than the Sun. We began these explorations by exploring convection in hydrodynamic simulations at a variety of rotation rates \\citep{Brown_et_al_2008}. These simulations capture the convection zone only, spanning from $0.72\\:R_\\odot$ to $0.97\\:R_\\odot$, and take solar values for luminosity and stratification but the rotation rate is more rapid. The total density contrast across such shells is about 25. In those simulations we found that the differential rotation generally becomes stronger as the rotation rate increases, while the meridional circulations appear to become weaker and multi-celled in both radius and latitude. These rapidly rotating stars have vigorous dynamos, and the magnetic fields created in the dynamos are often organized on global-scales into banded wreath-like structures \\citep{Brown_et_al_2010a}. Surprisingly, this organization occurs in the middle of the convection zone itself, rather than in a tachocline of penetration and shear between the convection zone and stable radiative zone beneath. Many of the wreath-building undergo quasi-cyclic reversals of magnetic polarity. Here we explore one of these cyclic dynamos (\\S\\ref{sec:case D5}), before putting it in context with other such dynamos (\\S\\ref{sec:parameter space}). ", "conclusions": "" }, "1101/1101.0940_arXiv.txt": { "abstract": "{We present photometric observations of the dwarf nova 1RXS J053234.9+624755. We performed a detailed analysis of the superoutburst that occurred in August 2009. We found the superhump period to be $P_{sh}=0.057122(14)$days. Based on the $O-C$ diagram we conclude that $P_{sh}$ increased during the plateau at the rate of $dP_{sh}/dt=(9.24 \\pm 1.4) \\cdot 10^{-5}$. Both the $O-C$ analysis and evolution of the superhumps light curve favour the model in which superhumps originate in a variable source located in the vicinity of the hot spot. In addition, the evolution of the light curve suggests that the superhump light source approaches the disc plane as the superoutburst declines. Detailed analysis of the superoutburst plateau phase enabled us to detect a signal which we interpret as apsidal motion of the accretion disc. We detected additional modulations during the final stage of the superoutburst characterized by periods of 104s and 188s which we tentatively interpret as quasi periodic oscillations. Estimations of $A_{0}$ and $A_{n}$ are in agreement with the dependence discovered by Smak (2010) between the amplitude of superhumps and the orbital inclination. } {accretion, accretion discs - binaries: cataclysmic variables, stars: dwarf novae, oscillations, stars: individual: 1RXS J053234.9+624755, 2MASS J05323386+6247 520,USNO-B1.0 1527-00176070, USNO-B1.0 1527-00176070} ", "introduction": "Over 30 years have passed since the discovery of superhumps in cataclysmic variables (Vogt 1974, Warner 1995). Different physical processes have been proposed to explain the origin of the superhump phenomenon. Models considered have included the ejection of matter from the white dwarf due to pulsation instabilities (Vogt 1974), periodic modulation of dissipation in the elliptical and precessing accretion disc (Osaki 1989), or models which explain superhumps in terms of the oscillations of a hot spot caused by uneven stream flows from the secondary (Smak 2009c). Nowadays we posses enough knowledge about the components of dwarf novae to describe these systems completely. It seems likely that under the strength of argument provided by Smak (2010), most doubts regarding superhumps will be dispelled soon. However, each new model should be tested by observations and therefore it is important to observe dwarf novae especially during superoutbursts, i.e. when superhumps occur predominantly. 1RXS J053234.9+624755 (1RXS J0532) is an example of a dwarf nova with documented superoutbursts. It was discovered quite unexpectedly as a counterpart of an X-ray source in the ROSAT all-sky bright catalog by Bernhard et al. (2005). Detailed investigation of the Sonneberg Plate Archive for this object found eight outbursts between 1990 and 2005 with a mean interval of 133.6d. Soon, further observations resulted in the detection of superhumps in the light curve of 1RXS J0532 and revealed that this star is a SU UMa type dwarf nova (Poyner \\& Shears 2006). Monitoring of superoutbursts showed the evolution from a clear tooth-shape light curve variation to more random flickering (Parimucha \\& Dubovsky 2006). Based on spectroscopic observations an orbital period of 0.05620(4) d was reported by Kapusta \\& Thorstensen (2006). An observational campaign during the 2005 superoutburst reported by Imada et al. (2009) showed the superoutburst accompanied by a precursor. The light curve of the precursor revealed a gradual increase in the amplitude of the light variation which was interpreted as developing superhumps. As this is in contradiction to the standard Thermal - Tidal Instability (TTI) model, further interpretation was based on the \"refined\" TTI model. The authors concluded that the existence of the precursor and the presence of growing superhumps during the precursor are related to the mass ratio of the system. This in turn determines the 3:1 resonance radius and the tidal truncation radius ratio which are supposed to be responsible for the behavior of the superhumps (Osaki 2005). Moreover, the photometric data provided an estimate of the length of the supercycle, $\\sim$450 days, and the superhump period, $P_{sh}=0.57169(6)$d (Imada et al. 2009). ", "conclusions": "\\begin{itemize} \\item[a.] It has been confirmed that 1RXS J053234+624755 is a SU UMa star. The measured superhump period is $P_{sh}=0.057122(14)$d. The superoutburst is characterized by a slight rebrightening in the later phase of the plateau. The amplitude of the superoutburst was determined to be 3.3 mag. \\item[b.] We determined the rate of superhump period change to be $\\dot{P}=9.5 \\cdot 10^{-5}$. This value is noticeably different than the rates obtained for for superoutbursts in 2005 and 2008 ($\\dot{P}$=$5.7 \\cdot 10^{-5}$ and $10.2 \\cdot 10^{-5}$, Imada et al. 2009 and Kato et al. 2009, respectively). \\item[c.] We proposed that the observed superhump light curve behavior can be explained if we assume that the distance of the superhump source from the disc plane decreases as the superoutburst declines. \\item[d.] Additional data analysis allowed us to detect a quasi-periodic signal at the frequency $f_1=460.34\\pm1.63$c/d (188s). This frequency, however, may be spurious and simply the result of high humidity during the observation. \\item[e.] Detailed analysis of the superoutburst plateau phase enabled us to detect oscillations with a period $P_{prec}/2=1.699\\pm0.005$d, which we interpret as the effect of an apsidal motion of the accretion disc. \\item[f.] Based on the double peaks in the emission lines, Kapusta \\& Thorstensen (2006) concluded that the orbital inclination, $i$, of 1RXS J0532 is not far from edge on. Given the prominence in the spectroscopic data of the double peaks in the emission lines, it is almost certain that $i$ is larger than $70$ deg. Taking into account the fact that there are no observed disc eclipses, we have assumed $i=75$ deg. Even though we were not able to observe this system during the maximum of the superoutburst, visual inspection of the amplitude of superhumps in Figure 4 in Poyner \\& Shears (2006) and Figure 3 in Imada et al. (2009) gave $A_{0}=0.28\\pm0.01$ mag. Using eq. (4) and (5) from Smak (2010) we obtained $A_{n}=0.151$ mag. These results are consistent with the dependence between superhump amplitude and orbital inclination discovered by Smak (2010). \\end{itemize} \\Acknow{This work was partly supported by the Polish MNiSW grant no. N203~301~335. Artur Rutkowski has been supported by 2221-Visiting Scientist Fellowship Program of TUBITAK. AR is also grateful to Prof. Zeki Eker for inspiring discussions. Skinakas Observatory is a collaborative project of the University of Crete and the Foundation for Research and Technology-Hellas. We would like to thank the Skinakas Observatory staff for their help and assistance. }" }, "1101/1101.2174_arXiv.txt": { "abstract": "{ Close to the Planck energy scale, the quantum nature of space-time reveals itself and all forces, including gravity, should be unified so that all interactions correspond to just one underlying symmetry. In the absence of a full quantum gravity theory, one may follow an effective approach and consider space-time as the product of a four-dimensional continuum compact Riemanian manifold by a tiny discrete finite noncommutative space. Since all available data are of a spectral nature, one may argue that it is more appropriate to apply the spectral action principle in this almost commutative space. Following this procedure one obtains an elegant geometric explanation for the most successful particle physics model, namely the standard model (and supersymmetric extensions) of electroweak and strong interactions in all its details, as determined by experimental data. Moreover, since this gravitational theory lives by construction at very high energy scales, it offers a perfect framework to address some of the early universe cosmological questions still awaiting for an answer. After introducing some of the main mathematical elements of noncommutative spectral geometry, I will discuss various cosmological and phenomenological consequences of this theory, focusing in particular on constraints imposed on the gravitational sector of the theory. } \\FullConference{Corfu Summer Institute on Elementary Particles and Physics - Workshop on Non Commutative Field Theory and Gravity,\\\\ September 8-12, 2010\\\\ Corfu Greece} \\begin{document} ", "introduction": "At energies much below the Planck scale, gravity can be considered as a classical theory, however as energies approach the Planck scale, the quantum nature of space-time becomes apparent, and the simple prescription, dictating that physics can be described by the sum of the Einstein-Hilbert and the Standard Model (SM) action ceases to be valid. At such high energy scales, all forces, including gravity, are expected to be unified so that all interactions correspond to one underlying symmetry. Thus, near Planckian energies, the appropriate formulation of geometry should be within a quantum framework, while the nature of space-time would change in such a way so that one can recover the low energy picture of diffeomorphism and internal gauge symmetries, which govern General Relativity (GR) and gauge groups on which the Standard Model is based, respectively. A promising attempt to obtain a quantum nature of space-time has been realised within the realm of NonCommutative Geometry (NCG). Noncommutative Geometry~\\cite{ncg-book1, ncg-book2} is a beautiful and rich mathematical theory, according which geometry can be described through the functions defined on the geometry, while the geometric properties of spaces can be described by the properties of functions defined on the spaces. An important new feature of noncommutative geometry is the existence of inner fluctuations of the metric, which correspond to the subgroup of inner automorphisms. Besides the mathematical beauty of NCG, which by itself explains why one may want to study this theory, NCG offers a variety of phenomenological consequences, which turn this theory into a fertile framework to address fundamental issues of early universe cosmology and high energy physics phenomenology. In what follows we will follow Connes' approach~\\cite{ncg-book1, ncg-book2} and consider a model of a two-sheeted space made from the product of a continuous space by a discrete space. This model led to a geometric explanation of the Standard Model; in particular the model shows that the vacuum Expectation Value of the Higgs field is related to the noncommutative distance between the two sheets. Within Connes' model of NCG, the Higgs field is conformally coupled to the Ricci curvature, while the generalised Einstein-Hilbert action contains in addition a minimally coupled massless scalar field related to the distance between the two sheets. Connes' approach is based upon a spectral action principle, stating that the bare bosonic Euclidean action for any noncommutative model based on a product (noncommutative) space is the trace of the heat kernel associated with the square of the noncommutative Dirac operator of the product geometry. Within noncommutative spectral geometry, we look for a hidden structure in the functional of gravity coupled to the SM at today's low energy scales, and avoid an extrapolation by many orders of magnitude to guess the appropriate structure of space-time at Planckian energy scales. Noncommutative spectral geometry offers an elegant approach to unification, based on the symplectic unitary group in Hilbert space, rather than on finite dimensional Lie groups. The model offers a unification of internal symmetries with the gravitational ones. All symmetries arise as automorphisms of the noncommutative algebra of coordinates on a product geometry. Due to the lack of a full quantum gravity theory, which {\\sl a priori} should define the geometry of space-time at Planckian energy scales, we will follow an effective theory approach and consider the simplest case beyond commutative spaces. Thus, below but close to the Planck energy scale, space-time will be considered as the product of a Riemanian spin manifold by a finite noncommutative space. At higher energy scales, space-time should become noncommutative in a nontrivial way, while at energies above the Planck scale the whole concept of geometry may altogether become meaningless. As a next, but highly nontrivial, step one should consider noncommutative spaces whose limit is the almost commutative space considered here. Unfortunately, the birth of geometry may remain an unsolved puzzle for still quite sometime. Let me draw the attention of the reader to the fact that the noncommutative spectral geometry approach discussed here, goes beyond the noncommutative geometry notion employed in the literature to implement the fuzziness of space-time by means of $[{\\bf x}^i, {\\bf x}^j]=i\\theta^{ij}$, where $\\theta^{ij}$ is an anti-symmetric, real, $d\\times d$ ($d$ is the dimension of space-time) matrix, and ${\\bf x}^i$ denote spatial coordinates. ", "conclusions": "Noncomutative spectral geometry provides an elegant way of expressing the full Standard Model of strong and electroweak interactions coupled to Einstein gravity, as pure gravity on a modified space-time geometry. The paradigm of metric noncommutative geometry studied here is of a spectral nature, an important notion in physics since all experimental data are indeed of a spectral type. This approach is fundamentally different than any other paradigm, which imposes a particular structure for geometrical spaces in the quantum gravity regime. Connes' model focuses on an almost commutative space, considering that at energies close but lower than Planckian energy scales, space can be described by the tensor product of a continuum manifold by a discrete space. Within the context of noncommutative spectral geometry, gravity and matter are treated in a similar way, leading to concrete relationships between matter and gravitational couplings. The asymptotic expansion of the gravitational sector of the theory leads to modifications to General Relativity which can be used to constrain the theory through astrophysical observations. Considering the energy lost by binary systems to gravitational radiation, we were able to restrict the value of the Weyl squared coupling in the bosonic action. Investigating Higgs driven inflation within noncomutative spectral action, we have shown that while the Higgs potential can lead to the slow-roll conditions being satisfied once the running of the self-coupling at two-loops is included, the constraints imposed from the CMB data make the predictions of such a scenario incompatible with the measured value of the top quark mass. Another massless scalar field, which naturally appears in the model, seems also not to lead to a successful era of slow-roll inflation. However, the arbitrary mass scale $\\Lambda$ in the spectral action for the Dirac operator can be made dynamical by introducing a dilaton field; this dilaton field may turn out to be a successful inflaton candidate. Noncommutative spectral geometry faces, to my opinion, at least two immediate research directions, essential in order to deduce further cosmological and phenomenological consequences of this paradigm. Firstly, one should compute higher order terms in the asymptotic expansion of the spectral action functional. Note that it is very difficult to compute exactly the spectral action in its nonperturbative form, even though some progress has been made however recently~\\cite{nonpert}. Since the action functional ${\\rm Tr}(f ({\\cal D}/\\Lambda))$ is not local -- its locality is only achieved when it is replaced by the asymptotic expansion -- at least the next term in the asymptotic expansion must be computed, in order to check the validity of the asymptotic expansion. It was recently shown~\\cite{nonpert} that for a space-time whose spatial sections are 3-spheres $S^3$, Wick rotated and compactified to a Euclidean model $S^3 \\times S^1$, the spectral action is given, for any test function, by the sum of two terms up to a remarkably tiny correction. Let me emphasise that for any low-energy astrophysical consequence of the noncommutative spectral geometry, {\\sl a priori} the full spectral action, and not only its asymptotic form, has to be considered. Secondly, it could be of great importance to find the running of the parameters appearing in the spectral action, since otherwise it is impossible to extract information for low-energy astrophysical events. It is worth repeating that the expressions for $\\kappa_0, \\alpha_0, \\gamma_0, \\tau_0, \\mu_0, \\lambda_0$ in terms of $f_0, f_2, f_4, \\mathfrak{a}, \\mathfrak{b}, \\mathfrak{c}, \\mathfrak{d}, \\mathfrak{e}$ and the conformal value for $\\xi_0$ are only valid at unification scale $\\Lambda$. It is simply incorrect to naively postulate that these equalities can hold at lower energy scales as such, by just considering the parameters $\\kappa, \\alpha, \\cdots$ as functions of the energy scale. At last but not least, one should consider less trivial noncommutative spaces whose limit is the almost commutative space considered in the original Connes' model discussed here. Nevertheless, besides these necessary further developments, it is fair to conclude by stating that noncommutative spectral geometry offers a beautiful mathematical construction which provides an elegant explanation for the most successful particle physics model at hand. \\vskip1.truecm It is a pleasure to thank the organisers of the Workshop on Non Commutative Field Theory and Gravity, held in the beautiful island of Corfu, for inviting me to present this work during a stimulating and interesting meeting. \\vskip1.truecm" }, "1101/1101.4943_arXiv.txt": { "abstract": "A new serendipitous XMM survey in the area of the Sloan Digital Sky Survey is described (XMM/SDSS), which includes features such as the merging of overlapping fields to increase the sensitivity to faint sources, the use of a new parametrisation of the XMM point spread function for the source detection and photometry, the accurate estimation of the survey sensitivity. About 40\\,000 X-ray point sources are detected over a total area of $\\rm 122\\,deg^2$. A subsample of 209 sources detected in the 2-8\\,keV spectral band with SDSS spectroscopic redshifts in the range $0.0341.5$ (erg/s) are selected to explore their distribution on the colour magnitude diagram. This is compared with the colour-magnitude diagram of X-ray AGN in the AEGIS field at $z\\approx0.8$. We find no evidence for evolution of the rest-frame colours of X-ray AGN hosts from $z=0.8$ to $z=0.1$. This suggests that the dominant accretion mode of the AGN population, which is expected to imprint on the properties of their host galaxies, does not change since $z=0.8$. This argues against scenarios which attribute the rapid decline of the accretion power of the Universe with time (1\\,dex since $z=0.8$) to changes in the AGN fueling/triggering mode. ", "introduction": "Understanding the evolution of Active Galactic Nuclei (AGN), which signpost accretion events onto Supermassive Black Holes (SMBHs), remains a challenge for modern astrophysics. Although observations have demonstrated beyond any doubt that the luminosity density of these systems has dropped by more than 1 order of magnitude from $z\\approx1$ to the present day \\citep[e.g.][]{Ueda2003, Hasinger2005, Ebrero2009, Aird2010}, the physical mechanisms that drive this rapid decline are still not well constrained. The lack of a physical description for the cosmological evolution of SMBHs has implications that go beyond the AGN community. Recent evidence indicates an intimate relation between the building of galaxies and the growth of the SMBH at their centres \\citep[e.g.][]{Ferrarese2000, Gebhardt2000}. Therefore without a better understanding of AGN evolution our picture for the buildup of galaxies will also be incomplete. Galaxy mergers have long been proposed as the mechanism that triggers AGN and drives their cosmological evolution. Numerical SPH (Smoothed Particle Hydrodynamic) simulations demonstrate that these violent events are very efficient in funneling gas to the nuclear galaxy regions \\citep[e.g.][]{Hernquist1989,Barnes1991,Barnes1996}, where it can be consumed by the SMBH \\citep{Springel2005,DiMatteo2005}. Consequently in most semi-analytic cosmological simulations of galaxy formation \\citep[e.g.][]{Cattaneo2005, Somerville2008, Wang2008} the merging of near equal mass gas rich galaxies is the primary mechanism for growing SMBHs. In this family of models the evolution of AGN is intimately related to the decline of the fraction of gaseous major mergers with redshift \\citep[e.g.][]{LopezSanjuan2009}. Alternative scenarios propose that the dominant mode of accretion onto SMBHs changes with time \\citep{Hasinger2008, Fanidakis2010}, thereby leading to the observed decline of the AGN space density since $z\\approx1-2$. It is suggested for example, that SMBHs form primarily in violent gaseous major merger events at high redshift \\citep[$z\\gtrsim1$; ``QSO-mode''][]{Hopkins2006}, while stochastic accretion (e.g. internal instabilities, minor interactions) dominates the growth of SMBHs \\citep[``Seyfert-mode''][]{Hopkins_Hernquist2006} at lower redshift ($z<0.5$) and produces, on average, lower luminosity systems. Another possibility proposed by \\cite{Fanidakis2010} is that the dominant AGN fueling mode shifts from disk instabilities at high redshift to halo gas accretion \\citep[also termed ``radio'' mode,][]{Croton2006} at low redshift. One of the testable predictions of those models is that the properties of AGN hosts, such as the the star-formation history and the morphology, change with redshift, from $z\\gtrsim1$ to $z<0.5$. In contrast, in the merger only driven evolution scenario, AGN should be hosted by galaxies that have similar properties at all redshifts. The study of AGN hosts as a function of redshift requires a homogeneously selected AGN sample over a wide redshift baseline, $z\\approx0$ to $z\\approx1$ and beyond. Only then can one compare directly the properties of AGN hosts at different epochs. Variations in the selection function with redshift are hard to account for and may lead to erroneous conclusions. At $z\\ga1$ in particular, X-ray observations, especially at energies $>2$\\,keV, are one of the most efficient and least biased methods for locating AGN with a selection function that is easy to quantify. The infrared is a promising wavelength regime for finding active SMBHs, although there are issues related to contamination of infrared selected AGN samples by starbursts \\citep[e.g.][]{Georgantopoulos2008, Donley2008, Pope2008, Georgakakis2010}. Optical spectroscopy is also a powerful tool for identifying AGN, but aperture effects, which are particularly severe at high redshifts, raise concerns about dilution of the AGN emission lines by the host galaxy stellar population \\citep[e.g.][]{Severgnini2003}. As a result, surveys with XMM and Chandra, both deep/pencil-beam and shallow/wide, have been the workhorse of the astronomy community for compiling AGN samples. Those surveys however, essentially probe active SMBHs close to the peak of the accretion power of the Universe, $z\\ga 1$, and lack the area to provide meaningful constraints on the AGN population at $z\\la 0.5$. The Sloan Digital Sky Survey \\citep[SDSS;][]{Abazajian2009} has identified the largest sample of low redshift ($z\\approx0.1$) AGN todate, using diagnostic emission line ratios \\citep{Kauffmann2004}. The selection function of that sample however, is very different from that of X-ray AGN in deep surveys, rendering the comparison difficult. For example, the SDSS AGN include a large number of LINERs (Low Ionisation Nuclear Emission line Region), which are controversial objects and may not be powered by accretion onto a SMBH \\citep[e.g.][]{Sarzi2010}. More relevant to high redshift X-ray AGN surveys are the serendipitous near all-sky AGN samples compiled in the nearby Universe by the high energy missions INTEGRAL \\citep[20-100\\,keV, total of 144 Seyfert AGN,][]{Beckmann2009} and SWIFT \\citep[15-195\\,keV, total of 266 Seyferts,][]{Tueller2010}. Although those AGN samples are selected at much higher rest-frame energies compared to X-ray sources detected by Chandra or XMM at $z\\approx1$ ($<15$\\,keV), they are more appropriate than SDSS AGN as a low redshift comparison sample. Ideally however, one would like to select nearby AGN at rest-frame X-ray energies comparable to those of Chandra/XMM surveys at $z\\approx1$. In this paper we describe the compilation of a large sample of X-ray selected AGN at $z\\approx0.1$ using a new serendipitous X-ray survey (hereafter referred to as XMM/SDSS) in the area of the SDSS based on archival XMM observations. The advantage of the low redshift X-ray AGN subset of the XMM/SDSS survey is that the selection function is almost identical to deep pencil-beam samples, thereby minimising differential selection biases. As a result the comparison between the AGN host galaxy properties between $z\\approx0.1$ and $z\\approx1$ is greatly facilitated. The XMM archive includes over 10 years worth of observations, which allow serendipitous surveys over many tens of square degrees on the sky. Therefore the volume probed by the XMM/SDSS sample at $z\\approx0.1$ is orders of magnitude larger than any wide-area survey carried out by either XMM or Chandra, thereby allowing detailed studies of the statistical properties of low redshift X-ray AGN. The SDSS is the field of choice for our serendipitous X-ray survey as it provides the essential follow-up observations for the identification of low redshift X-ray AGN. The SDSS optical data include 5-band photometry ($ugriz$) and extensive optical spectroscopy for galaxies at $z\\approx0.1$. Wide area surveys at various wavelengths in the SDSS are either underway or completed, thereby allowing panchromatic studies of the XMM/SDSS X-ray AGN over a wide redshift baseline. The ancillary data in the SDSS include (i) the UKIRT Infrared Deep Sky Survey \\citep[UKIDSS,][]{Lawrence2007}, which will cover $\\rm 4000 \\, deg^2$ of the SDSS in 4 near-infrared bands to $K=18.5$\\,mag, (ii) the AKARI all-sky survey in 6 mid- and far-infrared bands from $\\rm 9-200\\mu m$ \\citep{Ishihara2010}, (iii) the Herschel ATLAS far-infrared survey \\citep{Eales2010} and (iv) the FIRST survey, which provides deep (1mJy) radio continuum (1.4GHz) images of the entire SDSS \\citep{Becker1995}. The structure of the paper is as follows. In sections 2 to 9 we describe the automated pipeline developed to reduce the archival XMM observations that overlap with the SDSS. The data reduction steps include the construction of event files and images in different energy bands, the detection of sources, the calculation and application of astrometric corrections, the estimation of fluxes and the identification of X-ray sources with optical counterparts. The pipeline uses a new parametrisation of the XMM Point Spread Function (Appendix A), which is employed for source detection and photometry, and an accurate method for estimating the sensitivity of the X-ray survey to point sources. Section 10 demonstrates the utility of our serendipitous X-ray AGN sample by comparing, for the first time, the colour magnitude diagrams of X-ray selected AGN at $z\\approx0.1$ and $z\\approx0.8$. Section 11 discuss the results and conclusions. Throughout this paper we adopt $\\rm H_{0} = 100 \\, km \\, s^{-1} \\, Mpc^{-1}$, $\\rm \\Omega_{M} = 0.3$ and $\\rm \\Omega_{\\Lambda} = 0.7$. Rest frame quantities (e.g. absolute magnitudes, luminosities) are parametrised by $h=H_{0} / 100$. ", "conclusions": "We present a new serendipitous X-ray survey, the XMM/SDSS, which is based on archival XMM observations and covers $\\rm 122\\,deg^2$ in the SDSS DR7 footprint. The size and sensitivity of this survey are well suited for low redshift X-ray AGN studies. The XMM/SDSS source catalogue is combined with the SDSS optical spectroscopy to compile one of the largest hard X-ray (2-8\\,keV) AGN samples to date in the redshift interval $0.032$, the He abundance suggested $n=2$, the uniformity of {\\rm HI} suggested $n\\ge2$, and the cluster mass function gave $n=1$ to 2. Schmidt also suggested that with $n=2$, dense galaxies like ellipticals should now have less gas than low-density galaxies like the LMC. His final comment was ``It is hoped to study the evolution of galaxies in more detail in the future.'' Following \\cite{schmidt59}, many authors derived scaling relations between the average surface density of star formation, $\\Sigma_{\\rm SFR}$, and the average surface density of gas. \\cite{buat89} included molecular and atomic gas and determined star formation rates from the UV flux corrected for Milky Way and internal extinction. They assumed a constant ${\\rm H}_{\\rm 2}/{\\rm CO}$ ratio and a \\cite{scalo86} IMF. The result was a good correlation between the average star formation rate in a sample of 28 galaxies and the $1.65\\pm0.16$ power of the average total gas surface density. In the same year, \\cite{ken89} used H$\\alpha$ for star formation, and {\\rm HI} and {\\rm CO} for the gas with a constant ${\\rm H}_{\\rm 2}/{\\rm CO}$ conversion factor, and determined star formation rates both as a function of galactocentric radius and averaged over whole galaxy disks. For whole galaxies, the average H$\\alpha$ flux scaled with the average gas surface density to a power between 1 and 2; there was a lot of scatter in this relation and the correlation was better for {\\rm HI} than ${\\rm H}_{\\rm 2}$. More interesting was Kennicutt's (1989) result that the star formation rate had an abrupt cutoff in radius where the \\cite{toomre64} stability condition indicated the onset of gravitationally stable gas. Kennicutt derived a threshold gas column density for star formation, $\\Sigma_{\\rm crit}=\\alpha\\sigma\\kappa/(3.36G)$ for $\\alpha=0.7$; $\\sigma$ is the velocity dispersion of the gas; $\\kappa$ is the epicyclic frequency, and $G$ is the gravitational constant. In a second study, \\cite{ken98} examined the disk-average star formation rates using a larger sample of galaxies with H$\\alpha$, {\\rm HI}, and {\\rm CO}. He found that for normal galaxies, the slope of the SFR-surface density relation ranged between 1.3 to 2.5, depending on how the slope was measured; there was a lot of scatter. When starburst galaxies with molecular surface densities in excess of 100 $M_\\odot$ were included, the overall slope became better defined and was around 1.4. This paper also found a good correlation with a star formation rate that scaled directly with the average surface density of gas and inversely with the rotation period of the disk. This second law suggested that large-scale dynamical processes are involved. \\cite{hunter98} considered the same type of analysis for dwarf Irregulars and derived a critical surface density that was lower than the \\cite{ken89} value by a factor of $\\sim2$. This meant that stars form in more stable gas in dwarf irregulars compared to spirals. \\cite{boissier03} compared $\\Sigma_{\\rm SFR}$ and $\\Sigma_{\\rm gas}$ versus radius in 16 resolved galaxies with three theoretical expressions. The best fits were a SFR dependence on the gas surface density as $\\Sigma_{\\rm SFR}\\propto\\Sigma_{\\rm gas}^{2.06}$, a more dynamical law from \\cite{bp99} which gave the fit $\\Sigma_{\\rm SFR}\\propto\\Sigma_{\\rm gas}^{1.48}(V/R)$ for rotation speed $V$ and radius $R$, and a third type of law from \\cite{dopita94}, which fit to $\\Sigma_{\\rm SFR}\\propto\\Sigma_{\\rm gas}^{0.97}/\\Sigma_{\\rm tot}^{0.61}$. \\cite{boissier03} assumed that H$_{\\rm 2}$/{\\rm CO} varied with radius as the metallicity \\citep{boselli02}. Their conclusion was that the three laws are equally good, and that for the pure gas law, $n>1.4$. \\cite{boissier03} also looked for a star formation threshold in the Milky Way. They determined $\\Sigma/\\Sigma_{\\rm crit}$ using both pure-gas for $\\Sigma_{\\rm crit}$ and a gas+star $\\Sigma_{\\rm crit}$ from \\cite{ws94}. They found that the gas+star $\\Sigma_{\\rm crit}$ gave the best threshold for determining where star formation occurs. The gas alone was sub-threshold throughout the disk. \\cite{zs05} showed that a threshold like $\\Sigma_{\\rm crit}$ may be used to determine the gas fraction in galaxies. If all galaxies have $\\Sigma({\\rm HI})$ approximately at the critical $\\Sigma_{\\rm crit} =\\alpha \\kappa \\sigma/\\pi G$, which is proportional to $V/R$ from $\\kappa$, then $M_{\\rm gas} = \\int_R 2\\pi R\\Sigma_{\\rm crit}dR \\propto VR$. This was shown to be the case from observations. They also considered that the total mass is $M_{\\rm tot}\\propto V^2R$, in which case $M_{\\rm tot}/M_{\\rm gas}\\propto V$, the rotation speed. This was also shown to be confirmed by observations. In their interpretation, small galaxies are more gas-rich than large galaxies because all galaxies have their gas column densities close to the surface density threshold. For the Milky Way, \\cite{mis06} used {\\rm CO}BE/DIRBE observations to get both the gas and dust distributions and the SFR distribution. They found a gas-law slope of $2.18\\pm0.20$, which they claimed was similar to Kennicutt's (1998) bivariate fit slope $n=2.5$ for normal galaxies. \\cite{luna06} determined the Milky Way SFR from IRAS point sources and the {\\rm CO} surface density from a southern hemisphere survey (assuming constant H$_{\\rm 2}$/{\\rm CO}). They found star formation concentrated in low-shear spiral arms and suggested an additional dependence on shear. Overall they derived $\\Sigma_{\\rm SFR}\\sim\\Sigma_{\\rm gas}^{1.2\\pm0.2}$. \\cite{voro03} also suggested a shear dependence for the SFR based on observations of the Cartwheel galaxy, where there is an inner ring of star formation with high shear that is too faint for the normal Kennicutt law, given the gas column density. \\subsection{The Q Threshold}\\label{sect:q} A threshold for gravitational instabilities in rotating disks has been derived for various ideal cases. For an infinitely thin disk of isothermal gas, the dispersion relation for radial waves is $\\omega^2 = k^2\\sigma^2 - 2\\pi G\\Sigma k + \\kappa^2$. Solving for the fastest growth rate $\\omega$ gives the wavenumber at peak growth, $k = \\pi G\\Sigma/\\sigma^2$, and the wavelength, $\\lambda=2\\sigma^2/G\\Sigma$, which is on the order of a kiloparsec in main galaxy disks. The dominant unstable mass is $M\\sim(\\lambda/2)^2\\Sigma=\\sigma^4/G^2\\Sigma\\sim10^7\\;M_\\odot$ in local spirals. The peak rate is given by \\begin{equation} \\omega_{peak}^2 = -(\\pi G\\Sigma/\\sigma^2)^2 + \\kappa^2 = -(\\pi G \\Sigma/\\sigma^2)^2(1-Q^2)\\end{equation} which requires $Q\\equiv\\kappa\\sigma/\\pi G\\Sigma < 1$ for instability (i.e., when $\\omega_{peak}^2<0$). Disk thickness weakens the gravitational force in the in-plane direction by an amount that depends on wavenumber, approximately as $1/(1+kH)$ for exponential scale height $H$ \\citep[e.g.,][]{e87, ko07}. Typically, $k\\sim1/H$, so this weakening can slow the instability by a factor of $\\sim2$, and it can make the disk slightly more stable by a factor of 2 in $Q$. On the other hand, cooling during condensation decreases the effective value of the velocity dispersion, which should really be written $\\gamma^{1/2}\\sigma$ for adiabatic index $\\gamma$ that appears in the relation $\\delta P\\propto\\delta \\rho^\\gamma$ with pressure $P$ and density $\\rho$. If $P$ is nearly constant for changes in $\\rho$, as often observed, then $\\gamma\\sim0$. \\cite{myers78} found $\\gamma\\sim0.25$ for various thermal temperatures at interstellar densities between 0.1 cm$^{-3}$ and 100 cm$^{-3}$. Thus the effects of disk thickness and a soft equation of state partially compensate for each other. There is also a $Q$ threshold for the collapse of an expanding shell of gas \\citep{epe02}. Pressures from OB associations form giant shells of gas and cause them to expand. Eventually they go unstable when the accumulated gas is cold and massive enough, provided the induced rotation and shear from Coriolis forces are small. Considering thousands of initial conditions, these authors found that a sensitive indicator of whether collapse occurs before the shell disperses is the value of $Q$ in the local galaxy disk, i.e., independent of the shell itself. The fraction $f$ of shells that collapsed scaled inversely with $Q$ as $f\\sim0.5- 0.4\\log_{10} Q$. The Toomre $Q$ parameter is also likely to play a role in the occurrence of instabilities in turbulence-compressed gas on a galactic scale \\citep{e02}. Isothermal compression has to include a mass comparable to the ambient Jeans mass, $M_{\\rm Jeans}$, in order to trigger instabilities. The turbulent outer scale in the galaxy is comparable to the Jeans length, $L_{\\rm Jeans}$, which is about the galactic gas scale height, $H$. If the compression distance exceeds the epicyclic length, then Coriolis forces spin up the compressed gas, leading to resistance from centrifugal forces. So instability needs $L_{\\rm Jeans}\\le L_{\\rm epicycle}$, which means $Q\\le1$, since $L_{\\rm Jeans}\\sim H \\sim \\sigma^2/\\pi G\\Sigma$. The epicyclic length is $L_{\\rm epicycle}\\sim\\sigma/k$, so $L_{\\rm Jeans}/L_{\\rm epicycle}=Q$. The dimensionless parameter $Q$ measures the ratio of the centrifugal force from the Coriolis spin-up of a condensing gas perturbation to the self-gravitational force, on the scale where gravity and pressure forces are equal, which is the Jeans length. The derivation of $Q$ assumes that angular momentum is conserved, so the Coriolis force spins up the gas to the maximum possible extent. When $Q>1$, a condensing perturbation on the scale of the Jeans length spins up so fast that its centrifugal force pulls it apart against self-gravity. Larger-scale perturbations have the same self-gravitational acceleration (which scales with $\\Sigma$) and stronger Coriolis acceleration (which scales with $\\kappa^2/k$); smaller-scale perturbations have stronger accelerations from pressure. If angular momentum is not conserved, then the disk can be unstable for a wider range of $Q$ because there is less spin up during condensation. For example, the Coriolis force can be resisted by magnetic tension or viscosity and then the angular momentum in a condensing cloud will get stripped away. This removes the $Q$ threshold completely \\citep{chandra54,stephenson61,lyndenbell66,hh83}. In the magnetic case, the result is the Magneto-Jeans instability, which can dominate the gas condensation in low-shear environments like spiral arms and some inner disks \\citep{e87, e91, e94, ko01, ko02, kos02}. For the viscous case, \\cite{gammie96} showed that for $Q$ close to but larger than 1, i.e., in the otherwise stable regime, viscosity can make the gas unstable with a growth rate equal to nearly one-third of the full rate for a normally unstable ($Q<1$) disk. A dimensionless parameter for viscosity $\\nu$ is $\\nu\\kappa^3/G^2\\Sigma^2$, which is $\\sim11$ according to \\cite{gammie96}. This is a large value indicating that galaxy gas disks should be destabilized by viscosity. An important dimensionless parameter for magnetic tension is $B^2/(\\pi G\\Sigma^2)\\sim8$, which is also large enough to be important. Thus gas disks should be generally unstable to form small spiral arms and clouds, even with moderately stable $Q$, although the growth rate can be low if $Q$ is large. \\subsection{Modern Versions of the KS Law with $\\sim1.5$ slope} \\cite{ken07} studied the local star formation law in M51 with 0.5-2 kpc resolution using Pa-$\\alpha$ and 24$\\mu$+H$\\alpha$ lines for the SFR, and a constant conversion factor for {\\rm CO} to H$_{\\rm 2}$. There was a correlation, mostly from the radial variation of both SFR and gas surface density, with a slope of $1.56\\pm0.04$. There was no correlation with $\\Sigma({\\rm HI})$ alone, as this atomic component had about constant column density ($\\sim10\\;M_\\odot$). The correlation with molecules alone was about the same as the total gas correlation. \\cite{leroy05} studied dwarf galaxies and found that they have a molecular KS index of $1.3\\pm0.1$, indistinguishable from that of spirals, except with a continuation to lower central ${\\rm H}_{\\rm 2}$ column densities (i.e., down to $\\sim10\\;M_\\odot$ pc$^{-2}$). \\cite{heyer04} found a slope $n=1.36$ for $\\Sigma_{\\rm SFR}$ versus $\\Sigma({\\rm H}_{\\rm 2})$ in M33, where the molecular fraction, $f_{\\rm mol}$ is small. The correlation with the total gas was much steeper. More recently, \\cite{verley10} studied M33 again and got $\\Sigma_{\\rm SFR}\\propto\\Sigma_{\\rm H2}^n$ for $n=1$ to 2, and $\\Sigma_{\\rm SFR}\\propto\\Sigma_{\\rm total\\;gas}^n$ for $n=2$ to 4. The steepening for total gas is again because $\\Sigma_{\\rm {\\rm HI}}$ is about constant, so the slope from {\\rm HI} alone is nearly infinite. This correlation is dominated by the radial variations in both quantities, as it is a point-by-point evaluation throughout the disk. Radial changes in metallicity, spiral arm activation, tidal density, and so on, are part of the total correlation. \\cite{verley10} also try other laws, such as $\\Sigma_{\\rm SFR}\\propto \\left(\\Sigma_{\\rm H2} \\rho_{\\rm ISM}^{0.5}\\right)^n$, for which $n=1.16\\pm0.04$, and $\\Sigma_{\\rm SFR}\\propto\\rho_{\\rm ISM}^n$, for which $n=1.07\\pm0.02$. These differ by considering the conversion from column density to midplane density, using a derivation of the gaseous scale height. The first of these would have a slope of unity if the star formation rate per unit molecular gas mass were proportional to the dynamical rate at the average local (total) gas density. The second has the form of the original Schmidt law, which depends only on density. To remove possible effects of {\\rm CO} to H$_{\\rm 2}$ conversion, Verley et al. also looked for a spatial correlation with the 160 $\\mu$ opacity, $\\tau_{160}$, which is a measure of the total gas column density independent of molecule formation. They found $\\Sigma_{\\rm SFR}\\propto\\tau_{160}^n$ for $n=1.13\\pm0.02$, although the correlation was not a single power law but a 2-component power law with a shallow part (slope $\\sim0.5$) at low opacity ($\\tau_{160}<10^{-4}$) and a steep part (slope $\\sim2$) at high opacity. \\subsection{Explanations for the 1.5 slope} Prior to around 2008, the popular form of the KS law had a slope of around 1.5 when $\\Sigma_{\\rm SFR}$ was plotted versus total gas column density on a log-log scale. This follows from a dynamical model of star formation in which the SFR per unit area equals the available gas mass per unit area multiplied by the rate at which this gas mass gets converted into stars, taken to be the dynamical rate, \\begin{equation} \\Sigma_{\\rm SFR}\\sim \\epsilon\\Sigma_{\\rm gas} \\left(G\\rho_{\\rm gas}\\right)^{1/2}. \\end{equation} If the gas scale height is constant, then $\\Sigma_{\\rm gas}\\propto\\rho_{\\rm gas}$ and $\\Sigma_{\\rm SFR}\\propto\\Sigma_{\\rm gas}^{1.5}$. In the model of star formation where star-forming clouds are made by large-scale gravitational instabilities, this 1.5 power law would work only where the Toomre instability condition, $Q\\leq1.4$, is satisfied. Such a model accounts for the Kennicutt (1989, 1998) law with the $Q<1.4$ threshold. Several computer simulations have shown this dynamical effect. \\cite{li06} did SPH simulations of galaxy disks with self-gravity forming sink particles at densities larger than $10^3$ cm$^{-3}$. They found a $Q$ threshold for sink particle formation, and had a nice fit to the KS law with a slope of $\\sim1.5$. \\cite{tb06} ran ENZO, a 3D adaptive mesh code, with star formation at various efficiencies, various temperature floors in the cooling function, and various threshold densities. Some models had a low efficiency with a low threshold density and other models had a high efficiency with a high threshold density. Some of their models had feedback from young stars. They also got a KS slope of $\\sim1.5$ for both global and local star formation, regardless of the details in the models. \\cite{kravtsov03} did cosmological simulations using N-body techniques in an Eulerian adaptive mesh. He assumed a constant efficiency of star formation at high gas density, and star formation only in the densest regions ($n>50$ cm$^{-3}$, the resolution limit), which are in the tail of the density probability distribution function \\citep[pdf; cf.][]{e02, km05}. \\cite{kravtsov03} got the KS law with a slope of 1.4 for total gas surface density. \\cite{wn07} did a similar thing, using the fraction of the mass at a density greater than a critical value from the pdf ($\\rho_{\\rm crit}=10^3$ cm$^{-3}$) to determine the star formation rate. Their analytical result had a slope of 1.5. \\cite{harfst06} had a code with a hierarchical tree for tracking interacting star particles, SPH for the diffuse gas, and sticky particles for the clouds. They included mass exchange by condensation and evaporation, mass exchange from stars to clouds (via PNe) and from stars to diffuse gas (SNe), and from clouds into stars during star formation. New clouds were formed in expanding shells. Their KS slope was $1.7\\pm0.1$. They also got a drop in $\\Sigma_{\\rm SFR}$ at low $\\Sigma_{\\rm gas}$, not from a $Q$ threshold but from an inability of the gas to cool and form a thin disk \\citep[cf.][]{burkert92,ep94}. ", "conclusions": "The empirical star formation law on kpc scales is essentially one where star formation follows {\\rm CO}-emitting molecular gas with a constant rate per molecule, and the ratio of molecular to atomic gas scales nearly directly with the ISM pressure \\citep{bigiel08,leroy08}. The rate per molecule corresponds to a consumption time of molecular gas equal to about 2 Gyr. The place in a galaxy where the transition occurs between {\\rm HI} dominance in the outer part to H$_{\\rm 2}$ dominance in the inner part is at a pressure of $P=2.3\\pm1.5\\times10^4k_{\\rm B}\\; {\\rm K}\\; {\\rm cm}^{-3}$. There also tend to be characteristic gas and stellar column densities at this place, and a characteristic galactic orbit time for all of the galaxies observed. Beyond this radius is the atomic-dominated outer disk. There, the SFR scales directly with $\\Sigma_{\\rm {\\rm HI}}$, and the consumption time is about 100 Gyr. Theoretical models of these empirical laws include the atomic-to-molecular transition in individual clouds and a sum over clouds to give the galactic scaling laws. Star formation occurs only in the densest parts of the clouds, as determined by a combination of turbulence-compression and self-gravity. Numerous simulations of star formation in galaxies can reproduce these empirical laws fairly well. The simulations usually show a sensitivity to the Toomre $Q$ parameter, unlike the observations." }, "1101/1101.5204_arXiv.txt": { "abstract": "We have identified a complete, flux-limited, ($S_{\\rm 160} > 120$~mJy), sample of 160$\\mu$m-selected sources from $Spitzer$ observations of the 1-deg$^2$ ISO Deep Field region in the Lockman Hole. Ground-based UV, optical and near-infrared (NIR) photometry and optical spectroscopy have been used to determine colors, redshifts and masses for the complete sample of 40 galaxies. Spitzer-IRAC+MIPS photometry, supplemented by ISOPHOT data at 90$\\mu$m and 170$\\mu$m, has been used to calculate accurate total infrared luminosities, $L_{\\rm IR}(8-1000\\micron)$, and to determine the IR luminosity function (LF) of luminous infrared galaxies (LIRGs). The maximum observed redshift is $z \\sim 0.80$ and the maximum total infrared luminosity is log~($L_{\\rm IR}/L_{\\sun}) = 12.74$. Over the luminosity range log~($L_{\\rm IR}/L_{\\sun}) = 10-12$, the LF for LIRGs in the Lockman Hole Deep Field is similar to that found previously for local sources at similar infrared luminosities. The mean host galaxy mass, log~$(M/M_{\\sun}) = 10.7$, and dominance of HII-region spectral types, is also similar to what has been found for local LIRGs, suggesting that intense starbursts likely power the bulk of the infrared luminosity for sources in this range of $L_{\\rm IR}$. However for the most luminous sources, log~$(L_{\\rm IR}/L_{\\sun}) > 12.0$, we find evidence for strong evolution in the LF $\\propto (1+z)^{6 \\pm 1}$, assuming pure number density evolution. These ultraluminous infrared galaxies (ULIRGs) have a larger mean host mass, log~$(M/M_{\\sun}) = 11.0$, and exhibit disturbed morphologies consistent with strong-interactions/mergers, and they are also more likely to be characterized by starburst-AGN composite or AGN spectral types. ", "introduction": "Deep surveys at rest-frame far-infrared (FIR) wavelengths are important for identifying statistically complete samples of luminous infrared galaxies (LIRGs: $L_{\\rm IR}/L_{\\sun} > 11.0$) -- objects that appear to produce the bulk of the bolometric infrared luminosity density in the universe, and which are often ``hidden\" and misidentified in deep UV-optical surveys. Until recently, progress has been relatively slow in identifying complete samples of FIR sources selected at wavelengths $\\lambda = 100-200\\mu$m, which corresponds to the wavelength range where the majority of LIRGs at $z < 1$ emit their peak emission. The {\\it Infrared Astronomical Satellite (IRAS)} all-sky survey provided the first complete census of 60$\\mu$m-selected galaxies in the local universe \\citep[e.g.][]{soi89,san03b}, but lacked the sensitivity and long wavelength coverage to detect LIRGs at $z>0.05$. The {\\it Infrared Space Observatory (ISO)} provided increased sensitivity and longer wavelength coverage, but background instabilities often limited the determination of exact source positions. The {\\it Spitzer Space Telescope} eventually provided the combination of long wavelength sensitivity and background stability needed to detect sources at the $\\sim 100$mJy level with relatively accurate positions, and extensive follow-up observations are now underway to identify source properties and redshifts. \\indent In this paper, we report observations first begun as part of the U.S-Japan ISO-ISOPHOT Deep Survey of a $\\sim$1-deg$^2$ region in the Lockman Hole \\citep{kaw98,oya05}, and later expanded to include more recent infrared observations of the same region obtained as part of the Spitzer Wide-area InfraRed Extragalatic (SWIRE) survey \\citep{lon03}. Our final sample consists of MIPS-160$\\mu$m sources with fluxes greater than 120mJy. Multi-wavelength follow-up observations include Keck spectra of the majority of the sources, along with UV-optical photometry from the Sloan Digital Sky Survey \\citep[SDSS:][]{yor00} for all of our sources, NIR photometry from the 2-Micron All Sky Survey \\citep[2MASS:][]{skr06} for most of our sources, and VLA 1.4GHz continuum images of the ISO-ISOPHOT deep fields \\citep{yun03}. While other surveys \\citep[e.g. COSMOS:][]{sco07} may offer superior (particularly ancillary) data, the Lockman Hole survey discussed here subtends a total of 1.2~deg$^2$ and is thus comparable in total area to similar existing datasets. This survey therefore substantially increases the total area to date at these wavelengths and helps guard against cosmic variance. \\indent Sections 2 and 3 describe our field selection and identification of SWIRE MIPS-160 sources, respectively. Methods used for identifying optical counterparts are presented in \\S~3, along with UV-NIR photometry and optical spectroscopy for individual sources. The spectral energy distributions (SEDs) and calculated infrared luminosities ($L_{\\rm IR}$) for each source are presented in \\S~5, and these data are then used to determine the Infrared galaxy Luminosity Function (LF). Section 6 presents evidence for evolution in the LF at the highest infrared luminosities, as well as a discussion of the properties of the host galaxies, including morphology, colors, masses and spectral types. Our conclusions are presented in \\S~7. \\\\ \\\\ ", "conclusions": "\\subsection{Evidence For Possible Evolution in the Luminosity Function} The space density of galaxies with infrared luminosity (8 - 1000$\\mu$m) in the range log~$(L_{\\rm IR}/L_{\\sun}) = 10 - 12$ appears to be consistent between the RBGS and our LH sample. In particular, \\citet{san03b} fit a broken power-law to the RBGS sample. At log~$(L_{\\rm IR}/L_{\\sun}) = 9.5 - 10.5$ the RBGS is fit with $\\Phi (L) \\propto L^{-0.6 \\pm 0.1}$, and at log~$(L_{\\rm IR}/L_{\\sun})= 10.5 - 12.5$ the power-law is: $\\Phi (L) \\propto L^{-2.2 \\pm 0.1}$. The Lockman Hole data for luminosities log~$(L_{\\rm IR}/L_{\\sun}) < 12$, agree within their errors to these power-laws. This concurrence is to be expected given the relatively low redshifts sampled in these lower luminosity bins. At log~$(L_{\\rm IR}/L_{\\sun}) > 12.0$ the situation changes. The co-moving space density of ultraluminous infrared galaxies (ULIRGs: log~$(L_{\\rm IR}/L_{\\sun}) > 12.0$) in the log~$(L_{\\rm IR}/L_{\\sun})= 12.0 - 12.4$ luminosity bin is $\\sim 7\\times$ higher in the LH than in the RBGS. The median redshift of the ULIRGs in the LH in this luminosity bin is $z = 0.51$. In the highest luminosity bin, log~$(L_{\\rm IR}/L_{\\sun}) = 12.4 - 12.8$, the median redshift of the 4 LH galaxies in this bin is $z = 0.71$. To compare the co-moving space density with the RGBS in this bin we extrapolate the RBGS power-law to log~$(L_{\\rm IR}/L_{\\sun}) = 12.6$ and find that the density in the LH sample is $\\sim 11\\times$ higher. \\indent Our new results for the LF of the most luminous infrared galaxies in the LH are consistent with strong evolution in the co-moving space density of ULIRGs. If we assume pure space-density evolution of the form $(1+z)^{n}$, our new results for the LH imply $n \\sim 6 \\pm 1$. This is similar to what was found in an earlier study of the infrared luminosity function of ULIRGs by \\citet{kim98}, where the co-moving space density of ULIRGs in the IRAS 1-Jy sample (mean $z \\sim 0.15$), was found to be $\\sim2\\times$ larger than the local space density of ULIRGs in the RBGS (mean $z \\sim 0.05$), implying $n = 7.6 \\pm 3.2$ . Our new results are also consistent with a recent determination of the extragalactic 250$\\mu$m luminosity function by \\citet{dye10}, which shows a ``smooth increase\" with redshift of a factor of 3.6$\\times$ in the co-moving space density of luminous infrared sources between $z = 0$ and $z = 0.2$, corresponding to $n = 7.1$. \\indent Deeper far-infrared surveys currently underway with $Spitzer$ and $Herschel$ will eventually allow us to determine whether the strong evolution observed for the most luminous infrared extragalactic sources in the relatively nearby universe continues out to higher redshift. For now, we simply note that if we assume similar strong evolution, e.g. $(1+z)^6$, in the ULIRG population out to higher redshifts, our results would imply a co-moving space-density of ULIRGs that is $\\sim$700$\\times$ larger at $z \\sim 2$ compared to the value at $z = 0$. Is there evidence for such a large population of ULIRGs at high redshift? The answer seems to be yes. There is a population of faint submillimeter sources detected by the Submillimeter Common User Bolometer Array (SCUBA) on the James Clerk Maxwell Telescope (JCMT), which has been interpreted variously as exotic objects, or ULIRGs at high redshift \\citep{sma97, hug98, bar98, lil99}. \\citeauthor{lil99} argued that these objects are indeed ULIRGs at $z \\sim 2$. Subsequently, \\citet{cha05} measured a range of spectroscopic redshifts, $z=1.7-2.8$ for a sample of 73 submillimeter galaxies, and suggested an evolution in number density of three orders of magnitude for ULIRGs between $z=0$ and $z\\sim 2.5$. Our results for ULIRGs in the LH, when extrapolated out to $z = 2-2.5$ are then consistent with the hypothesis that the SCUBA submillimeter sources are indeed ULIRGs. \\subsection{Galaxy Properties} To achieve a better understanding of the processes responsible for the observed infrared emission and the nature of the galaxies in our MIPS-160$\\mu$m sample, we use our UV-NIR imaging data and optical spectra to determine galaxy morphology and masses, and spectral types, respectively. \\subsubsection{Imaging: Morphology and Masses} In order to develop a picture of the morphologies, and to gain an indication of the prevalence of merging/interacting galaxies in the sample we examine their UV-NIR images. We compile color composite images of the sources from those available through the Finding Chart section of the SDSS DR7 website, and show these in luminosity order in Figure \\ref{grij}. The zoom on these cutouts is scaled so that each box is 100kpc on a side. At high redshifts and thus high zoom, the image quality of the SDSS charts is low, so for sources with $z > 0.3$ we display stacked $g^\\prime r^\\prime i^\\prime$ images from KPNO when available. A brief description of the galaxy morphologies is presented in Table \\ref{morph}. Many of the higher luminosity sources with log$(L_{\\rm IR}/L_{\\sun}) > 11.5$ exhibit features suggestive of interactions/mergers, such as multiple cores and/or tidal tails. At luminosities lower than log$(L_{\\rm IR}/L_{\\sun}) < 11$ the large majority of sources appear to be mostly unperturbed spirals. These trends are consistent with previous studies of local samples of LIRGs and ULIRGs \\citep[e.g.][]{san96}, which have shown that strong interactions and mergers appear responsible for triggering the most luminous infrared sources. \\indent Stellar masses for each of the MIPS-160$\\mu$m sources are listed in Table \\ref{morph}. The masses were computed by fitting the UV-NIR SEDs using {\\it Le Phare} \\citep{ilb10} and assuming a Chabrier \\citep{cha03} initial mass function (IMF). The mass range is log~$(M/M_\\sun) \\sim 10.0 - 11.5$ corresponding to $\\sim 0.5-3 M^*$. Higher mass systems are more likely to be associated with higher infrared luminosity. \\subsubsection{Spectroscopy: Extinction, Abundances and Spectral Types} Our high-resolution Keck/ESI spectra, supplemented by SDSS spectra and four low-resolution ESI spectra, allow us to measure robust spectral types for 25 of the 160\\micron\\ sources in our sample. An example of these spectra was shown in Figure \\ref{spec}. Spectra for the 25 sources with spectral types as well as those from three sources for which we have data but were unable to measure spectral types are available in the online edition of the Journal. After carefully accounting for the effects of stellar absorption and applying an extinction correction (median $E(B-V) = 0.7$), we classify the spectra as \\ion{H}{2}-region-like, or star-forming (H); composite, or star formation + an AGN (C); Seyfert (S); or LINER (L), based on the classification scheme proposed by \\citet{kew06} (see Figure \\ref{stype}). Table \\ref{lumtab} lists the measured spectral types and extinctions for each galaxy when they are available. \\indent Dividing the subsample with spectral types into luminosity bins, we find that 15 of 19, (79\\%), of galaxies in the log$(L_{\\rm IR}/L_\\sun) < 11$ bin have \\ion{H}{2}-region-like spectral type, consistent with star formation as the dominant source of excitation. This is close to the \\ion{H}{2}-region-like fraction of nearby galaxies selected at 60\\micron\\ \\citep[$\\sim$70\\% -- ][]{vei95, yua10}. The other four galaxies include one Seyfert and one LINER and 2 objects with mixed types that suggest a ``composite\" starburst-AGN mixture of excitation. Only six galaxies with high resolution spectroscopy (and hence derived spectral types) have log$(L_{\\rm IR}/L_\\sun) > 11$. Three (3/6 = 50\\%) have \\ion{H}{2}-region-like spectral type, one is a Seyfert, one is a Seyfert/LINER, and one is a ``composite\" mixture of starburst/LINER excitation. Although the fraction of galaxies with \\ion{H}{2}-region-like spectral types decreases at higher infrared luminosity (similar to what is observed for nearby galaxies selected at 60\\micron), the number statistics in this high luminosity bin are too low to draw conclusions about the fractions of different spectral types. In an attempt to provide additional information on the spectral types of our high infrared luminosity sources, we have employed a new technique developed from studies of SDSS galaxies \\citep{smo08} that maps UV/optical continuum colors onto the spectral line diagnostic diagram. This method is described in the Appendix, where the SDSS photometry for all of our MIPS-160$\\mu$m sources is used to derive ``P1,P2\" photometric spectral types for each source, following the prescription given by \\citet{smo08}. These results both confirm the large \\ion{H}{2}-region-like fraction among the lower luminosity infrared sources, and show that composite and AGN spectral types appear to increase among the highest luminosity sources. \\indent Because our spectra also contain the [\\ion{O}{2}] $\\lambda\\lambda$3727, 3729 doublet, we are able to estimate gas-phase oxygen abundances for these systems. Where available, these are listed in Table \\ref{morph}, using the robust [\\ion{N}{2}]/[\\ion{O}{2}] diagnostic of \\citet{kew02}. To put these in context, we also used the measured $K_{\\rm s}$ data to compare to the luminosity-metallicity relation in the NIR \\citep{sal05}. For the 8 systems that have sufficient information (upper-branch $R_{23}$ gas abundances and measured luminosities; see \\citet{rup08} for more on the methodology), we find that 5 follow the $L-Z$ relation of normal galaxies. Three others have higher luminosities than the data threshold, and appear to be slightly below the $L-Z$ relation (by $0.1-0.2$~dex), as found for other infrared-selected objects at high luminosity \\citep{rup08}. \\subsubsection{Radio-FIR Correlation} The measured 1.4 GHz radio continuum fluxes are converted to 1.4 GHz radio power $L_{\\rm 1.4GHz}$ assuming a spectral index of $\\alpha =+0.75$\\footnote{Spectral index $\\alpha$ is defined as $S_\\nu = S_0 (\\frac{\\nu}{\\nu_0})^{-\\alpha}$.}, and they are plotted as a function of redshift on the left panel of Figure \\ref{fig:radio}. The observed 1.4~GHz radio power range between $10^{20}$ and $10^{24.3}$ W Hz$^{-1}$ (see Table \\ref{lumtab}), similar to the IR-selected galaxies in the local universe studied by \\citet{yun01b}, and none of the sources has sufficient radio power to be classified as a ``radio-loud'' object. The most luminous infrared sources also tend to be those with the largest radio luminosities, i.e. log$(L_{\\rm 1.4GHz}) = 23.0-24.4$, equivalent to the radio powers typically seen among Seyfert galaxies, and thus the presence of a low luminosity AGN cannot be ruled out based on these radio powers alone. The well-known correlation between the measured infrared luminosity and radio power for star forming galaxies is often quantified using the ratio commonly referred to as ``$q$-value'' \\begin{equation}q={\\rm log}~[({\\rm FIR}/3.75\\times 10^{12}~{\\rm Hz})/S_{1.4 {\\rm GHz}}]\\end{equation} where FIR is the far-infrared flux density and $S_{1.4 {\\rm GHz}}$ is in W m$^{-2}$ Hz$^{-1}$ \\citep{condon92,yun01b}. We computed these $q$-values using the $L_{\\rm 1.4GHz}$ derived above and $L_{\\rm FIR}$ computed from the best-fit SED models integrated between $\\lambda=40$ and 500 \\micron, where the wavelength range has been chosen to match the original definition of $L_{\\rm FIR}$ used to compute ``$q$\". As shown on the right panel in Figure \\ref{fig:radio}, the derived $q$-values of the LH 160 \\micron\\ sources fall between 1.6 and 3.0, suggesting that most of these sources follow the same radio-FIR correlation as the local star forming galaxy populations. Some of the low redshift ($z\\lesssim 0.15$) sources appear to have $q$-values on the high end of the local population. These are also the sources with the largest angular size, and the VLA measurements are likely under-estimates as a consequence. None of the LH 160 \\micron\\ sources has a $q$-value less than 1.6 and thus a clear evidence for a radio-loud AGN." }, "1101/1101.5803_arXiv.txt": { "abstract": "Strong gravitational lensing of an extended object is described by a mapping from source to image coordinates that is nonlinear and cannot generally be inverted analytically. Determining the structure of the source intensity distribution also requires a description of the blurring effect due to a point spread function. This initial study uses an iterative gravitational lens modeling scheme based on the semilinear method to determine the linear parameters (source intensity profile) of a strongly lensed system. Our `matrix-free' approach avoids construction of the lens and blurring operators while retaining the least squares formulation of the problem. The parameters of an analytical lens model are found through nonlinear optimization by an advanced genetic algorithm (GA) and particle swarm optimizer (PSO). These global optimization routines are designed to explore the parameter space thoroughly, mapping model degeneracies in detail. We develop a novel method that determines the L-curve for each solution automatically, which represents the trade-off between the image $\\chi^2$ and regularization effects, and allows an estimate of the optimally regularized solution for each lens parameter set. In the final step of the optimization procedure, the lens model with the lowest $\\chi^2$ is used while the global optimizer solves for the source intensity distribution directly. This allows us to accurately determine the number of degrees of freedom in the problem to facilitate comparison between lens models and enforce positivity on the source profile. In practice we find that the GA conducts a more thorough search of the parameter space than the PSO. ", "introduction": "Strong gravitational lens effects produce multiple distorted images of a background object and also provide magnification of lensed sources. Magnification may reveal unresolved features in lensed sources and provides a useful tool for studying cosmologically distant objects. Furthermore, gravitational lensing provides a unique method to determine the mass distribution of lensing objects, which can be most accurately modeled when the lens potential is probed in parallel at many points. Therefore accurate lens inversion methods for extended sources are important because they provide a large number of constraints on the lens mass distribution. Models of both the intensity profile of the source and the lens mass distribution are required to model a strong gravitational lens system. Analytical models of the source are sometimes used because they are typically described by a small number of parameters, and can ensure smoothness and positivity when used to model the source intensity distribution. However, the correct parameterization is not always clear for such models, and the choice of a specific parametric model biases the lens and source solutions. Authors have attempted to partially overcome this drawback by using complex but flexible parametric models specified by large parameter sets. The most extreme example is \\citet{Tyson}, who used an elaborate source model with more than $200$ parameters to fit the gravitational lens $CL0024 +1654$. In such cases, it may be simpler to use pixelized source models, which treat all pixels on the source independently. The difficulty with pixelized source models is that they require many more free parameters than even the most complex analytical models. The optimization of extended sources via pixelized intensity distributions is simplified using the versatile semilinear method developed by \\citet{WD03} and later expanded upon by a number of authors, including \\citet{TK04}. The semilinear method uses a pixelized source, and also incorporates the blurring due to the point-spread function (PSF) of the instrument used to obtain the data. Additive noise in the observational data is also taken into account by the semilinear method. In this paper we detail Mirage, a gravitational lens modeling code written in MATLAB and C. The present version of Mirage is designed to optimize the parameters of analytical lens models and pixelized sources, but work is underway to extend the code to handle non-parametric lens models as well. A modified version of the semilinear method forms the backbone of our lens modeling program. We use sophisticated global optimization methods to fit the lens parameters, and the semilinear method to determine the corresponding source light profile that best matches the data. As a final step, we employ the method of \\citet{BLGA} to enforce the positivity of the source while keeping the nonlinear lens parameters constant. This affords a method of comparison between distinct lens density models because the number of degrees of freedom is well-defined and fixed \\citep{BL06}. The global optimizers studied in this paper consist of a sophisticated genetic algorithm (GA), called Ferret \\citep{fiege04}, and an enhanced particle swarm optimizer (PSO), Locust, which are both components of the Qubist Global Optimization Toolbox by \\citet{Fiege}. This paper discusses a robust method for gravitational lens reconstructions, highlights the benefits of both types of optimization routines, and compares their performance. In Section \\ref{sec:lensing} we will review the gravitational lens inverse problem, the semilinear method and our new matrix-free approach to lens modeling. In Section \\ref{sec:fullopt} we discuss the details of the GA and PSO, as well as a variety of simulated data tests. Section \\ref{sec:results} presents our results using these methods, and our conclusions are summarized in Section \\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} The semilinear method provides an elegant way to describe gravitational lens inversion in terms of a least squares problem, but is limited to relatively small images and a narrow PSF. This is due to the fact that the semilinear method requires the inversion of a large matrix whose size increases as the fourth power of the number of source pixels, and the sparsity of this matrix is reduced as larger PSFs are used. Solving for lens parameters is a nonlinear optimization problem, which can be solved by global optimization techniques. We applied and compared the Ferret GA and Locust PSO to determine the nonlinear parameters of the lens model. The global optimization of lens parameters requires a lens inversion for each set of lens parameters tested, and $10^4-10^5$ such evaluations are required for a thorough exploration of the parameter space and mapping of the optimal region. This reinforces the need for fast lens inversion techniques that scale well with the size of the image and PSF. We addressed the need for a fast lens inversion algorithm by developing a matrix-free approach to solve the least squares lensing problem, based partly on recent developments in the image deblurring literature, which solves the problem without the need to explicitly build the lens or blurring operators. This approach is intended to complement the semilinear method when speed is of the essence, or when large images and broad, highly structured PSFs are used. We note that our approach can be extended to the case of a spatially variant PSF. Our analysis evaluated the convergence behavior of a matrix-free method using several local optimization methods. We found that the CGLS method is fastest to converge, but all linear optimization schemes suffer from over-fitting of noise if the optimization is not stopped at the critical iteration, which cannot be predicted a priori. We showed that steepest descent methods are more robust against over-fitting to noise at the expense of the speed of convergence. The number of degrees of freedom in the iterative optimization step is estimated using a Monte Carlo method, allowing us to draw connections to the work of \\citet{suyu06} that estimate the number of degrees of freedom using Bayesian statistics. We derived a formula for the number of degrees of freedom based on the filter factors of the Tikhonov regularization problem, which agrees with the expression found by \\citet{suyu06} using Bayesian analysis. We developed a novel method that computes the optimally regularized solution for each set of lens parameters by finding the point of maximum curvature in the trade-off curve between $\\chi^2$ and a measure of the amount of regularization in the solution, which we took to be the sum of the squares of source pixel intensities. The ambiguity of choosing a regularization parameter or stopping criteria is removed, because we automatically determine the optimal number of iterations (regularization constant) using the L-curve. We evaluate the fitness of lens parameter sets using the image $\\chi^2$ statistic. The convergence and parameter space mapping properties of the Ferret GA and the Locust PSO schemes were compared, and we determined that the GA explores the parameter space more thoroughly than the PSO. The GA obtained a more detailed optimal set of solutions, highlighting the degeneracy in the position angle of a Singular Isothermal Elliptical lens model due to the rotational symmetry of the lens. Both methods converge at a similar rate. As a final refinement step in the image reconstruction our approach uses the GA or PSO to directly solve for pixel intensities. This addition has the important benefit that the non-negativity of the source intensity profile can be enforced. It is notable that the Ferret GA was able to solve this bounded linear solution refinement problem, but the Locust PSO failed due to the high dimensionality of the search ($\\sim$2500 parameters). This analysis step shows stable convergence, and noise is introduced to the source very slowly. In practice this routine is relatively insensitive to stopping criteria. This paper serves as a foundation for future explorations, which will apply the techniques discussed here to data, and expand them to include non-parametric lens models, such as those discussed by \\citet{VK09} and \\citet{Saha}. Non-parametric lens density models are extremely valuable, since dark matter haloes may contain significant substructure \\citep{K05} that is not taken into account by analytical lens models. GAs have been applied to this problem previously, specifically by \\citet{Lies07}, using the work of \\citet{Diego} as a starting point. \\citet{Lies09} used such an approach to model the system SDSS $J1004+4112$. This approach could be used in conjunction with the semilinear method to model complicated non-parametric lens density distributions and reveal the details of lensed extended sources." }, "1101/1101.5342_arXiv.txt": { "abstract": "The \\textit{Chandra X-ray observatory} has proven to be a vital tool for studying high-energy emission processes in jets associated with Active Galactic Nuclei (AGN). We have compiled a sample of 27 AGN selected from the radio flux-limited MOJAVE (Monitoring of Jets in AGN with VLBA Experiments) sample of highly relativistically beamed jets to look for correlations between X-ray and radio emission on kiloparsec scales. The sample consists of all MOJAVE quasars which have over 100 mJy of extended radio emission at 1.4 GHz and a radio structure of at least 3$\\arcsec$ in size. Previous \\textit{Chandra} observations have revealed X-ray jets in 11 of 14 members of the sample, and we have carried out new observations of the remaining 13 sources. Of the latter, 10 have X-ray jets, bringing the overall detection rate to $\\sim$ 78$\\%$. Our selection criteria, which is based on highly compact, relativistically beamed jet emission and large extended radio flux, thus provides an effective method of discovering new X-ray jets associated with AGN. The detected X-ray jet morphologies are generally well correlated with the radio emission, except for those displaying sharp bends in the radio band. The X-ray emission mechanism for these powerful FR II (Fanaroff-Riley type II) jets can be interpreted as inverse Compton scattering off of cosmic microwave background (IC/CMB) photons by the electrons in the relativistic jets. We derive viewing angles for the jets, assuming a non-bending, non-decelerating model, by using superluminal parsec scale speeds along with parameters derived from the inverse Compton X-ray model. We use these angles to calculate best fit Doppler and bulk Lorentz factors for the jets, as well as their possible ranges, which leads to extreme values for the bulk Lorentz factor in some cases. When both the non-bending and non-decelerating assumptions are relaxed the only constraints on the kpc scale jet from the Chandra and VLA observations are an upper limit on the viewing angle, and a lower limit on the bulk Lorentz factor. ", "introduction": "Blazar jets are generated in active galactic nuclei (AGN) as a result of accretion onto supermassive black holes, and can transport energy over large distances. These outflows tend to show apparent superluminal speeds, and are oriented at very shallow angles with respect to the line of sight \\citep{AS80}. The blazar class encompasses flat spectrum radio quasars (FSRQs), which have Fanaroff-Riley type II jets (FR II; \\citealt{FR74}), and BL Lac objects, which are thought to have FR I type jets \\citep{UP95}. The AGN outflows that we discuss here are of the powerful FR type II class which have well collimated jets and bright terminal hotspots (e.g., \\citealt{PK08}). In terms of X-ray production in the jet, the inverse Compton radiation process is suggested to be more important in FSRQs than in the less powerful BL Lac sources \\citep{GT08,HK06}. Prior to the launch of the \\textit{Chandra X-ray observatory}, there were very few AGN jet detections in the X-ray band. The only major X-ray imaging telescopes in use were \\textit{Einstein} and \\textit{ROSAT}. Only a few very bright, nearby X-ray jets were known, e.g., M87, 3C 273, Centaurus A, and a few lesser known sources (see \\cite{RS04}, \\cite{HM05} and references therein). Since the launch of \\textit{Chandra} there have been approximately 50 new discoveries of X-ray jets that are spatially correlated to some extent with the radio emission. The excellent angular resolution of \\textit{Chandra} has revealed detailed structure in FR II jets, such as knots, lobes and hotspots \\citep{HC02}, and has opened up an entirely new subfield of AGN astronomy. Many X-ray emitting jets were discovered in early surveys by \\cite{RS04} \\& \\cite{HM05}. The quasars in these surveys were selected mostly from radio imaging surveys of FSRQs, but the surveys were not statistically complete. For our study we have chosen to assemble a complete sample of beamed FR II jets according to well defined selection criteria. These jets generally have high Doppler factors and relativistic speeds. The MOJAVE Chandra Sample (MCS) is a complete subset of compact radio jets selected from the MOJAVE sample. The latter sample consists of all 135 known AGN with $\\delta$ $>$ $-20^\\circ$, $|$b$|$ $>$ 2.5$^\\circ$ and VLBA 15 $\\sim$ GHz correlated flux density exceeding 1.5 Jy at any epoch between 1994.0 and 2004.0 (2 Jy for AGN below $\\delta = 0\\arcdeg$) \\citep{ML09a}. Since the long interferometric baselines of the VLBA are insensitive to large-scale unbeamed radio emission, the sample is heavily dominated by blazars. In Section 2, we describe how the MCS was selected using a set of criteria designed to maximize the chances of X-ray jet detection. The goals of our study are threefold. First, we seek to identify new X-ray jets for future follow up with \\textit{Chandra}, \\textit{Spitzer}, and the \\textit{Hubble Space Telescope (HST)}. Because of the large redshift range of the MCS (0.033 $\\leq$ z $\\leq$ 2.099), we can examine the effects of proposed X-ray mechanisms such as inverse Compton scattering off of cosmic microwave background (IC/CMB) photons by relativistic electrons in the jets, which is highly dependent on redshift. Second, we wish to characterize the ratio of X-ray emission to radio emission for a large complete sample of jets. Such information is vital for determining the respective roles that deceleration and bending play in determining why jets associated with some AGN are strong X-ray emitters. Lastly, we can use the detailed viewing angle and speed information of the AGN jets on parsec (pc) scales provided by the MOJAVE program to better model the X-ray emission mechanism(s). A total of 14 AGN in the MCS have previously been observed by \\textit{Chandra} \\citep{HM05,RS04,CO02,SJ06,JM06,WE07,WA01}; here we present new 10 kilosecond exposures on the other 13 sources, in which we have detected 10 new X-ray jets. A detailed analysis of the full 27 source sample will be presented in subsequent papers. Our paper is laid out as follows; we describe the MOJAVE Chandra Sample in Section 2, along with our data reduction method and selection criteria. In Section 3, we describe the jet observations for each specific source in which a jet was present in both the radio and X-ray images. In Section 4, we discuss overall source trends and provide additional ancillary information on selected sources. In Section 5, we discuss implications of the model with respect to the bulk Lorentz factor and viewing angle. We summarize our conclusions in Section 6. The limits for the derived Doppler factor and bulk Lorentz factors are given in the Appendix. Throughout this paper we use a standard cosmology with H$_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega$$_m$ = 0.27, and $\\Omega$$_{\\Lambda}$ = 0.73. ", "conclusions": "We have performed Chandra observations of a radio-core-selected sample of blazar jets. The selection criteria that we used to define our AGN sample has increased the overall fraction of correlations between X-ray and radio jets in radio selected AGN. Of the popular single zone models available (synchrotron or IC/CMB), we chose to apply the IC/CMB model to our sample, based on the earlier results of \\cite{HM05}. The detected X-ray jets are generally well correlated spatially with the radio jet morphology, except for those radio jets that display sharp bends. The wide range of apparent X-ray to radio ratios among the jets suggests that no single overall emission model can explain all of the X-ray morphologies. We are currently analyzing follow-up \\textit{Chandra} and \\textit{HST} observations of selected AGN to obtain multiwavelength spectra of jet knots (Kharb et al. 2011, in prep.), which will allow us to investigate possible synchrotron and IC models for the emission beyond what we have discussed in this paper. Our major findings are as follows: \\begin{itemize} \\item The selection criteria associated with the MCS has increased the detection rate from previous jet surveys \\citep{RS04,HM05} from a $\\sim$ 60\\% detection rate to a $\\sim$ 78\\% detection rate. \\item We have found that the 1.4 GHz, VLA-A array extended radio jet flux density, S$_{ext}$, is a strong predictor of X-ray jet emission in a core-selected sample such as the MCS, which is related to the correlation of the extended luminosity and the pc scale jet (apparent) speed. Above a value of 300 mJy we find a 100\\% X-ray detection rate, with $\\sim$ 57\\% detection rate for sources located below that threshold. This further reinforces the usefulness of our extended radio emission selection criteria for this sample. \\item The IC/CMB assumptions can produce calculated values for the jet bulk Lorentz factor, $\\Gamma$, which are larger than expected in some sources (eg. 0106+013 and 1849+670) under the assumption that the jet speed and direction are the same on both the pc and kpc scales. \\item Bending alone can not reconcile the large $\\Gamma$ values in these sources as it constrains the minimum $\\Gamma$ value on the kpc scale to the minimum value on the pc scale. This can still be quite large as seen in sources such as 0106+013 and 1849+670. \\item If we allow for the possibility of deceleration with out jet bending, the VLBI jet speeds and IC/CMB X-ray model can be reconciled, although jet bending is necessary in several cases. In this scenario the kpc scale relativistic jet bulk Lorentz factors typically range from $\\sim$ 1.7 to 7. \\item When both the non-bending and non-decelerating assumptions are relaxed the only constraints on the kpc scale jet from the Chandra and VLA observations are an upper limit on the viewing angle, and a lower limit on the bulk Lorentz factor. These typically range from $8^{\\circ} < \\theta < 20^{\\circ}$ and 1.6 $< \\Gamma_{min} <$ 3.5 for our sample. \\end{itemize} This research has made use of data from the MOJAVE database that is maintained by the MOJAVE team \\citep{ML09a}. This work was supported by Chandra Award GO8-9113A. We would also like to thank, the referee, Dan Harris for the insightful comments he provided that helped to improve this paper. \\clearpage" }, "1101/1101.0751_arXiv.txt": { "abstract": "High resolution spectropolarimetric observations by {\\em Hinode} have revealed the existence of supersonic downflows at the umbra-penumbra boundary of 3 sunspots \\citep{2011ApJ...727...49L}. These downflows are observed to be co-spatial with bright penumbral filaments and occupy an area greater than 1.6 arcsec$^2$. They are located at the center-side penumbra and have the same polarity as the sunspot which suggests that they are not associated with the Evershed flow. In this paper we describe the supersonic velocities observed in NOAA AR 10923 and discuss the photospheric as well as chromospheric brightenings that lie close to the downflowing areas. Our observations suggest that this phenomenon is driven by dynamic and energetic physical processes in the inner penumbra which affect the chromosphere, providing new constraints to numerical models of sunspots. ", "introduction": "\\label{intro} The Evershed flow \\citep[EF;][]{1909MNRAS..69..454E} is a distinct property of sunspot penumbrae which exemplifies their filamentary structure \\citep[][and references therein]{2003A&ARv..11..153S}. In the inner penumbra the EF starts as upflows \\citep{2006A&A...453.1117B, 2006ApJ...646..593R, 2009A&A...508.1453F} that turns into downflows in the mid and outer penumbra \\citep{1997Natur.389...47W, 1999A&A...349L..37S, 2003A&A...410..695M, 2004A&A...427..319B}. In addition to the EF, other types of mass motions exist in the penumbra, as reported recently by \\citet{2010A&A...524A..20K} using {\\em Hinode} observations. They detected small-scale downflowing patches with velocities of $\\sim$1~km~s$^{-1}$ which have the same polarity as the parent sunspot. Some of them also appear to be co-spatial with chromospheric brightenings. Based on their physical properties, \\citet{2010A&A...524A..20K} inferred that these weak downflows are different from the Evershed flow returning to the photosphere, which sometimes happens well within the penumbra \\citep{2004A&A...427..319B,2007ApJ...668L..91B,2008A&A...481L..21S}. \\begin{figure}[t] \\centering \\includegraphics[width=11cm,angle=90,bb=-7 100 595 792]{./louis_fig1.eps} \\caption{Continuum image of NOAA AR 10923 at 630~nm. The white dashed square represents the small region chosen for analysis and whose magnified image is shown in the inset. The white arrow points to disk center.} \\label{cont_image} \\end{figure} \\citet{2011ApJ...727...49L} observed a new type of downflows which are supersonic and occur at or near the umbra-penumbra boundary of sunspots. These downflowing patches are conspicuously large in size with areas ranging from 1.6--6 arcsec$^2$. They have the same polarity as the sunspot and occur along bright penumbral filaments. Their properties are different from those of the Evershed flow and the downflows reported by \\citet{2010A&A...524A..20K}, which indicates a different physical origin. The strong downflows possibly represent energetic and dynamic processes occurring in the inner penumbra which also affect the chromosphere. In this paper we restrict our discussion to the supersonic downflows observed in NOAA AR 10923 and briefly describe the photospheric and chromospheric activities associated with them. ", "conclusions": "\\label{orient} The downflows that are associated with the EF can sometimes be supersonic in the outer penumbra \\citep{2001ApJ...549L.139D,2004A&A...427..319B} or even beyond the sunspot boundary \\citep{2009ApJ...701L..79M}. Such a configuration represents mass flux returning to the photosphere and has a polarity opposite to that of the sunspot. The supersonic downflows we have observed have the same polarity as the parent sunspot and so they cannot be related to the Evershed downflows. One could assume that these strong downflows are the photospheric counterpart of some kind of inverse Evershed flow seen in the chromosphere. However, it is not clear how such a chromospheric phenomenon could produce supersonic downflows close to the umbra-penumbra boundary in the photosphere. The orientation of the filaments P1 and P2, bifurcating at the strong downflowing patch (Fig.~\\ref{unsharp}), resembles the post-reconnection configuration illustrated in Fig.~5c of \\citet{2008ApJ...686.1404R}, suggesting that the origin of the downflows is the slingshot effect associated with the reconnection of the filaments. The bisecting angles shown by the solid green lines were estimated to be $51\\deg$ and $46\\deg$. According to \\citet{2008ApJ...686.1404R}, the unwinding of filaments in a cork screw fashion can lead to reconnection, transient brightenings and twists in the penumbral filaments. This model was proposed as a possible mechanism for producing penumbral MJs. \\citet{2010ApJ...715L..40M} investigated the above scenario using numerical simulations and concluded that MJs occur in the intermediate region between nearly horizontal flux tubes and the relatively vertical background field of the penumbra. In this model, only parts and not the entire penumbral filament participate in the reconnection process. Slingshot reconnection may be a possible mechanism for producing the supersonic downflows and the photospheric as well as chromospheric brightenings, although there is no strict one-to-one correspondence between the two phenomena. The above process has to be different from the one producing MJs since their intensities and lifetimes are much smaller than the events described in this work. \\begin{figure}[t] \\centering \\vspace{-2pt} \\hspace{20pt} \\includegraphics[width=0.55\\textwidth,angle=90]{./louis_fig6.eps} \\vspace{-25pt} \\caption{Continuum image that has been unsharp masked using a $3 \\times 3$ pixel boxcar. The red arrow indicates the location where the filaments P1 and P2 appear to intersect each other. The solid green lines refer to the bisecting angles between the filaments P1 and P2. The blue contour has been drawn for LOS velocities greater than 2 km~s$^{-1}$.} \\label{unsharp} \\end{figure}" }, "1101/1101.5497_arXiv.txt": { "abstract": "The diffuse intracluster light (ICL) contains a significant fraction of the total stellar mass in clusters of galaxies, and contributes in roughly equal proportion as the hot intra-cluster medium (ICM) to the total baryon content of clusters. Because of the potential importance of understanding the origin of the ICL in the context of the formation and evolution of structure in the Universe, the field has recently undergone a revival both in the quality and quantity of observational and theoretical investigations. Due to cosmological dimming, the observational work has mostly concentrated on low redshift clusters, but clearly observations at higher redshifts can provide interesting clues about the evolution of the diffuse component. In this paper we present the first results of a program to characterize the ICL of intermediate redshift clusters. We find that at $z\\sim0.3$, the X-ray cluster RX J0054.0-2823 already has a significant ICL and that the fraction of the total light in the ICL and the brightest cluster galaxy (BCG) is comparable to that of similar clusters at lower redshift. We also find that the kinematics of the ICL is consistent with it being the remnant of tidally destroyed galaxies streaming in the central regions of the cluster, which has three central giant elliptical galaxies acting as an efficient ``grinding machine''. Our cluster has a bi-modal radial-velocity distribution and thus two possible values for the velocity dispersion. We find that the cluster fits well in the correlation between BCG+ICL fraction and cluster mass for a range of velocity dispersions, leading us to question the validity of a relevant correlation between these two quantities.% ", "introduction": "\\label{SECintro} \\footnotetext[1]{Based on observations collected at the European Organisation for Astronomical Research in the Southern Hemisphere, Chile. ESO 078A--0456; ESO 65.O--0425} The diffuse intracluster light (ICL) in rich clusters of galaxies is by now a mature field that was pioneered by \\cite{Zwicky1951,Zwicky1952} and quantified using innovative photoelectric and photographic techniques by \\cite{deVaucouleurs1970}, \\cite{Melnick1977} and \\cite{Thuan1977}. But it was not until the advent of low-noise panoramic CCD detectors that the field matured and evolved to the present state of the art \\citep{Zibetti2005,Gonzalez2007,Mihos2005,Krick2006}. However, even with state of the art CCDs on the most powerful telescopes, the task of detecting and quantifying the diffuse intergalactic component is far from straightforward. Not only are we dealing with signals that are much fainter than the sky, but it is also very difficult to disentangle the emission from bona-fide free floating intergalactic stars from the outer haloes of the brightest cluster galaxies (BCGs). Recently, there have been substantial improvements in the theoretical understanding of the origin of the ICL through detailed numerical simulations \\citep[eg.][and references therein]{Rudick2009,Dolag2009a,Puchwein2010,Murante2007}, and the field is evolving from being purely observationally driven, to a phase of rich interaction between theory and observations. Numerical simulations indicate that most of the ICL is accrued fairly recently in the merging history of clusters, probably at redshifts $z<1$. Some of the ICL is produced in earlier sub-mergers and falls into the clusters together with their parent groups and smaller clusters, while much of the intracluster light comes from stars stripped off massive galaxies that fall into the forming cluster. Recent state of the art hydrodynamical simulations \\citep{Puchwein2010} predict that the ICL should contain a significant fraction ($\\sim45\\%$) of the total stellar mass in clusters, well in excess of the values typically observed. Thus, observations of the diffuse intracluster light in clusters should yield information not only about the origin of the ICL, but also about the formation of the clusters themselves. Mapping the evolution of the ICL as a function of redshift should, at least in principle, provide a novel cosmological probe. In the course of an investigation of the diffuse intergalactic light in X-ray emitting clusters at intermediate redshifts, we detected a puzzling S-shaped arc-like structure in the ROSAT cluster RX J0054.0-2823 that we tentatively identified as the gravitationally lensed image of a background galaxy at a redshift between $z=0.5$ and $z=1.0$ \\citep{Faure2007}. The cluster is characterized by having three dominant D class galaxies, two of which are interacting forming a dumbbell cD system. Thus, although our lensing models reproduced surprisingly well the arc-structure, the possibility that it could be tidal debris related to these massive galaxies could not be excluded. In fact, numerical simulations of cluster formation show that galaxies and groups falling in the cluster potential actually retain their identities for rather long times, so the ICL should actually have a filamentary structure, which is actually observed in some clusters \\citep[see][ and references therein]{Mihos2005,Rudick2009}. We therefore embarked on an ambitious project to measure the redshift of the arc, which has a low surface brightness so, unless the spectrum is dominated by strong emission lines, required long integration times. We designed an observing strategy allowing us at the same time to observe the arc, the diffuse intracluster Light (ICL), and a substantial number of individual galaxies in the field taking advantage of the multi-object spectroscopic mode of the FORS2 instrument on Paranal. This allowed us to obtain the redshifts of more than 650 galaxies in addition to very deep observations of the arc, and very deep images and long-slit observations of the ICL. The analysis of the field population around the cluster are presented in two separate papers \\citep{Giraud2010,Giraud2011} that also contain the details of the observations and the data analysis procedures. The present paper is devoted to the investigation of the diffuse intra-cluster light component and the analysis of the dynamics of the cluster based on our deep images and deep spectroscopy. In Section~\\ref{SECobs} we present a detailed description of the observations and the data reduction techniques for the photometry and for the long-slit spectroscopy that has not been described in our previous papers. Section~\\ref{results} presents the dynamical and photometric properties of the cluster and the ICL and details the techniques that we used to detect and measure the ICL component and to define its photometric and spectroscopic properties, including those of the S-shaped arc. Finally, Sections~\\ref{SECanal} and \\ref{SECconcl} discuss the results and presents the conclusions of this work. Throughout this work we use the standard cosmology $H_0=75~\\mathrm{km~s^{-1}}$, $\\Omega_\\mathrm{m}=0.3$; $\\Omega_\\mathrm{\\Lambda}=0.7$. ", "conclusions": "\\label{SECconcl} We have presented deep photometry and spectroscopy of the diffuse intracluster light (ICL) of an intermediate redshift cluster at z=0.3. We show that despite the redshift we are able to measure both the photometric and kinematic properties of the ICL out to distances comparable to the distances reached by state-of-the-art observations of local clusters such as Virgo. This is because, being more compact, we have less problems to subtract the sky background and the contamination by large scale features like Galactic cirrus. And of course we are able to observe with 8m-class telescopes. We found that globally the BCG+ICL component is reasonably well characterized by a de Vaucouleurs $R^{1/4}$ law, and we detect a significant outward bluing of the diffuse component. At a distance between 40~kpc and 350~kpc from the cluster centre, the colour gradient flattens and the slope of the surface brightness profile steepens probably signaling the region of pure ICL light. We also find that the D(4000) index, which is an indicator of stellar populations, is shallower in these outer regions and consistent with the value observed for faint galaxies in the cluster. This, combined by the rather chaotic kinematics of the ICL at large radii, is consistent with the observations in lower-z clusters of filamentary structures in the ICL. Our cluster is characterized by a triplet of bright elliptical galaxies in its centre, two of which form a close dumbbell pair. We argue that this structure will eventually merge thus significantly increasing the BCG+ICL component. In spite of its dynamical youth, however, the fraction of BCG+ICL to total light in our cluster is consistent with the values observed for similar mass clusters at low redshifts. If real, this indicates that most of the ICL must be formed early in the merging history of clusters. However, the radial velocity histogram of our cluster is consistent with two different values for the velocity dispersion, and curiously the cluster fits the correlation between BCG+ICL fraction and velocity dispersion for both velocity dispersions equally well. This happens because there is a well known relationship between BCG+ICL fraction and radius, and we integrate the light out to a radius that scales with velocity dispersion. This leads us to suggest that the correlation between BCG+ICL and velocity dispersion may be much weaker than previously thought, which is consistent with the predictions of hydrodynamical numerical simulations. However, the simulations predict BCG+ICL fractions on average a factor of $\\sim2$ larger than observed (with BCG+ICL fractions up to $70\\%$ and with the ICL alone having $45\\%$ of the stellar mass), but unfortunately the theoretical papers do not provide sufficient details about the modeling to allow a better understanding of the discrepancy. In particular, it is not clear how these simulations model stellar evolution, nor whether the models actually reproduce the observed colours and D(4000) steps of the ICL. We are confident in our measurement of the BCG+ICL fraction in RX~J0054.0-2823 and therefore that the discrepancy does not lie with the observations. The main conclusions of our paper can be summarized as follows: \\begin{enumerate} \\item The S-shaped arc in the ROSAT cluster RX~J0054.0-2823 is not the gravitationally lensed image of a background galaxy, but the remnant of two faint cluster galaxies in the act of being tidally destroyed by the gravitational field of the three giant elliptical galaxies in the cluster centre; \\item In spite of the relatively large redshift of the cluster ($z\\sim0.3$) we were able to clearly detect its diffuse intracluster-light (ICL) component and to measure its broad-band colours and spectral signatures. The bluer colour and shallower spectral indexes of the ICL, and its chaotic motions are consistent with the idea that the ICL originates from tidally disrupted galaxies. \\item For the ICL we measured D(4000)=1.39$\\pm$0.02, typical of post-starburst (E+A) galaxies. This value is significantly different from D(4000)=1.83$\\pm$0.01 measured for the central galaxies, and is consistent with the bluer colour of the ICL; \\item The bi-modal radial velocity distribution of the cluster galaxies makes its velocity dispersion rather uncertain, despite the fact that we have velocities for close to 100 cluster galaxies. However, we find that the cluster fits the correlation between the fraction of the light in the BCG+ICL component and velocity dispersion found by GZZ for any value of the velocity dispersion in the range allowed by the observations. This led us to question the correlation given the rather large observational uncertainties in the velocity dispersions of the low-mass clusters in the GZZ sample. However, our cluster shows the same discrepancy between the observed BCG+ICL fraction and the values predicted by numerical simulations. \\item The BCG+ICL fraction of our cluster is undistinguishable from that of similar mass clusters at lower redshift in the GZZ sample. This indicates that most of the diffuse component in clusters is already in place at z=0.3. \\end{enumerate}" }, "1101/1101.0084_arXiv.txt": { "abstract": "{Asteroseismology has entered a new era with the advent of the NASA \\emph{Kepler} mission. Long and continuous photometric observations of unprecedented quality are now available which have stimulated the development of a number of suites of innovative analysis tools.} {The power spectra of solar-like oscillations are an inexhaustible source of information on stellar structure and evolution. Robust methods are hence needed in order to infer both individual oscillation mode parameters and parameters describing non-resonant features, thus making a seismic interpretation possible.} {We present a comprehensive guide to the implementation of a Bayesian peak-bagging tool that employs a Markov chain Monte Carlo (MCMC). Besides making it possible to incorporate relevant prior information through Bayes' theorem, this tool also allows one to obtain the marginal probability density function for each of the fitted parameters. We apply this tool to a couple of recent asteroseismic data sets, namely, to \\emph{CoRoT} observations of \\object{HD 49933} and to ground-based observations made during a campaign devoted to \\object{Procyon}.} {The developed method performs remarkably well at constraining not only in the traditional case of extracting oscillation frequencies, but also when pushing the limit where traditional methods have difficulties. Moreover it provides an rigorous way of comparing competing models, such as the ridge identifications, against the asteroseismic data.} {} ", "introduction": "\\label{sect:Intro} Seismology of solar-like stars is a powerful tool that can be used to increase our understanding of stellar structure and evolution. Solar-like oscillations in main-sequence stars and subgiants have been measured thanks to data collected from ground-based high-precision spectroscopy \\citep[for a review e.g.,][]{BeddingKjeldsenReviewNew} and, more recently, to photometric space-based missions such as \\emph{CoRoT} \\citep[e.g.,][]{Michel08}. Red giants also exhibit solar-like oscillations, although at lower frequencies, and hence require longer time series in order to resolve them \\citep[e.g.,][and references therein]{Ridder09}. The launch of the NASA \\emph{Kepler} mission \\citep[][]{Koch10} definitely marked a milestone in the field of asteroseismology. \\emph{Kepler} will particularly lead to a revolution in the seismology of solar-like oscillators, since it will increase by more than two orders of magnitude the number of stars for which high-quality observations will be available, while allowing for long-term follow-ups of a selection of those targets. The large homogeneous sample of data made available by \\emph{Kepler} opens the possibility of conducting a seismic survey of the solar-like part of the colour-magnitude diagram, which researchers in the field already started naming as \\emph{ensemble asteroseismology}. As of the time of writing of this article, first results arising from the \\emph{Kepler} asteroseismic programme had already been made available \\citep[][]{Bedding10,Chaplin10,Gilliland10,Hekker10,Stello10,JCDKjeldsenEtAl2010,Gemma}. The rich informational content of power spectra of solar-like oscillations allows fundamental stellar properties (e.g.~mass, radius, and age) to be determined, and the internal structure to be constrained to unprecedented levels provided that individual oscillation mode parameters are measured \\citep[e.g.,][]{Dalsgaard}. Furthermore, the measured stellar background signal provides us with valuable information on activity and convection. In the case of the highest signal-to-noise ratio ($S/N$) observations, for which it is possible to measure individual oscillation mode parameters, we expect asteroseismology to produce a major breakthrough on stellar structure and evolution, on topics as diverse as energy generation and transport, rotation and stellar cycles \\citep[e.g.,][]{Karoff09}. For the past few years significant work has been invested in making preparations for the mode parameter analysis of \\emph{Kepler} data. This analysis involves the estimation of individual and average oscillation mode parameters, as well as estimation of parameters that describe non-resonant signatures of convection and activity. Examples include the work conducted in the framework of the AsteroFLAG consortium \\citep[][]{asteroflag} and the work undertaken by the \\emph{CoRoT} Data Analysis Team \\citep[][]{DAT}. This consequently paved the way for the development of suites of analysis tools for application to \\emph{Kepler} data \\citep[][]{Birmingham,Huber,KaCa,Mathur,Mosser,ACPS}. In the present study we give continuity to this work by presenting a comprehensive guide to the implementation of a Bayesian peak-bagging\\footnote{The term ``peak-bagging'' has become the customary name for the examination of individual oscillation peaks in the field of asteroseismology. The origin of the name is explained in \\cite{Appourchaux2003B}.} tool that employs a MCMC. These techniques derive from the tools traditionally used in helioseismology and are in many ways an extension of the Maximum Likelihood Estimation (MLE) methods. This peak-bagging tool is to be applied to the power spectra of solar-like oscillators and used as a means to infer both individual oscillation mode parameters and parameters describing non-resonant features. Besides making it possible to incorporate relevant prior information through Bayes' theorem, this tool also allows one to obtain the marginal probability density function (PDF) for each of the model parameters (frequencies, mode heights, mode lifetimes, rotational splitting, inclination angle etc.). This is one of the main advantages of these MCMC techniques, as it not only performs well in low signal-to-noise conditions, but also provides reliable error bars on the parameters. Parameter space is sampled using a \\emph{Metropolis--Hastings} algorithm featuring a built-in \\emph{statistical control system} that allows to automatically set an appropriate instrumental law during the burn-in stage. Also included is \\emph{parallel tempering}, which increases the mixing properties of the Markov chain. The outline of the paper is as follows: We start in Sect.~\\ref{sect:pspec} by providing an overview of the theory behind the power spectrum of solar-like oscillations, introducing the assumptions and the set of parameters needed to model the spectrum to the level of detail required by modern asteroseismic data. In Sect.~\\ref{sect:Bayes} we describe the subjacent Bayesian statistical framework by highlighting the topics of parameter estimation and model selection. Section \\ref{sect:MCMC} is devoted to the modus operandi of advanced Markov chain Monte Carlo methods and their implementation. In Sect.~\\ref{sect:Application} we present a couple of examples where this tool has been applied to recent asteroseismic data sets, evidencing some of its capabilities and illustrating its functioning. A summary and discussion are presented in Sect.~\\ref{sect:Conclusions}. ", "conclusions": "\\label{sect:Conclusions} In this paper, we have presented the basic theory and methods behind the extraction of parameters from the power spectra of solar-like stars. In order to handle the ever rising quality and complexity of modern asteroseismic data, we have developed a tool (APT~MCMC) that enables us to constrain parameters associated with the subtlest features in the spectra. The algorithm has been extensively tested and performs extremely well, not only in the traditional case of extracting oscillation frequencies, but also when pushing the limit where traditional methods have difficulties, such as constraining linewidths, rotational splittings and stellar inclination angles. In this work we have focused on data in the signal-to-noise regime of current asteroseismic measurements. In the case of very high signal-to-noise ratios, other features in the power spectrum becomes important, such as mode asymmetries and rotational splittings dependent on $\\ell$, arising from differential rotation with radius. In future work these effects will be incorporated into the program and tested on solar data. One disadvantage of the method is that it can be quite computationally intensive, both to implement and run, when compared to traditional MLE fits. This is however balanced by the much added information outputted from the fits, specifically in the probability distributions of each parameter, making it easy to obtain accurate, reliable and realistic error bars on the results -- a feature seriously missing from the traditional methods. The parameter estimation also benefits enormously from the possibilities the Bayesian formalism provides with inclusion of prior information. This not only allows control of the fit to, for example, not allow unphysical parameter combinations, but also include information into the fit that is better constrained by other measurements (as we saw in Sect.~\\ref{sec:HD49933}). Another powerful feature of the method lies in the parallel tempering, which not only keeps the fits from getting stuck in local maxima, but also provides an objective way of comparing different competing models, as it provides a way of calculating the global likelihood. This can for example be utilized in the familiar problem of ridge identification in solar-like stars (see Sect.~\\ref{sec:Procyon}). A thing to keep in mind is also that the APT~MCMC algorithm is completely general, in the sense that it could be applied to other problems without modification. MCMC methods are being used in various branches of astrophysics: cosmology \\citep{Liddle}, extra solar planets \\citep{Gregory2005} and stellar model fitting \\citep{Bazot}, but in fact the methods would be applicable in any problem including parameter estimation. And as computational power continues to grow, the downsides are quickly becoming insignificant. What could to some extent also be seen as a disadvantage of these methods is that they can never be fully automated, in the sense that they will not be able to handle a large number of stars without human interaction. The whole fundamental idea behind the Bayesian formalism is that it relies on ''wise`` human inputs on the priors and model setup that should not be done in an automated way. If nothing else, take this as a positive reassurance: You will, as an astrophysicist, never be obsolete to computers or monkeys with keyboards." }, "1101/1101.5945_arXiv.txt": { "abstract": "{} {We investigate the behaviour of nonlinear, nonideal Alfv\\'en wave propagation within an inhomogeneous magnetic environment.} {The governing MHD equations are solved in 1D and 2D using both analytical techniques and numerical simulations.} {We find clear evidence for the ponderomotive effect and visco-resistive heating. The ponderomotive effect generates a longitudinal component to the transverse Alfv\\'en wave, with a frequency twice that of the driving frequency. Analytical work shows the addition of resistive heating. This leads to a substantial increase in the local temperature and thus gas pressure of the plasma, resulting in material being pushed along the magnetic field. In 2D, our system exhibits phase mixing and we observe an evolution in the location of the maximum heating, i.e. we find a drifting of the heating layer.} {Considering Alfv\\'en wave propagation in 2D with an inhomogeneous density gradient, we find that the equilibrium density profile is significantly modified by {\\emph{both}} the flow of density due to visco-resistive heating {\\emph{and}} the nonlinear response to the localised heating through phase mixing.} ", "introduction": "\\label{section1} Phase mixing, a mechanism for dissipating shear Alfv\\'en waves, was first proposed by Heyvaerts \\& Priest (\\cite{Heyvaerts1983}). The basic concept is straightforward: consider shear Alfv\\'en waves propagating in an inhomogeneous plasma, such that on each magnetic fieldline each wave propagates with its own local Alfv\\'en speed. Thus, after propagating a certain distance, these neighbouring perturbations will be out of phase, which will lead to the generation of smaller and smaller transverse spatial scales, and thus to the growth of strong currents. This ultimately results in a strong enhancement of the dissipation of Alfv\\'en-wave energy via viscosity and/or resistivity. Alfv\\'en wave phase mixing has also been studied extensively as a possible mechanism for heating the corona (e.g. Heyvaerts \\& Priest \\cite{Heyvaerts1983}, Browning \\cite{Browning1991}, Ireland \\cite{Ireland1996}, Malara {{et al.}} \\cite{Malara1996}, Narain \\& Ulmschneider \\cite{Narain1990}; \\cite{Narain1996}). Nakariakov {{et al.}} (\\cite{Nakariakov1997}) extended the model of Heyvaerts \\& Priest (\\cite{Heyvaerts1983}) to include compressibility and nonlinear effects. They found that fast magnetoacoustic waves are generated continuously by Alfv\\'en-wave phase mixing at a frequency twice that of the driven Alfv\\'en wave, and that these generated waves propagate across magnetic fieldlines and away from the phase mixing layer. Since there is a permanent leakage of energy away from the phase mixing layer, these fast waves can cause indirect heating of the plasma as they propagate away and dissipate far from the layer itself, thus spreading the heating across the domain. Nakariakov {{et al.}} (\\cite{Nakariakov1997}) found that the amplitude of these fast waves grows linearly in time, according to weakly-nonlinear, $\\beta=0$, analytical theory. Nakariakov {{et al.}} (\\cite{Nakariakov1998}) further extended this model to include a background steady flow and found that the findings of their \\cite{Nakariakov1997} model persist. Hood {{et al.}} (\\cite{Hood2002}) investigated the phase mixing of single Alfv\\'en-wave pulses and found that this results in a slower power-law damping (as opposed to the standard $\\sim \\exp{(-t^3)}$ for harmonic Alfv\\'en waves, i.e. Heyvaerts \\& Priest \\cite{Heyvaerts1983}). Hence, Alfv\\'enic-pulse perturbations will be able to transport energy to a greater coronal height than that of a harmonic Alfv\\'en wavetrain. Botha {{et al.}} (\\cite{Botha2000}) considered a developed stage of Alfv\\'en-wave phase mixing and found that the growth of the generated fast waves saturates at amplitudes much lower than that of the driven Alfv\\'en wave. They concluded that the nonlinear generation of fast waves (Nakariakov {{et al.}} \\cite{Nakariakov1997}) {{saturates}} due to destructive wave interference, and has little effect on the standard phase mixing model of Heyvaerts \\& Priest (\\cite{Heyvaerts1983}). The numerical simulations of Botha {{et al.}} assumed wave propagation in an ideal plasma, and thus heating was absent from their model. Tsiklauri and co-authors repeated these numerical experiments for both a {{weakly}} and strongly-nonlinear Alfv\\'enic pulse (Tsiklauri {{et al.}} \\cite{Tsiklauri2001}; \\cite{Tsiklauri2002}). De Moortel {{et al.}} (\\cite{DeMoortel1999}; \\cite{DeMoortel2000}) investigated how gravitational density stratification and magnetic field divergence changes the efficiency of phase mixing. They report that the resultant dissipation can be either enhanced or diminished depending on the specific choice of equilibrium. Ruderman {{et al.}} (\\cite{Ruderman1998}) investigated phase mixing in open magnetic equilibria, and found similar results using the WKB method. These \\cite{Ruderman1998} analytical results were repeated and corrected by numerical and semi-analytical work by Smith {{et al.}} (\\cite{Smith2007}). All these authors found the following: a diverging magnetic field enhances the efficiency of phase mixing, whereas gravitational stratification diminishes the mechanism. Hood {{et al.}} (\\cite{Hood1997a}; \\cite{Hood1997b}) derive analytical, self-similar solutions of Alfv\\'en-wave phase mixing in both open and closed magnetic topologies. Ofman \\& Davila (\\cite{Ofman1995}) found that in an inhomogeneous coronal hole with an enhanced dissipation parameter $(S=1000-10,000)$, Alfv\\'en waves can dissipate within several solar radii, which can provide significant energy for the heating and acceleration of the solar wind. This model was later extended to include nonlinear effects (Ofman \\& Davila \\cite{Ofman1997}). Parker (\\cite{Parker1991}) pointed out that phase mixing requires an ignorable coordinate, an assumption which is expected to be unphysical in the corona. Parker found that including all three coordinates results in the driven Alfv\\'en waves coupling with a fast magnetoacoustic mode, and that this elimates the growth in transverse spatial scales, and thus phase mixing is absent from the system. Instead, such a system exhibits resonant absorption (e.g. Lee \\& Roberts \\cite{Lee1986}; Hollweg \\& Yang \\cite{Hollweg1988}) on the surfaces where the phase velocity equals the Alfv\\'en velocity. However, Parker's conclusions are in disagreement with the work of Tsiklauri and co-authors who considered the interaction of an impulsively-generated, weakly-nonlinear MHD pulse with a one-dimensional density inhomogeniety, considered in the three-dimensional regime (i.e. without an ignorable coordinate) in both an ideal (Tsiklauri \\& Nakariakov \\cite{TN2002}) and resistive (Tsiklauri {{et al.}} \\cite{Tsiklauri2003}) plasma. Tsiklauri and co-authors found that phase mixing remains a relevant paradigm and that the dynamics can still be qualitatively understood in terms of the classic 2.5D models. Mocanu {{et al.}} (\\cite{Mocanu2008}) have revisited the Heyvaerts \\& Priest (\\cite{Heyvaerts1983}) model using anisotopic viscosity (i.e. incorporating the Braginskii \\cite{Braginskii1965} stress tensor) and report that this significantly increases the damping lengths, i.e. compared to those obtained for isotropic dissipation. More recently, Threlfall {{et al.}} (\\cite{Threlfall2010}) have investigated the effect of the Hall term on phase mixing in the ion-cyclotron range of frequencies. Another key concept that we shall invoke in this paper is the ponderomotive force: a nonlinear force proportional to spatial gradients in magnetic pressure, also referred to as the Alfv\\'en wave-pressure force. The ponderomotive effect has been considered in a solar context initially by Hollweg (\\cite{Hollweg1971}) and later by Verwichte and co-authors (Verwichte \\cite{VerwichteTHESIS1999}; Verwichte {{et al.}} \\cite{VNL1999}). Hollweg (\\cite{Hollweg1971}) {{considered}} linearly-polarised Alfv\\'en waves propagating in a direction parallel to the magnetic field, and found that the transverse behaviour of the Alfv\\'en wave was identical when comparing the linear and nonlinear (to second order) solutions, but that longitudinal wave velocity and density fluctuations appear in the nonlinear solutions driven by gradients in the wave magnetic-field pressure, i.e. the Alfv\\'en wave is no longer purely transverse and is compressive (through nonlinear coupling to magnetoacoustic waves). Thus, the ponderomotive force can be used as an extra acceleration mechanism and as an explanation for density fluctuations in the solar wind. Verwichte (\\cite{VerwichteTHESIS1999}) presented a mechanical analogy for the ponderomotive effect by considering the resulting motion of discrete particles (beads) on an oscillating string. Verwichte {{et al.}} (\\cite{VNL1999}) considered the temporal evolution of a weakly-nonlinear, Alfv\\'en wave in a $\\beta=0$ homogeneous plasma. These authors showed that the an initially-excited, gaussian-pulse perturbation in transverse velocity splits into two Alfv\\'en wave pulses, each propagating in opposite directions (as naturally expected). Furthermore, Verwichte {{et al.}} found that the ponderomotive force produces a shock in longitudinal velocity at the starting position. Note that in a cold plasma, there is no force to counteract this ponderomotive acceleration. In this paper, we investigate the nonlinear, nonideal behaviour of Alfv\\'en-wave propagation and phase mixing over long timescales. Thus, this paper can be seen as an extension of the model of Botha {{et al.}} (\\cite{Botha2000}) to include visco-resistive effects. We also seek to address a fundamental question of phase mixing: by considering nonlinear, nonideal phase mixing over long timescales, {\\emph{is it possible to observe a drifting of the heating layer?}} In other words, phase mixing will occur due to the density inhomogenity in our system, and this process will generate strong, localised heating due to enhanced dissipation. This localised heat deposition is expected to modify the equilibrium density profile, and thus may change the location of maximum heating. However, it is unclear what the actual result will be: we may observe a change in the location of the heating layer, or the phase mixing mechanism may break down, or the heating may bifurcate spatially (since our density profile will no longer be monotonic). Is is also unclear how this may affect the indirect heating of the plasma, due to the coupling to the fast magnetoacoustic mode (Nakariakov {{et al.}} \\cite{Nakariakov1997}). Such an investigation requires nonlinear and nonideal effects to be considered and, as we shall see, observed over long timescales. The work of Ofman {{et al.}} (\\cite{Ofman1998}) is also relevant here. These authors investigated a model of resonant absorption that incorporated the dependence of loop density on the heating rate, and studied the spatial and temporal dependence of the heating layer. Ofman {{et al.}} find that the heating occurs in multiple resonance layers, rather than the single layer of the classic resonant absorption models (e.g. Ionson \\cite{Ionson1978}; Ulmschneider {{et al.}} \\cite{Ulmschneider1991}; Ruderman \\& Roberts \\cite{RR2002}) and that these layers drift throughout the loop to heat the entire volume. Poedts \\& Boynton (\\cite{Poedts1996}) also investigated resonant absorption using nonlinear, resistive MHD simulations and found a spreading of the heat deposition, i.e. a broadening of the resonant layer due to changes in the background inhomogeneity. This paper has the following outline: $\\S\\ref{section2}$ describes the governing equations, assumptions and analytical and numerical details of our investigation, $\\S\\ref{section:1-D}$ investigates the 1D nonlinear, nonideal system (with no density inhomogeneity) and focuses on the underlying physical processes. $\\S\\ref{section:bulk_flow}$ details the bulk-flow phenomenon found to be present in our system, and investigates its density dependence. $\\S\\ref{section:2-D}$ considers a 2D model with an inhomogeneous density profile, and details the long-term evolution and coupled nature of the three MHD waves present in our system. The conclusions are presented in $\\S\\ref{section:conclusions}$ and there are two appendicies. ", "conclusions": "\\label{section:conclusions} We have investigated the nonlinear, nonideal behaviour of Alfv\\'en wave propagation and phase mixing within an inhomogeneous environment, over long timescales. The governing MHD equations have been solved in 1D and 2D environments using both analytical techniques and numerical simulations. In an ideal, one-dimensional study (no density inhomogeneity, $\\partial/\\partial x=0$) we find that by driving a linear Alfv\\'en wave (in ${\\rm{v}}_z$) into our numerical domain, we also generate two types of longitudinal wave: boundary-driven acoustic waves (propagating at speed $c_s$) and a nonlinear perturbation (propagating at ${\\rm{v}}_A$) which is driven by the ponderomotive force. Both these motions have a frequency twice that of the driven Alfv\\'en wave, and an exact mathematical solution for both wave types was derived. The acoustic wave is degenerate under the $\\beta=0$ approximation (see Appendix $\\ref{appendix:beta=0}$) but the ponderomotive wave is always present in a nonlinear system. We find that the addition of resistive and viscous effects naturally leads to visco-resistive damping (through the dispersion relation) and thus to visco-resistive heating, which in turn leads to the introduction of a new phenomenon: a bulk flow in the positive $y-$direction. Physically, this bulk flow is a direct response to the visco-resistive heating in the system: the heating increases the local temperature which increases the local thermal pressure. The resulting presure gradient drives the bulk flow. The bulk flow is purely in the positive $y-$direction because the boundary conditions are held fixed (${{\\rm{v}}_y}=0$). As a result of this out-flow, a significant reduction in density occurs since the boundary conditions do not allow for the plasma to be replaced by an in-flow. Note that the bulk flow is not due to the ponderomotive force, otherwise such a flow would be apparent in the ideal, nonlinear system. We find that the magnitude of the visco-resistive heating is strongly dependent upon $k_I$ and thus on our choice of $\\rho_0$; the greater the value of $\\rho_0$, the more visco-resistive heating that occurs. More visco-resistive heating leads to a stronger bulk flow in the longitudinal direction. We also investigated the effect of driving an Alfv\\'en wave in a nonideal plasma over very long timescales. We find that the equilibrium density profile changes substantially over the duration of our simulation (end of the simulation at $t=100,000\\tau_A$) specifically due to this bulk-flow phenomenon (which is itself directly due to the visco-resistive heating). For a $\\rho_0=1$ plasma, we find a decrease in density of about $4.6\\%$ at $y=0$ but for a $\\rho_0=5$ plasma we find a decrease of about $46.4\\%$. We also find that the rate-of-change of density at $y=0$ is decreasing as the simulation proceeds. This is because as the density decreases, the value of $|k_I|$ also decreases. This smaller value of $|k_I|$ reduces the amount of visco-resistive damping, which slows the rate-of-change of density, and so on. Thus, there is a strong feedback effect in our system, and we can explain the nature of the bulk flow and change of density profile over a range of density values. We conclude that the bulk-flow phenomenon is {\\emph{a natural consequence of driving an Alfv\\'en wave in a nonideal plasma}}. We then extended our analysis to include an inhomogeneous density profile in a two-dimensional system and, as before, we drive our system with sinusodial, linear Alfv\\'en waves. As expected from Heyvaerts \\& Priest (\\cite{Heyvaerts1983}), our simulation displays classic phase mixing. A cut showing ${{\\rm{v}}_z}$ as a function of height along $x=35$ provided an excellent match with the analytical predictions of Heyvaerts \\& Priest (\\cite{Heyvaerts1983}). Cuts along $x=0$ ($\\:\\rho_0=5$) and $x=100$ ($\\:\\rho_0=1$) were identical to those of the corresponding visco-resistive damped Alfv\\'en waves seen in our one-dimensional analysis. Our one-dimensional analysis also explained the longitudinal motions (${{\\rm{v}}_y}$) in our system and we again observed the boundary-driven acoustic waves (interpreted as slow magnetoacoustic waves in 2D, with speed $c_s$), our ponderomotive waves (propagating at ${{\\rm{v}}_A}$) and the bulk-flow phenomenon. In addition to Alfv\\'en waves and slow magnetoacoustic waves, our 2D system also contains fast magnetoacoustic waves (seen primarily in ${{\\rm{v}}_x}$). These fast waves are continuously generated by Alfv\\'en-wave phase mixing at a frequency twice that of the driven Alfv\\'en wave, and propagate across magnetic fieldlines and away from the phase mixing layer (as predicted by Nakariakov {{et al.}} \\cite{Nakariakov1997}; \\cite{Nakariakov1998}). These waves are also refracted towards regions of lower Alfv\\'en speed. Since there is a permanent leakage of energy away from the phase mixing layer, it is possible that these fast waves can cause indirect heating of the plasma as they propagate away and dissipate far from the phase mixing layer itself, thus spreading the heating across the domain (Nakariakov {{et al.}} \\cite{Nakariakov1997}). However, due to our choice of amplitude ($A=0.01$) we find that only a small fraction of the Alfv\\'en-wave energy is converted into fast waves and, thus, only a small amount of indirect heating occurs. Over long timescales, we find that both the slow and fast waves saturate, and always remain at amplitudes of the order $A^2$. Hence, these waves only play a secondary role in the heating of the plasma. We find that at the end of our simulation, the density profile has changed substantially. This is due to two effects: firstly, the bulk-flow phenomenon is naturally present in our system, with the density changes as predicted by our one-dimensional analysis. Secondly, there is a substantial decrease in density at the locations of phase mixing heating. The change in the background density profile does not affect the propagation of the Alfv\\'en wave (amplitude of the order $A$) to a great degree, and the behaviour of ${{\\rm{v}}_x}$ (fast waves) and ${{\\rm{v}}_y}$ (slow waves) do not have a strong feedback effect on the system (both amplitudes of the order $A^2$) and only play secondary roles in the system (i.e. $A=0.01$, so feedback effect is only of the order of $1\\%$). However, these conclusions may be different for Alfv\\'en waves driven with a much larger amplitudes. We have also investigated how the temperature profile changes during our simulation. We find a substantial increase in localised temperature, generated by the phase mixing mechanism, over the duration of our simulation, and find that the maximum temperature increases by a factor of about a hundred during the simulation. We also find that the location of the maximum temperature drifts during our simulation, thus providing an answer to one of the key questions asked in $\\S\\ref{section1}$, i.e. we find that drifting of the heating layer {\\it{is}} possible in a nonlinear phase mixing scenario. However, we also find that this drifting is very small and occurs over a very long timescale (our simulation runs over $100,000\\tau_A$), at least for the parameters we have considered in this paper. For typical coronal parameters, $\\tau_A=60$ seconds and so we are considering timescales of approximately $69.4$ days, which is too long to be physically important in the corona (i.e. other coronal processes act over much shorter timescales). However, it may be possible to increase the speed and magnitude of this drifting by significantly increasing the magnitude of heating in our system, i.e. attempt to increase the amount of enhanced dissipation from phase mixing. Thus, future studies could consider increasing the amplitude and frequency of the driven Alfv\\'en wave, or increasing the steepness of the gradient in our background density inhomogeneity." }, "1101/1101.0059_arXiv.txt": { "abstract": "We report first multicolor polarimetric measurements ($UBV$ bands) for the hot Jupiters HD189733b and confirm our previously reported detection of polarization in the $B$ band (Berdyugina et al.\\ 2008). The wavelength dependence of polarization indicates the dominance of Rayleigh scattering with a peak in the blue $B$ and $U$ bands of $\\sim$10$^{-4}\\pm10^{-5}$ and at least a factor of two lower signal in the $V$ band. The Rayleigh-like wavelength dependence, detected also in the transmitted light during transits, implies a rapid decrease of the polarization signal toward longer wavelengths. Therefore, the { nondetection by Wiktorowicz (2009), based on a measurement} integrated within a broad passband covering the $V$ band and partly $B$ and $R$ bands, { is inconclusive} and consistent with our detection in $B$. We discuss possible sources of the polarization and demonstrate that effects of incomplete cancellation of stellar limb polarization due to starspots or tidal perturbations are negligible as compared to scattering polarization in the planetary atmosphere. We compare the observations with a Rayleigh-Lambert model and determine effective radii and geometrical albedos for different wavelengths. { We find a close similarity of the wavelength dependent geometrical albedo with that of the Neptune atmosphere, which is known to be strongly influenced by Rayleigh and Raman scattering.} Our result establishes polarimetry as a reliable means for directly studying exoplanetary atmospheres. ", "introduction": "Direct detection of exoplanets with polarimetry opens the prospect for probing their atmospheres. A reflecting planet breaks the symmetry of the stellar radiation and conspicuously marks its presence in polarized light \\citep[e.g.,][]{fb10,berd11}. Thanks to the differential nature of polarimetry, its dynamic range exceeds that of any other technique, but it is still a challenge to detect low signals. It is known that polarization properties of the light are generally wavelength dependent \\citep[e.g.,][]{stenflo05}, with Thomson scattering being one exception. Rayleigh scattering being an important opacity source in upper layers of cool atmospheres results in polarization strongly increasing toward the blue with a $\\lambda^{-4}$ law. Combined with an angular dependence of polarization, this dictates an optimal range of wavelengths and scattering angles at which polarization can be successfully detected. There have been a few attempts to detect polarized reflected light from an exoplanet. \\citet{hou06} and \\citet[][hereafter L09]{luc09} observed hot Jupiters in $\\tau$\\,Boo and 55\\,Cnc in a red filter of 590\\,nm to 1000\\,nm (maximum at 800\\,nm). Due to relatively low statistics and incomplete orbital phase coverage, only standard deviations of Stokes $q$ and $u$ could be deduced: 5.1$\\cdot$10$^{-6}$ and 2.2$\\cdot$10$^{-6}$ for the two systems, respectively. \\citet[][hereafter B08]{berd08} observed another hot Jupiter, HD189733b, which is twice as close to the parent star as, e.g., $\\tau$\\,Boo\\,b. The Stokes parameters were measured in the Johnson $B$ band (370--550\\,nm, maximum at 430\\,nm). Despite a relatively low accuracy of individual measurements, a very high statistics of the data (about 100 nightly measurements) and an even distribution over orbital phases have enabled the detection of variable, phase-locked polarization of maximum $\\sim$2$\\cdot$10$^{-4}$ and best-fit amplitudes in Stokes $q$ and $u$ of (1.5$\\pm$0.3)$\\cdot$10$^{-4}$ and (1.1$\\pm$0.4)$\\cdot$10$^{-4}$, respectively, for binned data. \\citet[][hereafter W09]{wik09} also observed HD189733 but in the 400--675\\,nm band (centered at 550\\,nm). An upper limit of 7.9$\\cdot$10$^{-5}$ in the polarization degree (5$\\cdot$10$^{-5}$ in Stokes $q$) from only six nightly measurements was reported. At first glance, these measurements contradict to each other. However, as they were taken at different wavelengths, { the conclusion made from} their direct comparison was incorrect. Furthermore, only one of the six W09 measurements (which provided the upper limit) is at a phase near elongation where a significant polarization signal is expected based on our B08 ephemeris. Therefore, the { nondetection} of W09 is { inconclusive} and consistent with our detection. In this paper, we show that the detected amplitudes and upper limits across the optical spectrum indicate the dominance of Rayleigh scattering in the atmospheres of hot Jupiters. This is also favored by transit spectroscopy in the optical by \\citet[][hereafter, P08]{pont08} and \\citet{lec08a,lec08b} and in the near infrared by \\citet{sing09}. Here we report first multicolor polarimetric measurements in the $UBV$ bands for HD189733b. We confirm within the standard deviation the previously reported first detection of polarization in the $B$ band by B08 and determine reflecting properties of the planet. We also find the consistency between the data of W09 and our measurements in the $V$ band. ", "conclusions": "\\label{sec:con} The new $UBV$ polarimetric observations of HD189733b confirm the first detection of polarized reflected light from the hot Jupiter by B08 at the 10$^{-4}$ level with the standard deviation of 10$^{-5}$. Our data are consistent with upper limits obtained by others at different wavelengths. They clearly demonstrate the dominance of Rayleigh scattering in the planetary atmosphere at optical wavelengths longer than 400\\,nm. At shorter wavelengths an additional opacity mechanism { (e.g., Raman scattering)} plays a significant role. Our result establishes polarimetry as an important tool for studying directly exoplanetary atmospheres in the visible and near UV. In the near future it will be employed for non-transiting systems. We appreciate an extended review by an anonymous referee. SVB acknowledges the EURYI (European Young Investigator) Award provided by the ESF (see www.esf.org/euryi) and the SNF grant PE002-104552. This work is also supported by the Academy of Finland, grant 115417. Based on observations made with the Nordic Optical Telescope, La Palma, Spain." }, "1101/1101.5160_arXiv.txt": { "abstract": "We study the visibility of the Ly$\\alpha$ emission line during the Epoch of Reionization (EoR). Combining galactic outflow models with large-scale semi-numeric simulations of reionization, we quantify the probability distribution function (PDF) of the fraction of \\lya\\ photons transmitted through the intergalactic medium (IGM), $\\mathcal{T}_{\\rm IGM}$. Our study focusses on galaxies populating dark matter halos with masses of $M_{\\rm halo}=10^{10}M_{\\odot}$ at $z=8.6$, which is inspired by the recent reported discovery of a galaxy at $z=8.6$ with strong Ly$\\alpha$ line emission. For reasonable assumptions, we find that the combination of winds and reionization morphology results in $\\mathcal{T}_{\\rm IGM}\\gsim 10 \\%$ [50\\%], for the majority of galaxies, even when the Universe is $\\sim 80 \\%$ [60\\%] neutral by volume. Thus, the observed strong Ly$\\alpha$ emission from the reported $z=8.6$ galaxy is consistent with a highly neutral IGM, and cannot be used to place statistically significant constraints on the volume averaged neutral fraction of hydrogen in the IGM. We also investigate the implications of the recent tentative evidence for a observed decrease in the `Lyman Alpha Emitter fraction' among drop-out galaxies between $z=6$ and $z=7$. If confirmed, we show that a rapid evolution in $\\avenf$ will be required to explain this observation via the effects of reionization. ", "introduction": "\\label{sec:intro} The Wide Field Camera 3 on board of the {\\it Hubble Space Telescope} has enhanced our ability to observe galaxies at redshifts great than six, so-far obtaining $\\sim 100$ likely candidate galaxies at $z=$7, 8 \\citep[e.g.][]{Bouwens10,Bunker10,Finkelstein10,Yan10}. The {\\it James Webb Space Telescope} is expected to probe galaxies a few magnitudes deeper, and also to spectroscopically confirm the redshifts of the existing candidates. One of the key predicted properties of young, metal poor galaxies in the high-redshift Universe are prominent nebular emission lines, dominated by hydrogen Ly$\\alpha$ ($\\lambda=1216$\\hs\\AA, see e.g. Johnson et al. 2009, Pawlik et al. 2010). The first generation of galaxies are likely to have been strong Ly$\\alpha$ emitters, with equivalent widths possibly as high as EW$\\sim 1500$ \\AA \\hs\\citep[][also see Partridge \\& Peebles 1967]{S02,S03,J09b}. During the epoch of reionization (EoR), the Ly$\\alpha$ emission line may be difficult to observe, due to the large opacity of the intervening neutral intergalactic medium: for example, a source needs to be embedded in a $\\gsim$1 Mpc HII region to allow Ly$\\alpha$ photons to redshift far away from the line center before they reach the IGM \\citep[e.g.,][]{M98,Cen00}. For the galaxy to generate such a large HII region, its ionizing luminosity would have to be unphysically large (unless there is a bright quasar associated with the galaxy, see Cen \\& Haiman 2000). However, several effects have been shown to boost the detectability of the Ly$\\alpha$ flux: ({\\it i}) source clustering, which boosts the sizes of HII regions \\citep{Fur04,Mes2,McQuinn07,Iliev08} ({\\it ii}) the patchiness of reionization, which may give rise to significant fluctuations in the IGM opacity between different sightlines, as well as a steeper absorption profile \\citep[e.g.,][]{Mes1,McQuinn08}, and ({\\it iii}) radiative transfer effects through outflows of interstellar (ISM) \\ion{H}{I} gas, which can impart a redshift\\footnote{The peculiar velocity of a galaxy can also redshift Ly$\\alpha$ photons away from resonance before they escape into the surrounding intergalactic medium (Cen et al. 2005). However, this redshift is typically significantly smaller than the redshift imparted by galactic outflows (see Dijkstra \\& Wyithe 2010).} to the Ly$\\alpha$ photons before they emerge from galaxies \\citep[][but also see Barnes et al. 2011]{Santos04,DW10}. \\cite{Lehnert} recently reported a detection of strong Ly$\\alpha$ line emission from a $Y_{105}$ drop-out galaxy in Wide Field Camera 3 observations of the Hubble Ultra Deep Field. The Ly$\\alpha$ line implies that the galaxy is at $z\\sim 8.56$, which is the highest redshift of any spectroscopically-confirmed object to date. Interestingly, when taken at face value, the observed Ly$\\alpha$ line is strong, with an observed equivalent width (EW) of $\\sim 100$ \\AA\\hs (see \\S~\\ref{sec:discuss}). Motivated by this observation, we study the visibility of the Ly$\\alpha$ emission line during the EoR, and compute the total fraction of emitted Ly$\\alpha$ photons that the IGM transmits directly to the observer, $\\mathcal{T}_{\\rm IGM}$. We simultaneously include the inhomogeneous large-scale reionization morphology, peculiar velocity offsets of the galaxies and IGM, and radiative transfer through the galactic outflows that is calibrated by observations of Lyman Alpha Emitters (LAEs) at $z<6$ \\citep[][]{V06,V08,V10}. This, in combination with the fact that we compute the full $\\mathcal{T}_{\\rm IGM}$-PDF, clearly distinguishes our analysis from previous work. Our models do not include dust, which at the redshifts of interest ($z=7-9$) is likely a good approximation \\citep[][]{Stanway05,Bouwens10b,Hayes10,Blanc10}. The dust opacity to Ly$\\alpha$ photons inside galaxies is expected to increase towards lower redshift, as the cumulative dust content of the Universe increased with cosmic time. This expected evolution has a different sign than the IGM opacity which decreases towards lower redshift. Thus, our discussion regarding the redshift evolution of the IGM opacity is likely conservative. The outline of this paper is as follows: in \\S~\\ref{sec:model}, we describe our models for galactic outflows (\\S \\ref{sec:winds}) and IGM opacity (\\S \\ref{sec:IGM}). In \\S~\\ref{sec:result} we present the corresponding \\lya\\ transmission fractions. Within this context, we interpret the observations of Ly$\\alpha$ emitting galaxies at $z>6$ in \\S~\\ref{sec:discuss}. We compare our results with previous work in \\S~\\ref{sec:prevwork}. Finally, we present our conclusions in \\S~\\ref{sec:conc}. We adopt the background cosmological parameters ($\\Omega_\\Lambda$, $\\Omega_{\\rm M}$, $\\Omega_b$, $n$, $\\sigma_8$, $H_0$) = (0.72, 0.28, 0.046, 0.96, 0.82, 70 km s$^{-1}$ Mpc$^{-1}$), matching the five--year results of the {\\it WMAP} satellite \\citep{Komatsu08}. Unless stated otherwise, we quote all quantities in comoving units. ", "conclusions": "\\label{sec:conc} In this paper we have studied the visibility of the Ly$\\alpha$ emission line during reionization. We combine large scale semi-numerical simulations of cosmic reionization with empirically-calibrated models of galactic outflows. With these sophisticated tools, we compute the PDFs of the IGM transmission fraction, $\\mathcal{T}_{\\rm IGM}$. We find that winds cause $\\mathcal{T}_{\\rm IGM}\\gsim 10 \\%$ [50\\%], for the majority of galaxies, even when the Universe is $\\sim 80 \\%$ [60\\%] neutral by volume. This only requires wind speeds greater than $\\sim 25$ km s$^{-1}$, which are quite conservative judging by the observed \\lya\\ lines shapes at $z<5$ (Verhamme et al. 2008, also see \\S~\\ref{sec:prevwork}). Therefore, we conclude that the observed strong Ly$\\alpha$ emission from the reported $z=8.6$ galaxy is consistent with a highly neutral IGM. We also show that evoking reionization to explain the observed drop in the `LAE fraction' (see \\S~\\ref{sec:z7}) of drop-out galaxies between $z=6$ and $z=7$ \\citep{Stark10,Stark11}, requires a very rapid evolution of $\\mathcal{T}_{\\rm IGM}$, corresponding to $\\avenf \\sim 0 \\rightarrow 0.5$ over $\\Delta z=1$. Reionization models find such a rapid evolution unrealistic, which may indicate that either ({\\it i}) the current sample of drop-out galaxies at $z=7$ happened to populate a region of our Universe that was more neutral than average, ({\\it ii}) winds become weaker and/or have smaller covering factors towards higher redshifts, or ({\\it iii}) that the Universe at $z=6$ still contained a non-negligible volume fraction of neutral hydrogen. However, these conclusions are tentative as the available data still has large uncertainties. Regardless of these current observational uncertainties, our work underlines the point that Ly$\\alpha$ emission can be detected from galaxies in the earliest stages of reionization. This is a positive result for (narrowband) searches for high redshift Ly$\\alpha$ emitters such as the `Emission-Line galaxies with VISTA Survey' (ELVIS) \\citep[e.g.][]{N07}. On the other hand, if a neutral IGM is quite transparent to Ly$\\alpha$ photons, then a signature of reionization may be more difficult to extract from observations of Ly$\\alpha$ emitting galaxies. However, the {\\it redshift evolution} of quantities such as ({\\it i}) the `LAE fraction'--or more generally the Ly$\\alpha$ restframe equivalent width PDF-- among LBGs \\citep{Stark10,Fontana10,Stark11}, and ({\\it ii}) the UV and Ly$\\alpha$ luminosity functions of LAEs (Kashikawa et al. 2006), already provide interesting and useful constraints on models of reionization. Furthermore, the clustering signature of LAEs (\\citealt{McQuinn07}; Mesinger \\& Furlanetto 2008b, though see Iliev et al. 2008) is also affected by reionization, and it has already been shown that winds do not affect this prediction \\citep{McQuinn07}. \\vskip+0.5in {\\bf Acknowledgements} Support for this work was provided by NASA through Hubble Fellowship grant HST-HF-51245.01-A to AM, 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." }, "1101/1101.5356_arXiv.txt": { "abstract": "I review recent work on X-ray sources in Galactic globular clusters, identified with low-mass X-ray binaries (LMXBs), cataclysmic variables (CVs), millisecond pulsars (MSPs) and coronally active binaries by Chandra. Faint transient LMXBs have been identified in several clusters, challenging our understanding of accretion disk instabilities. Spectral fitting of X-rays from quiescent LMXBs offers the potential to constrain the interior structure of neutron stars. The numbers of quiescent LMXBs scale with the dynamical interaction rates of their host clusters, indicating their dynamical formation. Large numbers of CVs have been discovered, including a very faint population in NGC 6397 that may be at or beyond the CV period minimum. Most CVs in dense clusters seem to be formed in dynamical interactions, but there is evidence that some are primordial binaries. Radio millisecond pulsars show thermal X-rays from their polar caps, and often nonthermal X-rays, either from magnetospheric emission, or from a shock between the pulsar wind and material still flowing from the companion. Chromospherically active binaries comprise the largest number of X-ray sources in globular clusters, and their numbers generally scale with cluster mass, but their numbers seem to be reduced in all globular clusters compared to other old stellar populations. ", "introduction": "Observations by the first X-ray satellite Uhuru demonstrated a strong overabundance of X-ray sources in globular clusters compared to the rest of the galaxy, which was quickly attributed to dynamical formation mechanisms in the dense cores of globular clusters \\citep{Clark75}. The bright ($L_X>10^{36}$ \\ergss) X-ray sources in globular clusters have mostly been shown to be neutron stars (NSs) in low-mass X-ray binaries (LMXBs) through their X-ray bursts \\citep{Lewin95,intZand03}. Several mechanisms to create these systems have been suggested, including exchange of NSs into primordial binaries \\citep{Hills76}, tidal capture of another star by the NS \\citep{Fabian75}, and collisions of NSs with giants \\citep{Sutantyo75}. These mechanisms all have a similar dependence on the structural parameters of the host globular clusters, as they depend on the rate of close encounters, and most encounters take place in the nearly-constant-density core of the clusters. Comparisons of the close encounter rates between clusters often use the parametrization $\\Gamma=\\rho_c^{2}r_c^{3}/\\sigma=\\rho_c^{1.5}r_c^2$ \\citep{Verbunt87}, where $\\rho_c$ is the central density, $r_c$ is the cluster core radius, and $\\sigma$ is the central velocity dispersion (a lower $\\sigma$ promotes gravitational focusing in close encounters). Indeed, this quantity reasonably describes the probability of finding LMXBs in Galactic globular clusters \\citep{Verbunt87}, and (although limited by our optical spatial resolution) is also related to the probability of finding LMXBs in globular clusters in other galaxies \\citep{Jordan04,Sivakoff07, Jordan07,Peacock09}. Fifteen bright LMXBs are now known in 12 globular clusters, of which at least seven are transients showing bright ``outbursts'' and longer periods in which they are quite faint \\citep{Verbunt04,Heinke10}. Periods are now known for ten of these systems, with recent discoveries of the periods of M15 X-2 \\citep[22.6 min.,][]{Dieball05}, SAX J1748.9-2021 in NGC 6440 \\citep[8.5 hrs,][]{Altamirano08}, 4U 0513-40 in NGC 1851 \\citep[17 min.,][]{Zurek09}, NGC 6440 X-2 \\citep[57.3 min.,][]{Altamirano10}, and IGR J17480-2446 in Terzan 5\\citep[21.252 hours,][]{Papitto10,Strohmayer10}. \\citet{Deutsch00} pointed out the excess of ``ultracompact'' periods (now at least 5 of 15) below one hour among globular cluster sources (compared to the period distribution among LMXBs in the rest of the Galaxy), indicating degenerate companions. Such systems are most easily created through collisions of NSs with red giants \\citep[e.g.][]{Ivanova05}, indicating the importance of this mechanism. Two transient systems (both, coincidentally, in NGC 6440; Fig. 1) have been found to show coherent millisecond X-ray pulsations during outbursts, identifying the rotational period of the NSs \\citep{Altamirano08,Altamirano10}, while the new transient in Terzan 5, IGR J17480-2446 \\citep{Bordas10,Pooley10b}, is a slow (11 Hz) pulsar \\citep{Strohmayer10b}. NGC 6440 X-2 (and, so far, IGR J17480-2446) show pulsations throughout all outbursts, while SAX J1748.9-2021 shows them only occasionally, perhaps due to its higher average mass-transfer rate burying its magnetic field \\citep{Cumming01}. NGC 6440 X-2 shows unusual outburst behavior, with its outbursts lasting $<$4 days (above $10^{35}$ erg/s), reaching relatively low peak $L_X$s of $<2\\times10^{36}$ erg/s, and (most unusually) recurring every $\\sim$31 days during much of 2009 and 2010 \\citep{Heinke10}. As it is hard to identify such faint outbursts even in the sensitive RXTE PCA Galactic Bulge scans \\citep{Swank01}, it seems quite likely that other such weak transients may be missed by current instruments. \\begin{figure} \\includegraphics[height=.36\\textheight]{6440_images} \\includegraphics[height=.36\\textheight]{Ter5_color4_chopnew.eps} \\caption{{\\bf Left:} \\Chandra\\ X-ray images of the cluster NGC 6440 during an outburst of NGC 6440 X-2 (top left), one week later (top right), during an outburst of SAX J1748.9-2021 (bottom right), and from two quiescent periods (bottom left). Images on the right have $\\sim$1/10 the exposure time as those on the left. The cluster core radius, and the position of 6440 X-2, are indicated on all panels \\citep{Heinke10}. {\\bf Right:} Representative-color X-ray image of the extremely rich globular cluster Terzan 5 (1-2 keV red, 2-3 keV green, 3-6 keV blue). Quiescent LMXBs generally appear reddish, while the white sources are likely dominated by cataclysmic variables.} \\end{figure} Even fainter transients, the very faint X-ray transients ($10^{34}$ M$_K$ $>$ --18 mag (for which dwarf galaxies have B--K $\\sim$ 2). %However, a single gap exists %at a galaxy mass of $\\sim$100 million solar masses %for which no dynamical %studies are available. On the high mass side of the gap are %This gap covers the transition from galaxies classified as dwarf elliptical (dE) galaxies, that are dominated by stars in their inner regions. While the low mass side includes dwarf spheroidal (dSph) galaxies that are dark matter-dominated and ultra compact dwarf (UCD) objects that are star-dominated. Evolutionary pathways across the gap have been suggested but remain largely untested because the `gap' galaxies are faint, making dynamical measurements very challenging. With long exposures on the Keck %world's largest optical telescope using the ESI instrument we have succeeded in bridging this gap by measuring the dynamical mass for five dwarf galaxies with M$_K$ $\\sim$ --17.5 (M$_B$ $\\sim$ --15.5). With the exception of our brightest dwarf galaxy, they possess relatively flat velocity dispersion profiles of around 20 km s$^{-1}$. By examining their 2D scaling relations and 3D fundamental manifold, we found that the sizes and velocity dispersions of these gap galaxies reveal continuous trends from dE to dSph galaxies. We conclude that low-luminosity dwarf elliptical galaxies are dominated by stars, not by dark matter, within their half light radii. %, in their inner regions. %as we found no evidence for elevated dynamical-to-stellarmass ratios. This finding can be understood if internal feedback processes are operating most efficiently in gap galaxies, gravitationally heating the centrally-located dark matter to larger radii. Whereas external environmental processes, which can strip away stars, have a greater influence on dSph galaxies resulting in their higher %Whereas additional mechanisms of %tidal and ram pressure stripping appear are required to account dark matter fractions. %, which are located within %the halos of the giant Milky Way or Andromeda galaxies, %With local environment playing a role in modifying %dark matter fractions, o UCDs appear to be more similar to massive compact star clusters than to small galaxies. Our dynamical study of low mass dwarf elliptical galaxies provides further constraints on the processes that shape some of the smallest and most numerous galaxies in the Universe. % ", "introduction": "% %In the last few years several new %galaxies have been found orbiting the Milky Way and %M31. Classified as dwarf spheriodal (dSph) galaxies they are extremely %faint but have been claimed to reside within massive halos of dark matter. %Thus although they contain few stars they are not puny galaxies. %Another subclass of dwarf galaxy is the dwarf %elliptical (dE). Unfortunately, the Local Group only possesses three %examples of this subclass -- all of which may be interacting %with M31. Studying dEs beyond the Local Group is very challenging %due to their faintness. %Indeed internal kinematics for dE less %luminous than M$_g$ = --16 are very rare. %It is unknown whether %these dwarfs reside in massive dark matter halos like their dSph %cousins. %Forbes et al. (2008) %highlighted the fact that dEs with $-16 > M_K > -18$ lacked %velocity dispersion measures and hence estimates of their %dynamical mass. %At these magnitudes, B--K %varies from $\\sim$ 2.5 to 2.0 %and hence this corresponds to $-14 > M_B > -15.5$. %The stellar systems that make up %lower luminosities ($M_K > -16$), are a collection %of objects, i.e. Ultra Compact Dwarfs (UCDs), nuclei of galaxies and %massive globular clusters (GCs). Internal velocity dispersion data %exists for such systems. Thus a gap exists in our knowledge %between the most massive UCDs/GCs and the smallest dE galaxies. %The ACSVCS imaged 100 early-type galaxies in the %Virgo cluster with the ACS to a magnitude limit of M$_g$ $\\sim$ %--15.5. With the superior resolution of the ACS, over 2/3 of %faint dE galaxies revealed nuclei which typically contribute %0.3\\% to the total galaxy luminosity. % ", "conclusions": "Using the ESI instrument on the Keck telescope we have obtained internal kinematics for four dwarf ellipticals and one late-type dwarf galaxy. These galaxies were selected to have little, or no, nuclear component and indeed there is no evidence for a significant kinematically distinct nucleus in our data. The galaxies have K band magnitudes down to M$_K$ $\\sim$ --17 mag which places them in the `gap' between previous dynamical mass measurements of dE galaxies by Geha et al. (2003) and Chilingarian (2009) and Local Group dwarf spheriodal (dSph) galaxies and Ultra Compact Dwarf (UCD) objects located in the Virgo and Fornax clusters. We define the dynamical mass gap to be between 8 $\\times$ 10$^{7}$ $<$ M$_{dyn}$/M$_{\\odot}$ $<$ 5 $\\times$ 10$^{8}$. %With the exception of the brightest dE galaxy, we find that they %are all pressure-supported in their inner regions with %relatively flat We measure central velocity dispersions of around 20 km s$^{-1}$ for each galaxy. Supplemented by data from the literature for dE, dSph and UCD objects, we derive total stellar masses (mostly from K band magnitudes) and dynamical masses (using the formulation of Wolf et al. 2010). We find that the dE galaxies in the mass gap suggest a continuity from dE to dSph galaxies in terms of both their central velocity dispersion and effective radius with dynamical mass. This is also true when these systems are examined in the 3D space of radius, dynamical mass and stellar mass. Such trends are qualitatively similar to those expected for dwarf galaxies from cosmological formation models. Interestingly, the dE galaxies in the gap reveal dynamical-to-stellar mass ratios, within their half light radii, of unity. Thus they appear to be stellar-dominated in their inner regions similar to higher mass dE galaxies and do not reveal the need for large dark matter fractions as inferred for Local Group dSph galaxies. We speculate that any dark matter in the inner regions of these low mass dE galaxies has been `puffed-up' to larger radii by gravitational heating effects from supernova feedback. Probing the dynamics of even lower mass dEs may reveal a population of dwarf galaxies that are dominated by stars in their inner regions." }, "1101/1101.5210_arXiv.txt": { "abstract": "We reconsider the pixel-based, ``template'' polarized foreground removal method within the context of a next-generation, low-noise, low-resolution (0.5 degree FWHM) space-borne experiment measuring the cosmological $B$-mode polarization signal in the cosmic microwave background (CMB). This method was first applied to polarized data by the {\\sl Wilkinson Microwave Anisotropy Probe} ({\\sl WMAP}) team and further studied by Efstathiou et al. We need at least 3 frequency channels: one is used for extracting the CMB signal, whereas the other two are used to estimate the spatial distribution of the polarized dust and synchrotron emission. No extra data from non-CMB experiments or models are used. We extract the tensor-to-scalar ratio ($r$) from simulated sky maps outside the standard polarization mask (P06) of {\\sl WMAP} consisting of CMB, noise ($2~\\mu$K~arcmin), and a foreground model, and find that, even for the simplest 3-frequency configuration with 60, 100, and 240~GHz, the residual bias in $r$ is as small as $\\Delta r\\approx 0.002$. This bias is dominated by the residual synchrotron emission due to spatial variations of the synchrotron spectral index. With an extended mask with $f_{sky}=0.5$, the bias is reduced further down to $<0.001$. ", "introduction": "\\label{sec:introduction} Why study the $B$-mode polarization of the cosmic microwave background (CMB)? Detection of the primordial gravitational waves generated during inflation would give us a direct insight into the physical condition of the universe when the energy scale was close to the grand unification scale, $\\sim 10^{16}$~GeV \\citep[see][for a recent review and references therein]{liddle/lyth:PDP}. While a direct detection of the primordial gravitational waves using, e.g., laser interferometers, seems not possible with the present-day technology, an {\\it indirect} detection using the $B$-mode polarization of the CMB \\citep{seljak/zaldarriaga:1997,kamionkowski/kosowsky/stebbins:1997} may be possible in the near future (most optimistically, within a few years), provided that the energy scale of inflation at which the observed gravitational waves were generated was indeed as high as the grand unification scale. We often characterize the amplitude of gravitational waves (also known as tensor perturbations) using the so-called ``tensor-to-scalar ratio,'' which is conventionally defined as \\begin{equation} r\\equiv \\frac{2\\langle |h^{+}_{{\\mathbf k}}|^2+|h^{\\times}_{{\\mathbf k}}|^2\\rangle}{\\langle |{\\cal R}_{\\mathbf{k}}|^2\\rangle}, \\end{equation} where $h^{+}_{\\mathbf k}$ and $h^{\\times}_{\\mathbf k}$ are the Fourier transform of the amplitudes of two linear polarization states of gravitational waves, and ${\\cal R}_{\\mathbf k}$ is the primordial curvature perturbation, which is a scalar perturbation (hence the name, ``tensor-to-scalar ratio''). It is ${\\cal R}_{\\mathbf k}$ that seeded the observed structure in the universe, as well as the dominant component of the observed CMB temperature anisotropy \\citep[see][for a recent review and references therein]{weinberg:COS}. The dominant, scalar part of the temperature anisotropy generates radial and tangential polarization patterns around hot and cold spots \\citep{coulson/crittenden/turok:1994}. This is called the $E$-mode polarization, and has been detected with high statistical significance \\citep{brown/etal:2009,chiang/etal:2010,larson/etal:prep,komatsu/etal:prep,quiet:prep}. However, the $B$-mode polarization, which cannot be generated by the scalar perturbations but can be generated by the tensor perturbations, has not been found yet. The current 95\\% upper limit on the tensor-to-scalar ratio is $r<0.24$, which mainly comes from the upper limit on the tensor contribution to the temperature anisotropy on large angular scales \\citep{komatsu/etal:prep}. \\begin{figure}[t] \\centering \\noindent \\includegraphics[width=8.5cm]{f01.ps} \\caption{% $E$-mode and $B$-mode polarization power spectra. The diamonds, triangles, stars, and squares show the {\\sl WMAP} seven-year data \\citep{larson/etal:prep}, the QUaD final data \\citep{brown/etal:2009}, the BICEP two-year data \\citep{chiang/etal:2010}, and the QUIET 43~GHz data \\citep{quiet:prep}, respectively. The upper solid line shows the scalar $E$-mode power spectrum of the {\\sl WMAP} seven-year best-fit model. The dashed lines show the primordial $B$-mode power spectra with the tensor-to-scalar ratio of $r=0.24$, which corresponds to the current 95\\% upper limit \\citep{komatsu/etal:prep}, as well as of $r=0.03$ and 0.003. These lines are linearly proportional to $r$. The dotted line shows the secondary $B$-mode power spectrum expected to be generated by the weak gravitational lensing effect converting $E$ modes to $B$ modes \\citep{zaldarriaga/seljak:1998}. This line is fixed (by the {\\sl WMAP} seven-year best-fit model) and acts as noise for the primordial $B$-mode detection. The lensing contribution becomes comparable to the primordial bump at $l=10$ and 100 for $r=0.003$ and 0.03, respectively. } \\label{fig:clnow} \\end{figure} Given the upper limit on $r$, one can calculate the expected level of the $B$-mode power spectrum (see Figure~\\ref{fig:clnow}). For $r=0.24$, the $B$-mode power spectrum is smaller than the $E$-mode power spectrum by a factor of 10 at the first bump (created by electrons at $z\\lesssim 10$). At the second bump (created by electrons at $z\\simeq 1090$), the $B$-mode power spectrum is smaller than the $E$-mode power spectrum by a factor of 50. It is the smallness of the $B$-mode power spectrum that makes the detection of this signal challenging. There are three sources of noise for $B$-mode detection: (1) Detector noise; (2) Galactic foreground emission; and (3) Gravitational lensing. In this paper, we shall focus on the Galactic foreground. We use a map-based method for reducing the Galactic foreground, and study how the residual foreground limits a measurement of the primordial $B$-mode polarization. The foreground reduction technique we use is motivated by the ``template cleaning method'' used by the {\\sl WMAP} team \\citep{page/etal:2007,gold/etal:2009,gold/etal:prep}. This method was further investigated by \\citet{efstathiou/gratton/paci:2009} in the context of the {\\sl Planck} mission. We shall study this technique in the context of a next-generation, low-noise, low-resolution (0.5 degree FWHM) space-borne experiment. There is a large body of literature on the issue of polarized foreground cleaning for the $B$-mode detection. Our method is one specific (and relatively simpler) example. For the other methods in the literature, see review articles \\citep{dunkley/etal:2008,fraisse/etal:prep} and references therein. This paper is organized as follows. In Section~\\ref{sec:noise}, we show how the detector noise and the lensing noise influence the statistical errors on $r$. In Section~\\ref{sec:method}, we describe our method for estimating $r$ in the presence of the Galactic foreground and the dominant scalar $E$-mode polarization. In Section~\\ref{sec:simulation}, we describe our simulation including CMB, detector noise, and foreground. In Section~\\ref{sec:results}, we present the main results of this paper. We conclude in Section~\\ref{sec:conclusion}. ", "conclusions": "\\label{sec:conclusion} In this paper, we have studied the pixel-based foreground cleaning method within the context of a next-generation, low-noise CMB polarization satellite. This method was originally applied to polarized data by the {\\sl WMAP} team \\citep{page/etal:2007,gold/etal:2009,gold/etal:prep}, and further investigated by \\citet{efstathiou/gratton/paci:2009} in the context of {\\sl Planck}. Despite the simplicity of the method (namely, we have maps at 3 different frequencies, two of which are used for removing the synchrotron and dust emission), we are able to recover the input tensor-to-scalar ratio with only a small bias, $\\Delta r\\approx 0.002(<0.001)$ for the P06(extended) mask, which is dominated by the residual synchrotron emission. Further improvements should be straightforward: one can tune the Galactic mask, and divide the synchrotron fitting regions according to the actual distribution of the synchrotron spectral index in the Galaxy (rather than using the regular division shown in Figure~\\ref{fig:alpha-regions}). One may also increase the number of frequencies for measuring the spatial distribution of the synchrotron spectral index, provided that we have enough space on the focal plane. These will be investigated in the context of specific experimental designs such as {\\sl LiteBIRD}\\footnote{Light satellite for the studies of B-mode polarization and Inflation from cosmic background Radiation Detection; http://cmb.kek.jp/litebird}, and presented elsewhere. Our study suggests that a detection of the primordial $B$-mode polarization at the level of $r\\approx 10^{-3}$ should be possible with carefully optimized mask and $\\alpha$ regions. Note that our statistical error and systematic bias becomes comparable with $f_{sky}=50\\%$ mask case. However, let us mention one important caveat in our analysis. While our knowledge of the distribution and properties of the polarized synchrotron is fairly secure thanks to the {\\sl WMAP} data, our knowledge of the polarized dust emission, especially the spatial variation of the dust spectral index, is still highly limited. Therefore, the estimated bias in $r$ that we have presented in this paper cannot be too accurate. Fortunately, {\\sl Planck} will soon provide us with maps of the polarized dust emission with the unprecedented sensitivity; thus, we intend to revisit this issue once the {\\sl Planck} data become available." }, "1101/1101.0329_arXiv.txt": { "abstract": "We present the analytic forms for the spectra of the cosmological perturbations from an initially anisotropic universe for the high momentum modes in the context of WKB approximations, as the continuation of the work \\cite{km}. We consider the Einstein gravity coupled to a light scalar field. We then assume that the scalar field has the zero velocity initially and then slowly rolls down on the potential toward the origin. In the slow-roll approximations, the Kasner-de Sitter universe with a planar symmetry is a good approximation as the background evolution. Quantization of the perturbations in the adiabatic vacuum, which we call the anisotropic vacuum, is carried out. For non-planar high momentum modes whose comoving momentum component orthogonal to the plane is bigger than the Hubble parameter at the inflationary phase, the WKB approximation is valid for the whole stage of the isotropization. On the other hand, the planar modes whose comoving momentum component orthogonal to the plane is comparable to the Hubble parameter, is amplified during the process of the anisotropic expansion. In the final gravitational wave spectra, we find that there is an asymmetry between the two polarizations of the gravitational wave because the initial mode mixing does not vanish. ", "introduction": "Recent measurements by the WMAP satellite \\cite{komatsu,wmap5,anomaly1} have suggested that the observed map of cosmic microwave background (CMB) anisotropy is almost consistent with the Gaussian and statistically isotropic primordial fluctuations from inflation. issues on a few anomalies in the CMB temperature map on large angular scales found in the recent data have been controversial. The most well-known fact is that there seems to be the suppression of the observed power of CMB anisotropy on angular scales bigger than sixty degrees \\cite{anomaly1}. There are other observational facts that imply the effect which induces the violation of the rotational invariance. More precisely, the planarity of lower multipole moments, the alignment between the quadrupole ($\\ell=2$) and the octopole ($\\ell=3$), and the alignment of them with the equinox and the ecliptic plane \\cite{anomaly2} were announced. There are other observational facts implying the large-scale anisotropy, i.e., odd correlations of $\\ell=4\\sim 8$ multipoles with $\\ell=2,3$ multipoles \\cite{copi}, a very large, possibly non-Gaussian cold spot in 10 degree scale \\cite{cruz}, asymmetry of angular map measured in north and south hemispheres \\cite{eriksen} (see also the recent review \\cite{review} and more recent references therein). It also should be noted that some authors claim that there is no significant evidence for primordial isotropy breaking in five-year WMAP data~\\cite{Picon} (see also more recent arguments \\cite{gawe}). Indeed, to explain the origin of the anomalies, various solutions have been suggested, introducing a nontrivial topology \\cite{topology}, a local anisotropy based on the Bianchi type VII${}_h$ universe to explain the quadrupole/octopole planarity and alignment \\cite{jaffe}, non-linear inhomogeneities \\cite{moffat} and assuming an elliptic universe to explain the suppression of the quadrupole CMB power \\cite{eli}. More recently, in particular, models which introduce an explicit source to break the spatial isotropy, either during inflation or in the late time universe, have been proposed, e.g., by the dynamics of an anisotropic energy-momentum component during inflation \\cite{ack,anomaly,ys,wks}, by the large scale magnetic field \\cite{vector}, by the anisotropic cosmological constant \\cite{cc} or dark energy \\cite{de}. The first purpose of this paper is to proceed to investigate the possibility that such large scale anomalies are produced by preinflationary anisotropy and obtain the leading order corrections to the spectra for them. Cosmic nohair theorem ensures that in the presence of a positive cosmological constant an initially anisotropic universe exponentially approaches the de Sitter spacetime at the later time under the strong or dominant energy condition \\cite{wald}. Therefore, it is plausible that the initial universe is highly anisotropic. The future CMB measurements will detect the fluctuations of B mode polarization in CMB, which may contain the information on the primordial gravitational waves. They would give a new tool to contrain the anisotropic cosmological model. The cosmological perturbation theory in the Kasner phase was formulated in Ref. \\cite{tpu,gcp,gkp}. In general in an expanding (planar) Kasner phase one of two polarizations of gravitational waves is coupled with the scalar mode, but the other gravitational mode is decoupled. Thus, this coupling induces the asymmetry between propagations of two polarizations of the gravitational waves. If there are effects of the chiral symmetry breaking in the cosmic history, they would give rise to nonzero cross correlations between the fluctuations of the temperature and B mode, and of E and B modes \\cite{lue}. They will give us powerful and independent tests on the primordial parity violation. The gravitational waves from the universe with an isotropy breaking would provide distinguishable signatures in the future CMB experiments. We will estimate how the initial mode mixing gives rise to the asymmetry between the primordial power spectra of the two gravitational wave modes, although we will not go into details of the observational aspects. Note that such subjects have been argued in the context of the anisotropic inflation models in Ref. \\cite{aniso_inf}. The investigation of the higher order correlations as the bispectrum would provide us another interesting prediction to examine the the anisotropic universe (see e.g., \\cite{chen}). We briefly comment on some expectations on this point in the last section. In the isotropic case, the quantization of fluctuations is carried out well inside the Hubble horizon, where the effects of the cosmic expansion can be ignored. In order to compare with the standard prediction, it is natural to quantize field in the initial adiabatic vacuum, which we call the {\\it anisotropic vacuum}. There are two branches of the expanding Kasner solution with the planar symmetry. The initial adiabatic vacuum present only on one of the two branches where the expansion rate along the planar directions vanishes while that along the special axis is finite \\footnote{ The anisotropic vacuum is not specific to the Bianchi I model. In Ref. \\cite{km}, it has been shown that an anisotropic vacuum can also be defined for a Bianchi IX model.}. As a result, the initial spacetime structure can be seen as the product of two-dimensional Milne spacetime and two-dimensional Euclidean space. The scalar fluctuations decouple from the tensor fluctuations at the very initial time, so that the initial dynamics reduces to that in a system composed of three independent harmonic oscillators. In the other branch there is an initial singularity. Since the coupling diverges at the initial time, we cannot find an adiabatic vacuum. Therefore, here we focus on the first branch. For a given set of initial conditions, the power spectrum was investigated, rather by the numerical ways in Ref. \\cite{tpu,gcp,gkp}. The aim of our study is to obtain more analytic understandings on the spectra from an initially anisotropic universe. Our previous work discussed the spectrum of a massless scalar field, ignoring its coupling with the metric perturbations \\cite{km}. In this work, as the continuation, we will discuss the metric perturbations, in particular focusing on the importance of the tensor-scalar coupling. The paper is constructed as follows: In Sec. II, the background solution of our anisotropic model is introduced. In Sec. III, we present the formulation of the coupled perturbations in the background of Kasner de Sitter solution with the planar symmetries and their relation to the cosmic observables. In Sec. IV, we investigate the behaviors of the perturbations modes after setting initial conditions in the anisotropic vacuum. In Sec. V, we close the article after giving a brief summary. ", "conclusions": "In this paper, we have analytically investigated the corrections to the power spectra of the cosmological perturbations due to the preinflationary anisotropy of the universe, for the high momentum modes in the context of WKB approximations. The first motivation to consider the anisotropic universe is that even if the present universe is almost isotropic, it does not mean that it is also isotropic from the beginning. It would be more generic that the initial universe is highly anisotropic. The second motivation comes from observations. In recent years, several groups have reported the so-called low-$\\ell$ anomalies in large angular power of CMB fluctuations. They may be produced by the breaking of the rotational symmetry in the early universe. We considered the Einstein gravity coupled to a light scalar field. We assumed that this scalar field initially stays at a very large (super-Planck) field value, and then starts to roll down slowly. If the mass of the scalar field is small enough, the kinetic energy of the scalar field does not affect the spacetime dynamics significantly. Imposing the regularity of the spacetime at the initial time, one of two planar branches of the Kasner-de Sitter solution, whose initial geometry becomes a (Milne) patch of the Minkowski spacetime, is a good approximation for describing the cosmic isotropization. The cosmic isotropization takes place within a few Hubble times. During the subsequent inflationary stage, the scalar field rolls down toward the true minimum and plays the role of the inflaton. Then, we investigated the analytic expressions for the spectra of the cosmological perturbations produced in the above anisotropic background for the high momentum modes. In the anisotropic background, there are 2 scalar and 1 vector modes in terms of the two-dimensional flat space. In the isotropic limit, this vector mode in the anisotropic universe reduces to one of tensor polarizations in the flat 3-dimensional space, while two scalar modes reduce to one scalar and the other tensor polarization there. During the anisotropic phase, two scalar modes are coupled, resulting in the asymmetry between the spectra of two tensor polarizations obtained in the isotropic limit. Since at the initial times, the coupling is absent, we could define the adiabatic vacuum, which we call the anisotropic vacuum, and canonically quantize the perturbations. Our anisotropic vacuum is definitely different from the standard Bunch-Davis vacuum. In addition, the presence of the anisotropic vacuum is specific to our particular choice of the Kasner parameter. For the non-planar modes $k,k_1\\gg H_0$, for the sufficiently high momentum the WKB approximation is valid but the mixing angle does not vanish at the early time. At the leading order, the power spectra of all the perturbation modes contain the corrections due to the nonstandard propagation in the anisotropic background, and the effects of the initial mixing of modes. The former has the universal form, while the latter induces an asymmetry of two gravitational wave polarizations in the isotropic limit. The modifications of spectra appear in the oscillatory behaviors of the primordial spectrum. On the other hand, for the planar mode, i.e., $k\\gg k_1 \\sim H_0$, although the WKB approximation is broken at the very early times, the mode mixing does not take place significantly. For the modes of $k_1=0$, the adiabaticity parameter initially diverges and hence the anisotropic vacuum is not well defined. However, such modes are not relevant for the observations. One of the main results which we have obtained is that, irrespective of the non-planar or planar modes, for the high momentum modes the ratio of the power spectra between two tensor polarizations is given by \\bea \\frac{P_{h_+}}{P_{h_\\times}}\\approx 1% -\\frac{2m}{\\sqrt{3}H_0}. \\eea In the chaotic inflation typically $\\frac{m}{H_0}=O(0.1)$, and hence there is a difference of the spectra of two gravitational wave polarizations, which is of roughly ten percent. Before closing this article, it may be important to mention the effects of the primordial anisotropy on the higher order spectra, in particular on the bispectrum, and the primordial non-Gaussianities. In the model discussed in this paper the initial vacuum is not the standard Bunch-Davis vacuum, and the subsequent evolution is almost the same as that in the single field, slow-roll inflation. Thus, we expect that the folded shape bispectrum where three momenta satisfy $k_1+ k_2 \\sim k_3$ would become dominant \\cite{chen,chks} (see also \\cite{isca}), while the local shape bispectrum of $k_2\\sim k_3\\gg k_1$ would be negligible \\cite{jm}. However, the particular local type bispectrum where $k_2$ and $k_3$ almost lie in the plane of the $y$ and $z$ directions of Eq. (\\ref{metric}) while $k_1$ is nearly orthogonal to this plane ($k_2\\sim k_3\\gg k_1$), could be much different from the case of the isotropic universe, since as in the case of the spectra for the planar modes it could be sensitive to the anisotropy. For the non-planar case of $k_1+ k_2 \\sim k_3$, the amount of the non-Gaussianities would be determined by the coefficient $B_{\\tilde V}$ in Eq. (\\ref{hiko}) which represents the amplitude of the negative frequency mode; \\bea {\\rm Re} (B_{\\tilde V})\\sim Q(r_2) \\Big(\\frac{H_0}{k}\\Big)^{\\frac{3}{2}} \\sin\\big(\\sqrt{\\frac{k}{H_0}}+\\phi_0\\big), \\eea where $\\phi_0$ denotes some constant phase. Note that due to the effect of the initial mode mixing ${\\rm Re} (B_{\\tilde V})$ may be amplified by some factor of order $\\frac{m}{H_0}\\sim 0.1$. Therefore, the bispectrum could exhibit an oscillatory behavior, and contain the information on the primordial anisotropy through the factor $Q(r_2)$, which may be distinguishable if detected. The concrete evaluation of the bispectrum in the anisotropic universe and the detectability of the non-Gaussianities will be interesting issues and should be investigated in the future studies." }, "1101/1101.0603_arXiv.txt": { "abstract": "We present single-epoch radio afterglow observations of $24$ long-duration gamma-ray burst (GRB) on a timescale of $\\gtrsim 100$ d after the burst. These observations trace the afterglow evolution when the blastwave has decelerated to mildly- or non-relativistic velocities and has roughly isotropized. We infer beaming-independent kinetic energies using the Sedov-Taylor self-similar solution, and find a median value for the sample of detected bursts of about $7\\times 10^{51}$ erg, with a $90\\%$ confidence range of $1.1\\times 10^{50}-3.3\\times 10^{53}$ erg. Both the median and $90\\%$ confidence range are somewhat larger than the results of multi-wavelength, multi-epoch afterglow modeling (including large beaming corrections), and the distribution of beaming-corrected $\\gamma$-ray energies. This is due to bursts in our sample with only a single-frequency observation for which we can only determine an upper bound on the peak of the synchrotron spectrum. This limitation leads to a wider range of allowed energies than for bursts with a well-measured spectral peak. Our study indicates that single-epoch centimeter-band observations covering the spectral peak on a timescale of $\\delta t\\sim 1$ yr can provide a robust estimate of the total kinetic energy distribution with a small investment of telescope time. The substantial increase in bandwidth of the EVLA (up to 8 GHz simultaneously with full coverage at $1-40$ GHz) will provide the opportunity to estimate the kinetic energy distribution of GRBs with only a few hours of data per burst. ", "introduction": "\\label{sec:intro} The energy budget of gamma-ray bursts (GRBs) provides fundamental insight into the nature of the explosions, the resulting ejecta properties, and the identity of the central compact remnant (``engine''). While the isotropic-equivalent $\\gamma$-ray energy ($E_{\\rm\\gamma,iso}$) can be easily determined from a measurement of the burst fluence and redshift, a complete accounting of the energy budget requires detailed observations of the afterglow emission. The afterglow observations provide a measure of the isotropic-equivalent blastwave kinetic energy ($E_{\\rm K,iso}$), as well as the explosion geometry (quantified by a jet opening angle, $\\theta_j$). The resulting beaming corrections, $f_b^{-1}\\equiv 1-{\\rm cos}(\\theta_j)$, can be substantial, approaching three orders of magnitude in some cases \\citep{fks+01,pk02,bkf03,bfk03}. To properly determine $E_{\\rm K,iso}$ and $f_b$ it is essential to observe the afterglows from radio to X-rays over timescales of hours to weeks, clearly a challenging task. This is particularly a problem for the subset of ``dark'' GRBs for which the lack of detected optical emission, or large extinction, prevent a determination of $E_{\\rm K,iso}$ and likely $f_b$ (e.g., \\citealt{bkb+02,pfg+02}). Over the past decade detailed afterglow observations have been obtained at a great cost of telescope time for about $20$ long-duration GRBs, with the basic result that the beaming corrections are large and diverse, leading to typical true energies of $E_\\gamma\\sim E_K\\sim 10^{51}$ erg \\citep{fks+01,pk01,pk02,bkf03,bkp+03,bfk03}. More recently, it has been recognized that some nearby long GRBs have much lower energies, $E_{\\rm iso}\\sim 10^{49}-10^{50}$ erg, and appear to be quasi-isotropic \\citep{kfw+98,skb+04,skn+06}. Similarly, some bursts appear to have large beaming-corrected energies of $\\sim 10^{52}$ erg \\citep{cfh+10,cfh+10b}. The existence of these highly energetic bursts depends at least in part on the ability to correctly infer their large beaming corrections. Indeed, the inference of jet opening angles from breaks in the afterglow light curves has become controversial in recent years due to conflicting trends in optical and X-ray light curves \\citep{lrz+08,rlb+09}. Similarly, in some cases a two-component jet has been inferred, with a narrow core dominating the $\\gamma$-ray emission and a wider component dominating the afterglow emission \\citep{bkp+03,rks+08}. Numerical simulations suggest that off-axis viewing angles can also lead to shallow breaks that may be missed or mis-interpreted \\citep{vzm10}. In addition to potential difficulties with the inference of $f_b$, the $\\gamma$-ray and kinetic energies measured from the early afterglow emission only pertain to the relativistic ejecta. The existence of a substantial component of mildly relativistic ejecta can only be determined from observations at late times when such putative material can refresh the forward shock. Clearly, the existence of substantial energy in a slow ejecta component will place crucial constraints on the activity lifetime of the central engine. Such late-time observations also have the added advantage that they probe the blastwave when it has decelerated to non-relativistic velocities and hence roughly approaches isotropy \\citep{fwk00,lw00}. This allows us to use the well-established Sedov-Taylor self-similar solution, with negligible beaming corrections, to estimate the total kinetic energy of both the decelerated ejecta and any additional initially non-relativistic material. Since the peak of the afterglow spectrum on these timescales is located in the radio band, the lack of optical afterglow emission (e.g., due to extinction) does not have an effect on the ability to determine $E_K$. This approach was first exploited by \\citet{fwk00} to model the late-time radio afterglow emission of GRB\\,970508 (at $\\delta t\\gtrsim 100$ d) from which the kinetic energy was inferred to be $E_K\\sim 5\\times 10^{50}$ erg. \\citet{bkf04} used the same approach to model the radio afterglow emission of GRB\\,980703 on timescales of $\\gtrsim 40$ d, and to re-model GRB\\,970508. They found kinetic energies of $E_K\\sim 3\\times 10^{51}$ erg for both bursts. Finally, \\citet{fsk+05} modeled the radio emission from GRB\\,030329 at $\\delta t\\gtrsim 50$ d and found $E_K\\sim 10^{51}$ erg. Only 3 bursts have been studied in this fashion so far because only those events have well-sampled radio light curves on the relevant timescales of $\\delta t\\gtrsim 100$ d. However, the kinetic energy can still be estimated using the same methodology even from fragmentary late-time radio observations. Such an approach will naturally result in larger uncertainties for each burst, but it can be applied to a much larger sample of events. Here we present such an analysis for $24$ long-duration GRBs with radio observations at $\\gtrsim 100$ d, but with only $1-3$ data points (at 1.4 to 8.5 GHz) per burst. Using these observations we infer robust ranges for the kinetic energy of each burst and for the population as a whole. The plan of the paper is as follows. The radio observations are summarized in \\S\\ref{sec:obs}. The model for synchrotron emission from a Sedov-Taylor blastwave, and the various assumptions we employ are presented in \\S\\ref{sec:model}. In \\S\\ref{sec:res} we detail the resulting kinetic energies and the range for the overall sample, and we compare these results to multi-wavelength analyses of early afterglows in \\S\\ref{sec:comp}. We conclude with a discussion of future prospects. ", "conclusions": "\\label{sec:comp} The key results of our analysis are that the median energy for the 11 bursts with self-consistent solutions is $E_K\\approx 7\\times 10^{51}$ erg, while the $90\\%$ confidence range is $1.1\\times 10^{50}-3.3\\times 10^{53}$ erg. The median value is about a factor of 3 times higher than previous calorimetric measurements for GRBs 970508, 980703, and 030329, for which energies of $3\\times 10^{51}$, $3\\times 10^{51}$, and $10^{51}$ erg, respectively, were determined \\citep{bkf04,fsk+05}. Similarly, the inferred energies are somewhat larger than the distributions of beaming-corrected $\\gamma$-ray and kinetic energies inferred from broad-band multi-epoch studies (Figure~\\ref{fig:ecomp}). From various such analyses, the median $\\gamma$-ray energy is $\\langle E_\\gamma\\rangle \\approx 8\\times 10^{50}$ erg \\citep{fks+01,bfk03,fb05}, while the median kinetic energy is $\\langle E_K\\rangle\\approx 5\\times 10^{50}$ erg (e.g., \\citealt{pk01,pk02,yhs+03}); see Figure~\\ref{fig:ecomp}. In both cases the $90\\%$ range spans about 2.5 orders of magnitude, somewhat narrower than our inferred $90\\%$ confidence range for $E_{\\rm ST}$. The extension to larger energies found in our analysis mainly reflects the lack of spectral peak determinations for the bursts with single-frequency observations (see Figure~\\ref{fig:hist}). These large energies can be generally eliminated with a measurement of the synchrotron peak in the GHz frequency range (e.g., Group A bursts; Figure~\\ref{fig:ecomp}). In the context of our results we note that recent numerical work by \\citet{zm09} led these authors to conclude that the timescale to reach isotropy is $\\sim 10^2$ yr rather than $\\sim 1$ yr as indicated by the analytic formulation of \\citet{lw00} which we follow here. As a result, they note that using the Sedov-Taylor formulation may lead to an erroneous estimate of the kinetic energy. However, inspection of the resulting potential disrepancies reveals that this effect is at most a factor of 2 {\\it as long as self-consistency between the inferred energy and density and the transition to the Sedov-Taylor phase is ensured} (see their Figure 10). The discrepancies become larger if the wrong timescale is assumed for the transition to non-relativistic expansion, but this quantity is not a free parameter. Indeed, our distributions of $\\beta$ values point to self-consistency for most bursts, and allow us to reject objects that are potentially still relativistic. Since the potential systematic uncertainty of about a factor of 2 is significantly smaller than the overall spread in allowed energy for each burst, we do not consider this to be an obstacle to our analysis, or to future work on the energy scale using late-time radio measurements. As clearly demonstrated in Figure~\\ref{fig:hist}, the most constrained energy determinations require a measurement of the synchrotron spectral peak (Group A); the absence of such a constraint requires additional assumptions about the circumburst density and results in a much wider energy range. Indeed, this is the key reason for the wider range of allowed high energy solutions ($\\gtrsim 10^{52}$ erg) compared to the results for $E_\\gamma$ and $E_K$ (Figure~\\ref{fig:ecomp}). Observations of GRBs 970508, 980703, and 030329 demonstrate that the spectral peak is typically located at $\\sim {\\rm few}$ GHz on a timescale of $\\sim 150$ d. Thus, observations in the $1-10$ GHz range on a timescale of $\\sim {\\rm few}$ hundred days should allow us to determine the peak flux and frequency. This will in turn provide an energy estimate with a similar level of precision to the results of early-time broad-band modeling. This is a fortuitous conclusion since with the full frequency coverage of the Expanded VLA (EVLA) it will soon be possible to cover this entire range in a few hours of observations to a sensitivity that is about an order of magnitude better than the VLA. As we demonstrated here, such a modest investment of observing time ($2-3$ hours per burst) can yield a robust estimate of the GRB energy distribution, {\\it regardless} of the ability to measure jet opening angles. Pursuing these observations for all bursts with a measured redshift will require only $\\sim 50-100$ hr of EVLA time per year. Indeed, with such observations we should be able to constrain the energy distribution to a comparable level as existing studies within a single year given that about 30 GRBs with known redshifts just from 2009 are now available for EVLA observations (a similar sample is available from 2008 bursts). In the longer term, the large number of objects will allow us to test the energy distribution as a function of redshift, at least over the range $z\\sim 1-3$ where the bulk of the detected bursts occur \\citep{bkf+05,jlf+06}. Similarly, this approach will be particularly useful for bursts that lack detailed optical or X-ray light curves due to observational constraints or dust extinction, and for bursts with controversial estimates of the jet opening angles." }, "1101/1101.3436_arXiv.txt": { "abstract": "We show how higher derivatives of the expansion factor can be developed into a null diagnostic for concordance cosmology ($\\Lambda$CDM). It is well known that the Statefinder -- the third derivative of the expansion factor written in dimensionless form, $a^{(3)}/aH^3$, equals unity for $\\Lambda$CDM. We generalize this result and demonstrate that the hierarchy, $a^{(n)}/aH^n$, % can be converted to a form that stays pegged at unity in concordance cosmology. This remarkable property of the Statefinder hierarchy enables it to be used as an extended null diagnostic for the cosmological constant. The Statefinder hierarchy combined with the growth rate of matter perturbations defines a {\\em composite null diagnostic} which can distinguish evolving dark energy from \\lcdm. ", "introduction": "Despite its radical connotations, there is mounting observational evidence in support of a universe that is currently accelerating \\cite{observations}. Theoretically there appear to be two distinct ways in which the universe % can be made to accelerate \\cite{DE_review}: (i) through the presence of an additional component in the matter sector, which, following \\cite{ss06}, we call {\\em physical dark energy}. Physical DE models possess large negative pressure and lead to the violation of the strong energy condition, $\\rho+3P \\geq 0$, which forms a necessary condition for achieving cosmic acceleration. Prominent examples of this class of models include the cosmological constant `$\\Lambda$', Quintessence, the Chaplygin gas, etc. (ii) The universe can also accelerate because of changes in the gravitational sector of the theory. These models (sometimes referred to as {\\em geometrical DE} or modified gravity) include $f(R)$ theories, extra-dimensional Braneworld models, etc. Due to its elegance and simplicity the cosmological constant, with $P = -\\rho$, occupies a privileged place in the burgeoning pantheon of DE models. Although the reasons behind the extremely small value of $\\Lambda$ remain unclear, concordance cosmology ($\\Lambda$CDM) does appear to provide a very good fit to current data (although possible departures from $P = -\\rho$ have also been noted \\cite{sss09}). Given the success and simplicity of concordance cosmology it is perhaps natural to discuss diagnostic measures which can be used to compare a given DE model with $\\Lambda$CDM. `Null measures' of concordance cosmology proposed so far include the $Om$ diagnostic \\cite{Om,zunkel_clarkson}, and the Statefinders \\cite{statefinder}. While $Om$ involves measurements of the expansion rate, $H(z)$, the Statefinders are related to the third derivative of the expansion factor. In a spatially flat $\\Lambda$CDM universe, the Statefinders and $Om$ remain pegged at a fixed value during our recent expansion history $(z \\lleq 10^3)$. In this letter we introduce the notion of the `{\\em Statefinder hierarchy}' which includes higher derivatives of the expansion factor $d^na/dt^n, n \\geq 2$. We demonstrate that, for concordance cosmology, all members of the Statefinder hierarchy can be expressed in terms of elementary functions of the deceleration parameter $q$ (equivalently the density parameter $\\om$). This property singles out the cosmological constant from evolving DE models and allows the Statefinder hierarchy to be used as an extended null diagnostic for $\\Lambda$CDM. ", "conclusions": "\\label{concl} In this paper we have shown that a simple series of relationships links the Statefinder hierarchy in concordance cosmology with the deceleration/density parameters. These relationships can be used to define {\\em null tests} for the cosmological constant. Including information pertaining to the growth rate of perturbations increases the effectiveness of this hierarchy of null diagnostics. Our results demonstrate that lower order members of the Statefinder hierarchy already differentiate quite well between concordance cosmology on the one hand, and Braneworld models and the Chaplygin gas, on the other. Since, for $n\\geq 3$, the $n^{\\rm th}$ Statefinder $S_n$ contains terms proportional to $w^{(n-2)}/H^{n-2}$, higher members of the hierarchy will contain progressively greater information about the evolution of the equation of state of DE. From the observational perspective, however, one might note that a determination of the $S_n$ statefinders involves prior knowledge of the $(n-1)^{\\rm th}$ derivative of $H(z)$. Thus only lower order Statefinders, $S_n$, $n \\leq 4$, together with the $Om$ diagnostic, may prove compatible with the quality of observational data expected in the near future." }, "1101/1101.1096_arXiv.txt": { "abstract": "Dust production among post-main-sequence stars is investigated in the Galactic globular cluster 47 Tucanae (NGC 104) based on infrared photometry and spectroscopy. We identify metallic iron grains as the probable dominant opacity source in these winds. Typical evolutionary timescales of AGB stars suggest the mass-loss rates we report are too high. We suggest that this is because the iron grains are small or elongated and/or that iron condenses more efficiently than at solar metallicity. Comparison to other works suggests metallic iron is observed to be more prevalent towards lower metallicities. The reasons for this are explored, but remain unclear. Meanwhile, the luminosity at which dusty mass loss begins is largely invariant with metallicity, but its presence correlates strongly with long-period variability. This suggests that the winds of low-mass stars have a significant driver that is not radiation pressure, but may be acoustic driving by pulsations. ", "introduction": "\\label{IntroSect} Stellar mass loss is of critical importance to the later stages of stellar evolution. Stars evolving up the red giant branch (RGB) lose mass via stellar winds. The remaining mass of a star leaving the RGB tip is the primary determinant of its position on the horizontal branch (HB; e.g.\\ \\citealt{Rood73,Catelan00}). Subsequent mass loss on the asymptotic giant branch (AGB) ejects the star's entire hydrogen envelope, creating a post-AGB star and (perhaps) a planetary nebula (PN). The timescale and end-point of AGB evolution may therefore not be determined by the nuclear burning rate, but by mass loss \\citep{vLGdK+99,BSvL+09}. Winds from these stars are the most significant producers of interstellar dust and therefore provided the heavy elements present in Population I stars, including the Sun and Solar System \\citep{Gehrz89,Sedlmayr94,Zinner03}. Debate exists around many finer points of the mass loss process, particularly its variation with fundamental parameters such as luminosity, temperature and metallicity. Firstly, the amount of mass loss that occurs on the RGB remains largely unmeasured. RGB mass loss appears to determine the `second parameter' (after metallicity) that defines the morphology of HBs in globular clusters (GCs) \\citep{FPB97,LC99,Catelan00}. Secondly, it is not proven where dust formation stars, and whether RGB stars can form significant amounts of dust in their ejecta. If many RGB stars do produce dust, it would add a large number of dust factories ejecting dust into a galaxy's interstellar medium (ISM) (cf.\\ the missing dust source in many galaxies; e.g.\\ \\citealt{MBZ+09}). \\citet{ORFF+07} suggest that dust production occurs episodically over a large portion of the RGB, while \\citet{MvLD+09} have presented evidence that dust production is confined to the RGB tip, though gaseous mass loss is still present all along both the RGB and AGB \\citep{MvL07,MAD09,DSS09}. Dust can also help drive a wind through radiation pressure on grains (e.g.\\ \\citealt{Lewis89}), though it is not certain whether winds which lack high-opacity amorphous carbon can be accelerated in this manner \\citep{Willson00,Woitke06b,Hoefner07,LZ08,MSZ+10}. Dustless stars typically lose mass via magneto-acoustic processes in their chromospheres \\citep{HM80,DHA84,MvL07}. Finally, the composition of circumstellar dust and the chemical processes by which it forms are only partly defined. Of particular interest is the r\\^{o}le that metallic iron grains play in dust formation, and how metallicity affects dust production and composition \\citep{MSZ+10}. In this study, we aim to address these points. To do this, we explore dust production in the Galactic globular cluster 47 Tucanae (NGC 104). Globular clusters are unique laboratories with which we can study late-stage evolution of low-mass stars: globular clusters contain stars that are among the oldest in the Universe, and most clusters are comprised of a single (or at least dominating) stellar population with little internal variance in age or chemical composition. Comparisons between clusters therefore probe differences in mass loss created by metallicity, while comparisons within clusters probe differences created by stellar temperature and luminosity. Throughout this paper, we will refer to similar studies: \\citet{MvLD+09}, which covers the cluster $\\omega$ Centauri (NGC 5139; hereafter Paper I); \\citet{BMvL+09}, which covers the cluster NGC 362 (hereafter Paper II); and a work which accompanies this paper (submitted ApJS; hereafter Paper III). The cluster 47 Tuc itself is among the Galaxy's most massive and, co-incidentally, one of the closest to us. Its distance has been estimated photometrically as 4190--4700 pc \\citep{Harris96,PSvWK02,SHO+07}\\footnote{The updated Harris cluster catalogue can be found at: http://www.physics.mcmaster.ca/\\protect\\~{}harris/mwgc.dat} and kinematically as 4000 pc \\citep{MAM06}, with 4500 pc being the median distance determination in the literature. Its mass is 6--9 $\\times 10^5$ M$_\\odot$ \\citep{SR75,MM85,MSS91}. Its metallicity is determined to be roughly [Fe/H] = --0.7, though various studies have placed it anywhere from [Fe/H] = --0.89 to --0.66 \\citep{Harris96,CG97,MB08,WCMvL10}. The cluster suffers from little interstellar reddening ($E(B-V) = 0.04$ mag --- \\citealt{Harris96}) and is estimated to be over 11 Gyr in age \\citep{CSD92,ZRO+01,GSA02,PSvWK02,GBC+03,KTR+07,SHO+07}. The cluster hosts an ionised intra-cluster medium (ICM; \\citealt{FKL+01}). This ionised medium is undetected in any other cluster (M15 has a neutral medium with a dusty component; \\citealt{ESvL+03,vLSEM06,BWvL+06}). Integrating the observed intracluster electron density over the cluster core yields $\\sim$0.1 M$_\\odot$ of H {\\sc ii}, representing $\\gtrsim$5\\% of the intracluster hydrogen \\citep{SWFW90,FKL+01}. The total stellar mass loss from 47 Tuc's stars therefore gives us insight into the fate of gas and dust expelled by stars in this cluster. Paper III more-precisely examines which stars have mid-infrared excess, indicative of dust, finding again that we cannot corroborate the presence of dust around the fainter giants. In this work, we use the data and results from Paper III to derive mass-loss rates for individual stars and (by summation) the cluster as a whole, and investigate the composition of the circumstellar dust produced by analysing archival mid-infrared (mid-IR) spectra. ", "conclusions": "\\label{SectConc} In this work, we have used spectral energy distributions and infrared spectroscopy to establish mass-loss rates and investigate mineralogy around the dust-producing stars in 47 Tuc. We summarise our conclusions as follows: \\begin{list}{\\labelitemi}{\\leftmargin=1em \\itemsep=0pt} \\item Mass loss in the cluster is dominated by its most luminous stars (V1, V8, V2, V3 and V4), however substantial mass loss appears to be taking place in moderate-luminosity stars too (x03, V18, V13). \\item The total mass-loss rate is estimated to be $\\sim$1.2 $\\times$ 10$^{-5}$ M$_\\odot$ yr$^{-1}$, mainly in the form of metallic iron, but this is probably over-estimated by a factor of 4--8 due to unconstrained properties of the dust grains. \\item It remains unclear whether radiation pressure on dust is capable of driving a wind alone, due to the low projected terminal velocities of the stellar winds. The existence of circumstellar dust appears to correlate with the presence of stellar pulsation. We theorise that such pulsation may inject energy into the wind acoustically. \\item A dust-to-gas ratio closer to unity and smaller or more needle-like iron grains are suggested as methods to decrease the total mass-loss rate while allowing for an increase in the terminal velocity. \\item Variations in the wind chemistry broadly concur with those found in \\citet{MSZ+10}, with silicates found preferentially in high-luminosity, high-metallicity winds, while iron is found preferentially in low-luminosity, low-metallicity winds. Iron may condense preferentially at low metallicity. Alternatively, iron grains may be coated by silicates in dense, metal-rich winds, but this coating process may not be effective in winds around lower-luminosity and metallicity stars. \\end{list} \\vspace{5 mm} \\noindent {\\bf Acknowledgments:} We are grateful to Thomas Lebzelter for sharing his original data with us. This paper uses observations made using the \\emph{Spitzer Space Telescope} (operated by JPL, California Institute of Technology under NASA contract 1407 and supported by NASA through JPL (contract number 1257184)); observations using \\emph{AKARI}, a JAXA project with the participation of ESA; and data products from the Two Microns All Sky Survey, which is a joint project of the University of Massachusetts and IPAC/CIT, funded by NASA and the NSF." }, "1101/1101.0165_arXiv.txt": { "abstract": "We present deep spectroscopic observations of the classical T Tauri stars DF Tau and V4046 Sgr in order to better characterize two important sources of far-ultraviolet continuum emission in protoplanetary disks. These new {\\it Hubble Space Telescope}-Cosmic Origins Spectrograph observations reveal a combination of line and continuum emission from collisionally excited H$_{2}$ and emission from accretion shocks. H$_{2}$ is the dominant emission in the 1400~$\\lesssim$~$\\lambda$~$\\lesssim$~1650~\\AA\\ band spectrum of V4046 Sgr, while an accretion continuum contributes strongly across the far-ultraviolet spectrum of DF Tau. We compare the spectrum of V4046 Sgr to models of electron-impact induced H$_{2}$ emission to constrain the physical properties of the emitting region, after making corrections for attenuation within the disk. We find reasonable agreement with the broad spectral characteristics of the H$_{2}$ model, implying $N(H_{2})$~$\\sim$~10$^{18}$ cm$^{-2}$, $T(H_{2})$~=~3000$^{+1000}_{-500}$ K, and a characteristic electron energy in the range of $\\sim$ 50~--~100 eV. We propose that self-absorption and hydrocarbons provide the dominant attenuation for H$_{2}$ line photons originating within the disk. For both DF Tau and V4046 Sgr, we find that a linear fit to the far-UV data can reproduce near-UV/optical accretion spectra. We discuss outstanding issues concerning how these processes operate in protostellar/protoplanetary disks, including the effective temperature and absolute strength of the radiation field in low-mass protoplanetary environments. We find that the 912~--~2000~\\AA\\ continuum in low-mass systems has an effective temperature of $\\sim$~10$^{4}$~K with fluxes 10$^{5-7}$ times the interstellar level at 1 AU. ", "introduction": "Classical T Tauri stars (CTTSs) are characterized by broad H$\\alpha$ emission lines and ultraviolet (UV) spectra dominated by atomic and molecular features that are attributed to gas-rich disks~\\citep{furlan06}. The ages of CTTS disks (0.1~--~12~Myr; Isella et al. 2009; Kastner et al. 2008) indicate that the gaseous processes contributing to their observational characteristics are contemporaneous with and likely intimately connected with giant planet formation, which is thought to be mostly completed on similar timescales~\\citep{alibert05}.\\nocite{isella09,kastner08} Terrestrial planet formation is thought to occur on timescales of 10~--~100 Myr~\\citep{kenyon06}, when the majority of the primordial gas is in the disk has been dissipated. Observations of the gas and dust in CTTS disks therefore probe the physical and chemical state of gas giant forming protoplanetary systems. Multi-wavelength spectroscopy is a powerful tool for making quantitative measurements of the dust and gas in these systems. Dust in protoplanetary disks is seen most clearly through the mid- and far-infrared (IR) excess flux produced by warm grains~\\citep{furlan06}, which overwhelms the narrow molecular lines when observed at low-resolution~\\citep{najita10}. Emission from hot gas is produced by magnetic activity in the atmospheres of the central stars and the accretion shocks near the stellar surface~\\citep{krull00,ardila02}, while the protoplanetary material itself can be probed through molecular observations of these systems. Protoplanetary disks have a multi-phase physical structure (see, e.g., the reviews by Woitke et al. 2009 and Dullemond \\& Monnier 2010)\\nocite{woitke09,dullemond10} seen in molecular line observations from the far-UV to the millimeter. \\citet{dutrey01} reviewed the mm-wave CO measurements of outer protoplanetary disks, but understanding the complex CO excitation structure with radius requires additional tracers of the bulk molecular material within 100 AU of the star~\\citep{greaves04}. Emission from OH, CO, CO$_{2}$, H$_{2}$O, and other biologically important species have been used to trace warm gas ($T$~$\\sim$~few~$\\times$~10$^{2}$~--~2000 K) in the inner disk~\\citep{najita03,salyk08,bethell09}. Metal-bearing molecules trace molecular hydrogen (H$_{2}$), the primary constituent of gas giant planets. However, H$_{2}$ can be observed best in the far-UV (912~--~1650~\\AA) bandpass, where the dipole-allowed molecular emission spectrum is primarily photo-excited (``pumped'') by stellar Ly$\\alpha$ photons~\\citep{ardila02,herczeg02}. The Ly$\\alpha$-pumping route requires that the second excited vibrational level ($v$~=~2) of H$_{2}$ have an appreciable population, which requires that the molecules reside in a warm ($T(H_{2})$~$>$~2000 K) gas layer close to the star~\\citep{herczeg06}. Most likely, this photoexcited H$_{2}$ is not physically associated with the bulk of the colder H$_{2}$ that goes into planet formation. The homonuclear nature of H$_{2}$ means that rovibrational transitions are dipole forbidden, making direct detection of the cool-H$_{2}$ component observationally challenging. Electron-impact excitation has been suggested as a means of directly probing the H$_{2}$ in the planet-forming regions of the disk~\\citep{bergin04}. This mechanism requires stellar X-rays to create a distribution of photoelectrons that pump the molecules into excited electronic states. The H$_{2}$ then relaxes, producing a characteristic far-UV cascade spectrum of discrete emission lines and quasi-continuous spectral features during the dissociation of H$_{2}$~\\citep{ajello84,abgrall97}. \\citet{ingleby09} proposed this mechanism to explain a portion of the low-resolution far-UV spectra from a sample of classical and weak-lined (low gas content) T Tauri stars observed with the {\\it Hubble Space Telescope}-ACS and -STIS. Low-resolution and intermediate S/N in existing far-UV observations make it difficult to separate the electron-impact H$_{2}$ signal from Ly$\\alpha$-pumped H$_{2}$, atomic emission, CO, and the underlying accretion continuum. We lump these emissions together under the term ``continuum'' due to their mostly unresolved structure in previous studies. Fortunately, the observational picture has improved dramatically with the installation of the Cosmic Origins Spectrograph (COS) on $HST$. $HST$-COS combines very low detector backgrounds with moderate spectral resolution ($\\Delta$$v$~$\\approx$~17 km s$^{-1}$). The low background permits measurements of the true continuum shape, including an assessment of the contribution from a hot accretion continuum. The resolution of COS also permits the identification and separation of the spectral components, enabling a more robust comparison with models of the processes that govern the planet-forming regions of CTTS disks. In this paper, we use COS to study in detail the far-UV continuum emission from these objects for the first time. We focus on two objects that we propose are prototypes for electron-impact H$_{2}$ emission (V4046 Sgr) and continuum emission produced in the hot accretion shock (DF Tau). CO $A$~--~$X$ band emission also contributes to the far-UV line and continuum spectrum in these systems. First results on CO emission will be presented in Paper II~(K. France~--~in preparation). In \\S2, we describe the COS observations and data reduction. In \\S3, we describe empirical fits to the data and compare these fits with model H$_{2}$ spectra. \\S4 describes our measurements of the collisionally excited H$_{2}$ and accretion continua. We also discuss discrepancies between the new COS observations and predictions of our electron-impact H$_{2}$ models. We then discuss in \\S5 the relevance of these processes to the physical state of the disk and the local far-UV radiation field. Finally, we present a brief summary of our work in \\S6. ", "conclusions": "\\subsection{Electron-Impact H$_{2}$ Emission in V4046 Sgr} In \\S3.3, we described how the $\\lambda$~$<$~1400~\\AA\\ discrete emission lines from the electron-impact H$_{2}$ excitation cascade are absorbed and/or scattered out of the line-of-sight. In order to measure the properties of the collisionally excited H$_{2}$ spectrum, we performed a $\\chi^{2}$ minimization analysis to compare the binned V4046 Sgr data at $\\lambda$~$>$~1400~\\AA\\ with the grid of electron-impact H$_{2}$ models described in \\S3.2. The free parameters are [$N$(H$_{2}$), $T$(H$_{2}$), $E_{e}$] and a linear fit to the underlying continuum, which we will discuss in the following subsection. We find that the electron-impact excited H$_{2}$ has a temperature of $T(H_{2})$~=~3000$^{+1000}_{-500}$ K, a column density of $N$(H$_{2}$)~$\\sim$~10$^{18}$ cm$^{-2}$, an electron energy of $E_{e}$~$\\sim$~50~--~100~eV, but the only firm constraint is on the temperature. By restricting the fits to $\\lambda$~$>$~1400~\\AA, we avoid the 1200~--~1300~\\AA\\ region where our models predict a dependence on $N$(H$_{2}$). We find a minimum near $E_{e}$~$\\sim$~50~--~100 eV, but higher energies are generally consistent with the data for $N$(H$_{2}$)~=~10$^{16-19}$ cm$^{-2}$. In Figure 7, we show representative fits to the $\\lambda$~$>$~1400~\\AA\\ data. We find that the best fit parameters not only display the lowest reduced $\\chi^{2}$, but they also show reasonable fits to the eye. For comparison, we also display a low energy case ($E_{e}$~=~14 eV). \\begin{figure} \\begin{center} \\hspace{-0.25in} \\epsfig{figure=f7.eps,width=2.65in,angle=90} \\caption{\\label{cosovly} The best-fit model of the electron-impact excited H$_{2}$ emission in the spectrum of V4046 Sgr. Due to the attenuation of the short-wavelength discrete emission, we restrict the fits to $\\lambda$~$>$~1400~\\AA. The best fit model parameters are $T(H_{2})$~=~3000$^{+1000}_{-500}$ K, N($H_{2}$) $\\sim$~10$^{18}$ cm$^{-2}$, and electron energies in the range $E_{e}$~$\\sim$~50~--~100 eV. The H$_{2}$ column density is only weakly constrained by this analysis. We also show the low energy model ($E_{e}$~=~14 eV) for the best fit temperature and column density. The underprediction of the model flux near 1600~\\AA\\ is due to not including CO emission from the disk.} \\end{center} \\end{figure} A major complication to all of these fits is the spectral overlap between the electron-impact emission spectrum of H$_{2}$ and CO $A$~--~$X$ band emission. With our broad binned approach, we can avoid strong discrete lines from the stellar atmosphere and accretion shock, as well as photoexcited H$_{2}$. The CO emission is distributed in wide vibrational band structures with non-zero intraband flux. Using the binned continuum bands described above, we necessarily include CO emission. We will present models of this emission in Paper II. For now, we note that the strongest contaminant to the pure H$_{2}$ electron-impact spectrum is the CO (0~--~1) $\\lambda$1597 \\AA\\ band that fills in the spectral valley between the H$_{2}$ emission peaks at 1575 and 1608~\\AA. This excess emission tends to push the fits to higher temperatures, where the dissociation continuum broadens, or to very low electron energies ($E_{e}$~$<$~14 eV), where dissociation is less important. This may explain some of the discrepancies between our results and the best-fit parameters [$T(H_{2})$~$\\sim$~5000 K; $E_{e}$~$\\sim$~12 eV] presented by~\\citet{ingleby09}. \\subsection{Accretion Continuum} Figure 3 shows the spectra of our targets from the far-UV to the optical. The black dotted optical spectrum has been scaled from an archival observation to show that the linear fit to the far-UV data alone (dashed red line) can plausibly also describe the shape of the $\\lambda$~$<$~3600~\\AA\\ Balmer continuum used to measure the accretion rates in low-mass pre-main sequence stars with disks~\\citep{calvet98,herczeg09}. We fit the far-UV spectra with both a power-law and a straight line. For both spectra a simple linear fit of the form $F_{\\lambda}$~=~$m$$\\lambda$~+~$b$ produced the best match to the data. We find [$m$,$b$] = [0.0393, -44.0250] and [0.0026, 0.0918] for DF Tau and V4046 Sgr, respectively, where the units of $m$ and $b$ are FEFU \\AA$^{-1}$ and FEFU, where 1 FEFU~=~1 $\\times$~10$^{-15}$ erg cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$. The linear slope for DF Tau is $\\sim$~15 times that found for V4046 Sgr. While DF Tau is known to have a much larger accretion rate (2~$\\lesssim$~$\\dot{M}$$_{acc}$~$\\lesssim$~1300~$\\times$~10$^{-9}$ $M_{\\odot}$ yr$^{-1}$; Hartigan et al. 1995; Herczeg et al. 2006; Herczeg \\& Hillenbrand 2008)\\nocite{hartigan95,herczeg06,herczeg08} than V4046 Sgr ($\\dot{M}$$_{acc}$~$\\sim$~3~$\\times$~10$^{-11}$ $M_{\\odot}$ yr$^{-1}$; G{\\\"u}nther and Schmitt 2007)\\nocite{gunther07} a direct comparison between accretion rates from the literature and the far-UV accretion continuum slopes is unreliable given the different measurement methods and temporal variability. We can estimate the mass accretion rates directly from the COS observations using empirical scaling relations between the \\ion{C}{4} luminosity ($L_{CIV}$) and $\\dot{M}$$_{acc}$. The largest systematic source of uncertainty to this approach is the relative contribution from the stellar transition region emission to the observed \\ion{C}{4} profile. We use equation 2 from~\\citet{krull00} to estimate the mass accretion rate in these systems at the time of our continuum observations, assuming that the \\ion{C}{4} surface flux in excess of a saturated stellar component ($F_{CIV}$~$>$~10$^{6}$ erg cm$^{-2}$ s$^{-1}$) is produced at the accretion shock. We fit the \\ion{C}{4} profiles with a multi-component line-profile employing the appropriate COS LSF. Assuming distances of 140 and 70 pc for DF Tau and V4046 Sgr, respectively, we find total $L_{CIV}$~=~1.81~$\\pm$~0.09~$\\times$~10$^{30}$ erg s$^{-1}$ and 1.76~$\\pm$~0.04~$\\times$~10$^{29}$ erg s$^{-1}$. The error bars are the measurement errors, which are very small for these bright lines. These are the time-averaged luminosities over the total G160M observing time. The \\ion{C}{4} fluxes were found to be slowly varying over the $\\approx$ 3 hours of G160M exposures. We measure \\ion{C}{4} count rate changes of +7.0 and -13.6\\% for DF Tau and V4046 Sgr from the COS microchannel plate two-dimensional spectrograms, evaluated in 50s time intervals. These \\ion{C}{4} luminosities can be explained by mass accretion rates of 7.8~$\\times$~10$^{-8}$ $M_{\\odot}$ yr$^{-1}$ and 1.3~$\\times$~10$^{-8}$ $M_{\\odot}$ yr$^{-1}$ for DF Tau and V4046 Sgr, respectively. While our accretion rate for DF Tau is only slightly larger than the range quoted by~\\citet{herczeg08}, our measurement for V4046 Sgr is considerably larger than the value quoted by~\\citet{gunther07}. The large discrepancy with the~\\citet{gunther07} accretion rate may be attributable to a larger contribution from transition region \\ion{C}{4} in V4046 Sgr, however it seems unlikely that values of the mass accretion rate as low as $\\dot{M}$$_{acc}$~$\\sim$~3~$\\times$~10$^{-11}$ $M_{\\odot}$ yr$^{-1}$ are consistent with our observations of V4046. This is not entirely surprising as CTTS mass accretion rates based on X-ray observations typically differ from optical/UV determinations by more than an order of magnitude due to absorption in the high-density regions where shock X-rays are produced (see e.g., Sacco et al. 2010 and references therein).\\nocite{sacco10} The relevant accretion measurements are given in Table 2. We suggest that future observations for which far-UV, near-UV, and optical observations are obtained in close time proximity are required to measure accurately both the far-UV continuum and the accretion rate from the Balmer continuum. These observations will tie together the observed far-UV accretion continuum in these objects with the well-calibrated accretion rate diagnostic. \\subsection{Outstanding Issues Regarding Electron-Impact Excitation of H$_{2}$} There is some uncertainty in the assumption that electron-excited H$_{2}$ is the dominant molecular continuum emission mechanism for V4046 Sgr and other young circumstellar disks. Our models and analysis support a picture in which intermediate energy electrons excite H$_{2}$ in the inner disk, which then radiates and is observed superimposed upon an underlying continuum from the accretion shock. Since X-rays are invoked to create the photoelectrons~\\citep{draine78} needed to collisionally excite the H$_{2}$, we must ask if X-rays are sufficient for this mechanism to operate. We can measure the total flux from electron-excited H$_{2}$ in V4046 Sgr, $F_{H2}$~$\\sim$ 8~$\\pm$~4~$\\times$~10$^{-13}$ erg cm$^{-2}$ s$^{-1}$ (approximately 2~$\\times$ the integrated 1400~--~1660~\\AA\\ H$_{2}$ flux shown in Figure 7 in order to account for the implied $\\lambda$~$<$~1400~\\AA\\ emission, \\S3.3). If we assume an inner disk electron-impact H$_{2}$ emitting region (0.2~$\\leq$~$a_{H2}$~$\\leq$~0.5 AU; see \\S5.2 for additional discussion), the collisionally excited H$_{2}$ luminosity is $L^{cont}_{H2}$~$\\sim$~9~$\\pm$~5~$\\times$~10$^{28}$ erg s$^{-1}$, or % $\\sim$~6~$\\pm$~3~$\\times$~10$^{39}$ photons s$^{-1}$ for a fiducial H$_{2}$ wavelength of 1400~\\AA. High resolution $Chandra$ spectra of V4046 Sgr show that the accretion shock flux is approximately half of the observed value in TW Hya~\\citep{gunther06,gudel07}. Scaling from the observed 0.45~--~6.0 keV X-ray luminosity ($L_{X}$) of TW Hya~\\citep{kastner02}, we estimate the total X-ray luminosity for V4046 Sgr is $L_{X}$~$\\sim$~1~--~5~$\\times$~10$^{30}$ erg s$^{-1}$, or 0.6~--~3~$\\times$~10$^{39}$ photons s$^{-1}$ for a fiducial X-ray energy of 1 keV. 1 keV is the approximate peak energy of the V4046 Sgr X-ray spectral energy distribution~\\citep{gunther07}. The grain photoelectric yields for the sub-micron sized grains that are present at the inner edge of transitional dust disks~\\citep{eisner06} are of order unity for photons with $E$~$\\gtrsim$~1~keV~\\citep{weingartner06}. Thus, we find that V4046 Sgr produces sufficient X-ray luminosity to create enough photoelectrons to excite the observed H$_{2}$ emission from the inner disk. While this electron-impact H$_{2}$ interpretation is energetically plausible for the observed emission, it may not be adequate to describe the full far-UV emission spectrum of V4046 Sgr. Absorption by a hydrocarbon screen appears to reproduce both the observed absorption of the short-wavelength discrete lines and the unobscured H$_{2}$ emission at $\\lambda$~$\\gtrsim$~1400~\\AA, but verification of this result requires additional study. A more speculative scenario involves a very high column density, high temperature self-absorbing screen of H$_{2}$ that could produce nearly continuous absorption in the COS G130M band while not absorbing the long wavelength emission arising from dissociative transitions. Another possible scenario is a ``non-traditional'' electron-impact H$_{2}$ emission process in which dissociative yields of order unity are achieved. This process would produce the quasi-continuous long wavelength emission without most of the energy being distributed among the discrete states of the Lyman and Werner bands. As the canonical H$_{2}$ dissociation fraction from the Lyman and Werner bands is $\\sim$~0.15~\\citep{stecher67,shull82,abgrall97}, large dissociative yields would require substantially fine-tuned molecular and electron distributions. It seems unlikely that either of these scenarios are physically possible. The other prominent long-wavelength emission source in these disks is CO, although CO alone probably cannot account for the observed far-UV molecular spectra of CTTS disks. We point out that a spectrograph such as COS is essential for a comprehensive analysis of such systems. The low background and high-sensitivity are both important because the continua in these systems have characteristic flux levels of a few~$\\times$~10$^{-15}$ erg cm$^{-2}$ s$^{-1}$ \\AA$^{-1}$ or less. Moderate spectral resolution is also required for the detection of individual discrete lines predicted by electron-impact models. Moreover, when protoplanetary systems are observed $without$ heavy internal reddening, spectral coverage at $\\lambda$~$<$~1300~\\AA\\ is important because this region probes the discrete line emission from the Werner bands of H$_{2}$ whose spectrum is a more sensitive diagnostic of the electron energy distribution. Without this wavelength region, spectral fits to the data will not account for the Werner lines, biasing the results toward very low energy electron distributions. As the threshold energies for the Lyman and Werner bands are 11.37 and 12.41 eV~\\citep{ajello84}, respectively, Lyman band continuum emission in the absence of discrete Werner band lines requires a highly-tuned electron distribution. Because Ly$\\alpha$ production from the dissociation of H$_{2}$ drops sharply at 14.7eV~\\citep{ajello91}, discrete lines should be even more important relative to the continuum at very low electron energies. \\citet{herczeg06} presented high-quality $HST$-STIS E140M observations of DF Tau. While the spectral resolution of these data is sufficient to analyze the brighter Ly$\\alpha$ pumped H$_{2}$ emission, the accretion and H$_{2}$ continua are too confused by the instrumental noise floor to make a clear separation. Conversely, \\citet{ingleby09} presented a large sample of CTTS targets using the ACS-SBC, but the low spectral resolution makes the detection of discrete molecular features impossible. Neither of these instruments has sufficient throughput in the $\\lambda$~$<$~1300~\\AA\\ band to measure the Werner emission lines of H$_{2}$. We have presented here the highest-quality observations of the far-UV continuum in CTTSs acquired to date. Because the gas-rich CTTS phase is roughly co-temporal with the epoch of giant planet formation, our observations represent an important portion of the energetic radiation environment present during the formation of exoplanetary systems. We characterize the spectra of hot accretion emission and electron-excited H$_{2}$ in two prototypical CTTSs. We find that the far-UV accretion continuum can be described by a simple linear fit, which connects to the optical Balmer continuum given assumptions about the optical activity level at the time of the far-UV observations. Accretion introduces an additional source of far-UV continuum emission that is 5~--~7 orders of magnitude stronger than the ISRF and has an effective temperature of $\\sim$~10$^{4}$ K in DF Tau. We also find that the electron-impact H$_{2}$ spectrum in V4046 Sgr is characterized by $T$(H$_{2}$)~=~3000 K and $N$(H$_{2}$)~$\\sim$~10$^{18}$ cm$^{-2}$, with $E_{e}$~$\\sim$~50~--~100 eV. Finally, we propose that hydrocarbon absorption is the most likely explanation for the differential reddening inferred from the attenuation of collisionally produced H$_{2}$ emission lines below 1400~\\AA." }, "1101/1101.0679_arXiv.txt": { "abstract": "% Gnu Data Language (GDL) is an open-source interpreted language aimed at numerical data analysis and visualisation. It is a free implementation of the Interactive Data Language (IDL) widely used in Astronomy. GDL has a full syntax compatibility with IDL, and includes a large set of library routines targeting advanced matrix manipulation, plotting, time-series and image analysis, mapping, and data input/output including numerous scientific data formats. We will present the current status of the project, the key accomplishments, and the weaknesses - areas where contributions are welcome! ", "introduction": "GDL is written in C++ and can be compiled on systems with GCC ($\\ge$ 3.4) and X11 or equivalents. The code, under \\htmladdnormallink{GNU GPL}{http://www.gnu.org/licenses/gpl.html}, is hosted by \\htmladdnormallink{SourceForge}{http://sf.net/projects/gnudatalanguage/}. The library routines make use of numerous open-source libraries including: \\htmladdnormallink{readline}{http://tiswww.case.edu/php/chet/readline/rltop.html}, the GNU Scientific Library (\\htmladdnormallink{GSL}{http://www.gnu.org/software/gsl/}), the \\htmladdnormallink{PLplot plotting library}{http://plplot.sourceforge.net/}, a Fourier transform package (\\htmladdnormallink{FFTw}{http://www.fftw.org/}), and others. Since recently (GDL 0.9rc4 release) GDL features multi-threaded matrix operations if compiled using an OpenMP-enabled compiler (e.g. GCC $\\ge$ 4.2.) PLplot and GSL are the only mandatory dependencies. Data input/output is managed using \\htmladdnormallink{ImageMagick}{http://www.imagemagick.org/}, \\htmladdnormallink{NetCDF}{http://www.unidata.ucar.edu/software/netcdf/}, \\htmladdnormallink{HDF}{http://hdf4.org/products/hdf4/} and \\htmladdnormallink{HDF5}{http://hdf4.org/HDF5/} libraries. FITS files can be read and written using the \\htmladdnormallink{Astron Library}{http://idlastro.gsfc.nasa.gov/}. GDL features a Python bridge (Python code can be called from GDL, GDL can be compiled as a Python module). ", "conclusions": "The core components of GDL (i.e. interpreter, library routines API, key data manipulation and plotting functionality) are stable and do not pose efficiency problems (no significant discrepancy from IDL performance). Large number of routines are available, several widely used external libraries (Astron, MPfit) can be used in GDL. We hope to consolidate the users community, to gather feedback in form of bug reports, feature requests, test routines, documentation and patches (several GDL modules have been provided by scientists who wrote the functions for their own work). The current status is stable and complete enough for numerous applications but still much work is needed. The major axis on development for the next year are: \\begin{itemize} \\item aggregating a more efficient community \\item giving pre-compiled versions in major Linux distributions closer to the CVS version \\item finishing a whole test suite (like the GSL one) to avoid regression and bugs \\item having a wider set of graphical keywords \\item better Postscript output (or other format(s) for graphical outputs) for publications \\item development of documentation \\end{itemize}" }, "1101/1101.0812_arXiv.txt": { "abstract": "Clusters of galaxies have long been used as laboratories for the study of galaxy evolution, but despite intense, recent interest in feedback between AGNs and their hosts, the impact of environment on these relationships remains poorly constrained. We present results from a study of AGNs and their host galaxies found in low-redshift galaxy clusters. We fit model spectral energy distributions (SEDs) to the combined visible and mid-infrared (MIR) photometry of cluster members and use these model SEDs to determine stellar masses and star-formation rates (SFRs). We identify two populations of AGNs, the first based on their X-ray luminosities (X-ray AGNs) and the second based on the presence of a significant AGN component in their model SEDs (IR AGNs). We find that the two AGN populations are nearly disjoint; only 8 out of 44 AGNs are identified with both techniques. We further find that IR AGNs are hosted by galaxies with similar masses and SFRs but higher specific SFRs (sSFRs) than X-ray AGN hosts. The relationship between AGN accretion and host star-formation in cluster AGN hosts shows no significant difference compared to the relationship between field AGNs and their hosts. The projected radial distributions of both AGN populations are consistent with the distribution of other cluster members. We argue that the apparent dichotomy between X-ray and IR AGNs can be understood as a combination of differing extinction due to cold gas in the host galaxies of the two classes of AGNs and the presence of weak star-formation in X-ray AGN hosts. ", "introduction": "Galaxy formation and evolution has long been a subject of considerable interest, with early work dedicated to exploring the physical processes responsible for star-formation \\citep{whip46}, explaining the genesis of the Milky Way \\citep{egge62}, and examining the evolution of galaxies in clusters \\citep{spit51}. Models for the evolution of galaxies in clusters gained strong observational constraints with the discovery of an apparent evolutionary sequence among local clusters \\citep{oeml74}. The discovery that the fraction of blue, spiral galaxies in relaxed galaxy clusters increases from $z=0$ to $z\\approx0.4$ quickly followed \\citep{butc78,butc84}. The dearth of spiral galaxies in the high-density regions at the centers of galaxy clusters is known as the morphology-density relation \\citep{dres80,post84,dres97,post05}. This relation places additional, strong constraints on evolutionary models for cluster galaxies. That star-forming galaxies are also rare in the centers of clusters had been previously suggested by the results of \\citet{oste60} and was subsequently observed in other work \\citep{gisl78,dres85}. The impact of environment on the frequency and intensity of star-formation at a wide variety of density scales has been measured using numerous visible \\citep{abra96,balo97,kauf04,pogg06,pogg08,vdLi10} and mid-infrared (MIR; \\citealt{sain08,bai09}) diagnostics. Star-forming galaxies are consistently found to be more common and to have higher star-formation rates (SFRs) in lower density environments and at higher redshift \\citep{kauf04,pogg06,pogg08}. The observed trends in star-formation with environment are usually attributed to variations in the sizes of gas reservoirs, either the existing cold gas or the hot gas that can cool to replenish the cold gas as it is consumed. Given that AGNs also consume cold gas to fuel their luminosity, similar patterns might be expected among AGNs. Indeed, recent work reveals strong dependencies of the luminosities and types of AGNs on environment (e.g.\\ \\citealt{kauf04,pope06,cons08,mont09}) for AGNs selected via visible-wavelength emission-line diagnostics. Von der Linden et al.\\ (2010) find fewer ``weak AGNs'' (primarily LINERS) among red sequence galaxies near the centers of clusters compared to the field, but they find no corresponding dependence among blue galaxies. Intriguingly, while \\citet{mont09} independently report a decline in the fraction of low-luminosity AGNs toward the centers of low-redshift clusters, they find an {\\it increase} in the fraction of LINERs in higher density environments. The difference is likely a result of evolution. \\citet{mont09} found qualitatively different behavior between their main $z\\sim1$ sample and the result produced when they applied their analysis to SDSS clusters. These results indicate that the variation of galaxy properties with local environment may influence the types of AGNs observed and that evolution in the relationship between some AGN classes and their host galaxies is important. Understanding the environmental mechanism that transforms star-forming galaxies into passive galaxies in clusters may help relate gas reservoirs in cluster galaxies to galaxy evolution as well as to AGN feeding and feedback. Several mechanisms to cause the transformation from star-forming to passive galaxies have been proposed. These include ram-pressure stripping of cold gas \\citep{gunn72,quil00,roed05}, strangulation \\citep{lars80,balo00,kawa08,mcca08} and galaxy harassment \\citep{moor96,moor98,lake98}. Each mechanism operates on a different characteristic timescale and has its greatest impact on galaxies of different masses and at different radii. In principle, the transition of galaxy populations from star-forming to passive as a function of environment can probe the relative importance of these processes. However, such approaches suffer from practical difficulties. For example, \\citet{bai09} argue that the similarity of the $24\\mu m$ luminosity functions observed in galaxy clusters and in the field suggests that the transition from star-formation to quiescence must be rapid, which implies that ram pressure stripping must be the dominant mechanism. Von der Linden et al.\\ (2010), by contrast, find a significant trend of increasing star-formation with radius up to $5 R_{200}$ from cluster centers. They conclude that preprocessing at the group scale is important, which is inconsistent with ram pressure stripping as the driver of the SFR-density relation. \\citet{pate09} find a similar trend for increasing average SFR with decreasing local density down to group-scale densities ($\\Sigma_{gal}\\approx1.0\\ {\\rm Mpc^{-2}}$) near RX J0152.7-1357 ($z=0.83$). The importance of preprocessing in group-scale environments reported by these authors suggests that strangulation rather than ram pressure stripping drives the SFR-density relation. The starkly different conclusions reached by \\citet{bai09} compared to \\citet{pate09} and \\citet{vdLi10}, despite their common use of star-forming galaxies to examine the influence of environment, highlight the difficulties inherent in such studies. Attempts to distinguish between various environmental processes become still more difficult with cluster samples that span a wide range in redshifts. The epoch of cluster assembly ($0\\leq z\\lesssim 1.5$, e.g.\\ \\citealt{berr09}) coincides with the epoch of rapidly declining star formation (e.g.\\ \\citealt{mada98,hopk06}) and AGN activity (e.g.\\ \\citealt{shav96,boyl98,shan09}), which makes it difficult to disentangle rapid environmental effects from the global reduction in the amount of available cold gas. \\citet{dres83} found early evidence for an increase in AGN activity with redshift, and the Butcher-Oemler effect had already provided evidence for a corresponding increase in SFRs. In the last decade, the proliferation of observations of high-redshift galaxy clusters at X-ray, visible and infrared wavelengths has yielded similar trends in the fraction of both AGNs \\citep{east07,mart09} and star-forming galaxies \\citep{pogg06,pogg08,sain08,hain09} identified using a variety of methods. These newer results have also examined cluster members confirmed from spectroscopic redshifts rather than relying solely on statistical excesses in cluster fields, which permits more detailed study of the relationships between galaxies and their parent clusters. The wide variety of AGN selection techniques employed in more recent studies represents an important step forward in understanding the dependence of AGNs on environment. Several recent papers have used X-rays to study the frequency and distribution of AGNs in galaxy clusters (\\citealt{mart06}, henceforth M06; \\citealt{mart07,siva08,arno09,hart09}) and their evolution with redshift \\citep{east07,mart09}. \\citet{mart09} found that the AGN fraction among cluster members increases with decreasing local density and increases dramatically ($f_{AGN}\\propto(1+z)^{5.3\\pm1.7}$) with redshift. They also found that X-ray identification produces a much larger AGN sample than visible-wavelength emission line diagnostics: only 4 of the 35 X-ray sources identified as AGNs by M06 would be classified as AGNs from their visible-wavelength emission lines. Similar results have been found when comparing radio, X-ray and mid-IR AGN selection techniques for field AGNs (e.g.\\ \\citealt{hick09}). The different AGN selection techniques identify different AGN populations and suffer from distinctive selection biases. Both X-ray and visible-wavelength techniques can miss AGNs due to absorption, either in the host galaxy or in the AGN itself; however, X-ray selection can find lower luminosity AGNs and AGNs behind larger absorbing columns compared to emission line selection. Mid-infrared selection techniques suffer from relatively poor angular resolution, so they are mainly sensitive to AGNs that outshine their host galaxies in the band(s) used to perform the AGN selection. The X-ray and visible techniques can also be contaminated by emission from the host galaxy. While the identification of X-ray sources with $L_{\\rm X}>10^{42}\\ {\\rm erg\\ s^{-1}}$ as AGNs is unambiguous, X-ray luminosities in the $10^{40}$--$10^{42}\\ {\\rm erg\\ s^{-1}}$ range can be produced by low-mass X-ray binaries (LMXBs), high-mass X-ray binaries (HMXBs), and thermal emission from hot gas. Both visible-wavelength and MIR indicators are subject to contamination from young stars, which produce emission lines and heat dust near star-forming regions until it emits in the MIR. Even the interpretation of the well-established Baldwin-Phillips-Terlevich diagram \\citep{bald81} can be controversial in the transition region between star-forming galaxies and AGNs. These difficulties motivate the use of multiple techniques to obtain a complete census of AGN and to correctly identify potential imposters. In this paper, we extend the work of Martini et al.\\ (2006, 2007) by supplementing their X-ray imaging and visible-wavelength photometry with MIR observations from the Spitzer Space Telescope. We use these data to select AGNs independent of their X-ray emission. We also measure the properties of AGN host galaxies by fitting their visible to MIR spectral energy distributions (SEDs). We discuss our visible and MIR data reduction and photometry in Section \\ref{secObs}. Section \\ref{secMethods} details our techniques for identifying AGNs and measuring galaxy properties, and we describe the results in Section \\ref{secResults}. We discuss the implications for the relationship between AGNs and their host galaxies in Section \\ref{secDiscuss}. Throughout this paper we use the WMAP 5-year cosmology---a $\\Lambda$CDM universe with $\\Omega_{m}=0.26$, $\\Omega_{\\Lambda}=0.74$ and $h=0.72$ \\citep{dunk09}. ", "conclusions": "We have used {\\it Spitzer} imaging of galaxy clusters to identify AGNs and to measure the masses and star-formation rates of their host galaxies. We find that AGNs identified by this technique have very little overlap with AGNs identified in X-rays. We compared the host galaxies of AGNs identified using the two methods and determined that, while their masses and SFRs are indistinguishable, IR AGNs reside in galaxies with higher sSFRs than both X-ray AGN hosts and the parent sample of cluster galaxies. The hosts of X-ray AGNs have sSFRs that are somewhat lower than but consistent with the sSFRs seen in cluster galaxies as a whole. The difference between X-ray AGN hosts and normal cluster galaxies is significant only when comparing their positions in visible color-magnitude and MIR color-color diagrams. X-ray AGN hosts are rarely found in the regions of both diagrams associated with vigorous star-formation. We also find that accretion rates of both X-ray and IR AGNs correlate strongly with SFR in their host galaxies. This suggests that X-ray and IR AGNs are physically similar and are fueled by the same mechanism. We hypothesize that the larger sSFRs seen in IR AGN hosts indicate larger cold gas fractions in these galaxies, and suggest that this could account for the apparent dichotomy between X-ray and IR AGNs. A moderately large cold gas column density of $10^{23}\\ {\\rm cm^{-2}}$ could suppress the X-ray emission from the IR AGNs enough that we would be unable to detect them. The presence of IR AGNs but not X-ray AGNs in galaxies with very red optical colors, indicative of strong absorption, lends credence to this hypothesis. It might also be verifiable directly by deep X-ray observations of either AC 114 or Abell 1689 to search for X-ray emission from IR AGNs and to determine if such X-ray emission shows evidence for absorption intrinsic to the host galaxy. For example, the most luminous IR AGN with no X-ray counterpart in Abell 1689 could be detected by Chandra with ${\\rm S/N}=3$ per resolution element at $4\\ keV$---the energy cutoff for objects with $N_{\\rm H}=10^{23}\\ {\\rm cm^{-2}}$---in $160\\ ks$. This would allow a crude model spectrum to be constructed and the intrinsic absorption column to be measured. Finally, we have obtained NIR spectra of several IR AGN in Abell 1689, which we will examine for high-ionization emission lines that would unambiguously indicate the presence of an AGN. Following \\citet{mart07}, we compared the radial distributions of AGNs and all cluster members. We eliminated one AGN with a spectroscopic redshift from the literature that incorrectly identified a background quasar as a cluster member. Without this object, the significance of their result that luminous X-ray AGNs ($L_{\\rm X}>10^{42}\\ {\\rm erg\\ s^{-1}}$) are more concentrated than cluster members as a whole is reduced to $\\sim90\\%$\\ confidence. While this result is no longer significant, it would be worthwhile to extend the present sample using archival {\\it Chandra} imaging of additional clusters to either confirm or refute that X-ray luminous AGNs are more concentrated than the galaxy populations of their parent clusters. It is unlikely, however, that a similar exercise using IR AGNs would yield a positive result, as the radial distribution of IR AGNs agrees very closely with the distribution of cluster galaxies." }, "1101/1101.1739_arXiv.txt": { "abstract": "{The Galactic microquasar GRS 1915+105 exhibits at least sixteen types of variability classes. Transitions from one class to another could take place in a matter of hours. In some of the classes, the spectral state transitions (burst-off to burst-on and vice versa) were found to take place in a matter of few to few tens of seconds.} {In the literature, there is no attempt to understand in which order these classes were exhibited. Since the observation was not continuous, the appearances of these classes seem to be in random order. Our goal is to find a natural sequence of these classes and compare with the existing observations. We also wish to present a physical interpretation of the sequence so obtained using two component advective flow model of black hole accretion.} {In the present paper, we compute the ratios of the power-law photons and the black body photons in the spectrum of each class and call these ratios as the `Comptonizing efficiency' (CE). We sequence the classes from the low to the high value of CE. The number of photons were obtained by fitting the spectra of two independent sets of data of each class with disk blackbody and power-law components, after making suitable correction for the absorption in the intervening medium.} {We clearly find that each variability class could be characterized by a unique average Comptonizing efficiency. The sequence of the classes based on this parameter seem to be corroborated by a handful of the observed transitions caught by Rossi X-ray timing explorer and the Indian payload Indian X-ray Astronomy Experiment and we believe that future observation of the object would show that the transitions can only take place between consecutive classes in this sequence. Since the power-law photons are produced by inverse Comptonization of the intercepted soft-photons from the Keplerian disk, a change in CE actually corresponds to a change in geometry of the Compton cloud. Thus we claim that the size of the Compton cloud gradually rises from very soft class to the very hard class.} {} ", "introduction": "The enigmatic stellar mass black hole binary GRS 1915+105 (Harlaftis \\& Greiner, 2004) was first discovered in 1992 by the WATCH detectors (Castro-Tirado et al. 1992) as a transient source with a significant variability in X-ray photon counts (Castro-Tirado et al. 1994). In the RXTE era, GRS 1915+105 was monitored thousands of times in the X-ray band and the scientific results reveal a unique nature of this compact object. The radio observation with VLA suggests apparent superluminal nature of its radio jets. Radio observation constrains that its maximum distance is no more than $13.5$ kpc and that the jet axis makes an angle of $70^{\\circ}$ with the line of sight (Mirabel \\& Rodriguez, 1994). Continuous X-ray observation of GRS 1915+105 reveals that the X-ray intensity of the source changes peculiarly in a variety of timescales ranging from seconds to days (Greiner et al., 1996, Morgan et al., 1997). Quasi-Periodic Oscillations (QPOs) are observed in a wide range of frequencies. QPOs in this source are associated with different types of X-ray variabilities and their timing properties are correlated with spectral features (Muno et al., 1999, Sobczak et al., 1999, Rodriguez et al., 2002, Vignarca et al., 2003,). The origin of QPO frequencies between $0.5$ to $10$ Hz is identified to be due to the oscillation of the Comptonized photons, presumably emitted from the post-shock region of the low angular momentum (sub-Keplerian) flow (Chakrabarti \\& Manickam, 2000, hereafter CM00; Rao et al., 2000). Small scale variabilities of GRS 1915+105 are identified with local variation of the inner disk (Nandi et al., 2000, Chakrabarti \\& Manickam, 2000; Migliari \\& Belloni, 2003). Several observers have reported that this object exhibited many types of variability classes (Yadav et al., 1999; Rao, Yadav \\& Paul, 2000; Belloni et al. 2000, Chakrabarti \\& Nandi, 2000; Naik et al. 2002a). Depending on the variation of photon counts in different arbitrary energy bands (hardness ratio) and color-color diagram of GRS 1915+105, the X-ray variability of the source was found to have fifteen arbitrarily named ($\\alpha,\\ \\beta,\\ \\gamma,\\ \\delta,\\ \\phi,\\ \\chi_1, \\chi_2, \\chi_3, \\chi_4,\\ \\mu,\\ \\nu,\\ \\lambda,\\ \\kappa,\\ \\rho,\\ \\theta$) classes. In a 1999 observation of RXTE, the existence of another class $\\omega$ was reported (Klein-Wolt et al., 2002; Naik et al. 2002a). In the so-called $\\chi$ (i.e., $\\chi_1$ to $\\chi_4$) class, the strong variability as is found in other classes is absent. The classes named $\\chi_1, \\chi_3$, $\\beta$ and $\\theta$ are associated with the presence of strong radio jets (Naik \\& Rao, 2000, Vadawale et al., 2003). To understand the above features from a dynamical point of view, we carried out a correlation study in between temporal and spectral features of this source in different classes. A preliminary report is presented in Pal, Nandi \\& Chakrabarti (2008, hereafter PNC08). While a large number of papers have been published in the literature on GRS 1915+105, to our knowledge, there is no work which actually asked the question: are these classes arbitrary, or they appear in a given sequence? The problem lies in the fact that no satellite continuously observed GRS 1915+105. Sporadic observations caught the object in sporadic classes. In the present paper, we try to show that these classes could be parameterized by a common parameter, namely, the ratio between the number of photons in the power-law component and the black body component. We call this as the Comptonizing Efficiency or CE. Since the number of photons in the power-law component depends on the degree of interception by the so-called hot electron cloud or Compton cloud (Sunyaev \\& Titarchuk, 1980, 1985), different classes are therefore parameterized by the average size of the Compton cloud. Along with the dynamical evolution of CE, we compute the spectrum and the power density spectrum (PDS) for each of the classes. To accomplish the computation of CE, we separate out the photons $\\gamma_{BB}$ of the black body component and the photons $\\gamma_{PL}$ from the power-law component and take the running ratio CE as a function of time to study how the Compton cloud itself varies in a short time scale. Our findings reveal that the Compton cloud is highly dynamic. We present possible scenarios of what might be occurring to it in different variability classes. The paper is organized as follows: in the next Section, we present a general discussion of the observation, our criteria of selection of data for analysis and analysis technique. In \\S 3, we discuss the procedure of calculation that is adopted to calculate the photon numbers. In \\S 4, the results are presented. In \\S 5, we present a unifying view where we show that the Comptonizing efficiency may be a key factor to distinguish among various classes. Finally, in \\S 6, we make concluding remarks. ", "conclusions": "In this paper, we have analyzed all known types of light curves of the enigmatic black hole GRS1915+105 and computed the dynamical nature of the energy and the power density spectra. We did not characterize these classes by conventional means, such as using hardness ratios defined in certain energy range since such a characterization does not improve our view about the physical picture. Furthermore, characterization using certain energy range is possible only in a case by case basis, and is not valid for the black holes of all masses. Instead, we asked ourselves whether we can distinguish one class from another from physical point of view purely in a model independent way. We observed that independent of what the nature of the Compton cloud is, the weighted mean $$ of Comptonizing efficiency (CE) obtained every 16 seconds of the binned data, increases monotonically as the class varies. This pattern we find was verified with two sets of data covering all the variability classes. So, is not arbitrary -- it is characteristics of a class. Based on the values of , it is observed that the classes belong to three states: Classes I-IV in softer states, Classes V-XII in intermediate states, and Classes XIII-XVI belong to the harder state. When the weighted average value of CE is monotonically arranged, we obtain a sequence which appears to be followed during the actual transitions. Indeed, when comparing with available data of PCU of RXTE and the Indian payload IXAE onboard IRS-P3, the transitions from one class to another as reported in the literature do follow our sequence. Given that a large variation of CE occurs in a given variability class, it is puzzling why the sequence of the classes obtained by us should follow changing the CE values {\\it averaged} over the whole class. It is possible that a specific value of $$ actually forces the system to be in a given class, just as a parameter such as wind speed decides the mean angle of an oscillating pendulum. The excursion of CE in that class could be due to totally different physical process and not necessarily due to interception of soft photons by the CENBOL alone. May be the mass-loss rate of CENBOL is playing a role, which in turn depends on the shock strength. In the two component model of the Chakrabarti-Titarchuk (CT95), the CENBOL and associated outflow play the role of the Compton cloud. In this model, the increase in Compton efficiency can be affected in several ways: (i) by increasing the shock location increases the size of the CENBOL which intercepts larger number of soft photons and/or (ii) by increasing the accretion rate of the sub-Keplerian component, which increases the optical depth and scatter more soft photons to produce power-law photons. In CT95 and C97 it was shown that harder states are produced by both the effects mentioned above. Indeed, we find that the classes with harder states have more CE. Shifting of the shock locations is possible by changes in viscosity or changes in cooling rate in the post-shock region. The time scales of such effects in a sub-Keplerian flow could take hours. Time scale of changing the global sub-Keplerian flow rates could be comparable to the free-fall time from the outer edge, i.e., of the order of a day or so. On the other hand, in Classes VII and VIII etc., the CE changes in a matter of minutes. Such a short time variability of Comptonizing efficiency is possible if the base of the outflow is abruptly cooled and returned back to the disk increasing the accretion rate of the Keplerian/sub-Keplerian rates locally (CM00). In this paper, we find that the $$ is really important in deciding the sequence. However, for a given class, the degree of excursion of CE in a given variability type must depend on another parameter, such as the outflow rate which in turn depends on the shock strength. A simple estimate (C99) suggests that the shock strength decides the outflow rate and hence the time taken by the base of the outflow to reach unit optical depth $\\tau$ for Compton scattering (CM00). For a very strong shock, this outflow rate is very weak, as is evidenced by weak (few tens of miliJansky) radio flux even for `radio-loud' classes. For intermediate shock strength the outflow rate is higher, and it is easier to have $\\tau=1$ in a short time scale (CM00) and fractional change in CE also becomes high. In classes XII-XVI, we not only see CE to be very high, the fractional change in CE is very small as well. The aspect of classification in terms of outflow rate is being looked into. The analysis is in progress and will be reported elsewhere." }, "1101/1101.3350_arXiv.txt": { "abstract": "The Press-Ryden-Spergel (PRS) algorithm is a modification to the field theory equations of motion, parametrized by two parameters ($\\alpha$ and $\\beta$), implemented in numerical simulations of cosmological domain wall networks, in order to ensure a fixed comoving resolution. In this paper we explicitly demonstrate that the PRS algorithm provides the correct domain wall dynamics in $N+1$-dimensional Friedmann-Robertson-Walker (FRW) universes if $\\alpha+\\beta/2=N$, fully validating its use in numerical studies of cosmic domain evolution. We further show that this result is valid for generic thin featureless domain walls, independently of the Lagrangian of the model. ", "introduction": "The dynamics of cosmological domain walls has been investigated using both high-resolution numerical simulations and a semi-analitical velocity-dependent one-scale (VOS) model \\cite{Press,PinaAvelino:2006ia,Avelino:2006xy,Avelino:2006xf,Battye:2006pf,Avelino:2008ve,Avelino:2009tk}. Most of these studies were motivated by the suggestion \\cite{Bucher:1998mh} that a frozen domain wall network could be responsible for the observed acceleration of the Universe (see also \\cite{Carter:2004dk,Battye:2005hw,Battye:2005ik,Carter:2006cf}). Although, current observational constraints on the equation of state parameter of dark energy strongly disfavor domain walls as a single dark energy component \\cite{Komatsu:2010fb,Frieman:2008sn}, they are unable to rule out a substantial impact of a frustrated domain wall network on the acceleration of the Universe around the present time. However, analytical and numerical results strongly support the conjecture that no frustrated domain wall network, accounting for a significant fraction of the energy density of the Universe today, could have emerged from realistic phase transitions. These results, on their own, seem to rule out any significant contribution of domain walls to the dark energy budget. However, they rely heavily on the validity of the so-called Press-Ryden-Spergel (PRS) algorithm used in cosmological domain wall network simulations. Domain walls have a constant physical thickness and, consequently, their comoving thickness decreases proportionally to the inverse of the cosmological scale factor. In numerical studies of cosmological domain wall evolution the rapid decrease of the comoving domain wall thickness would be serious problem since it would imply that domain walls could only be resolved during a small fraction of the simulation dynamical range. The PRS algorithm is a modification to the field theory equations of motion, implemented in numerical simulations of cosmological domain wall evolution, allowing for a fixed comoving resolution. It has been argued that the PRS algorithm \\cite{Press}, provides the correct domain wall dynamics in $3+1$ dimensions, as long as $\\alpha+\\beta/2=3$ ($\\alpha$ and $\\beta$ are the PRS algorithm parameters of ref. \\cite{Press}). Although this claim is strongly supported by numerical tests it has never been proven that the same Nambu-Goto effective action is recovered in the thin wall limit. In this paper we eliminate this shortcoming, extending the analysis to generic thin featureless domain walls in FRW universes with an arbitrary number of spatial dimensions. ", "conclusions": "In this paper we explicitly demonstrated that the PRS algorithm provides the correct dynamics of thin featureless domain walls in FRW universes with an arbitrary number, $N$, of spatial dimensions, if $\\alpha+\\beta/2=N$. Our results fully justify the use of the PRS algorithm in numerical studies of cosmological domain wall network evolution. Although, fixing the comoving thickness of the domain walls, using the PRS algorithm, increases artificially the impact of the junctions on the overall network dynamics during the course of the simulations, this effect is negligible for the light junctions usually considered in such simulations." }, "1101/1101.4570_arXiv.txt": { "abstract": "{}{Based on a rather complete model atom for neutral and singly-ionized iron, we evaluate non-local thermodynamical equilibrium (non-LTE) line formation for the two ions of iron and check the ionization equilibrium between \\ion{Fe}{i} and \\ion{Fe}{ii} in model atmospheres of the cool reference stars.} {A comprehensive model atom for Fe with more than 3\\,000 measured and predicted energy levels is presented. As a test and first application of the improved model atom, iron abundances are determined for the Sun and five stars with well determined stellar parameters and high-quality observed spectra. The efficiency of inelastic collisions with hydrogen atoms in the statistical equilibrium of iron is estimated empirically from inspection of their different influence on the \\ion{Fe}{i} and \\ion{Fe}{ii} lines in the selected stars.} {Non-LTE leads to systematically depleted total absorption in the \\ion{Fe}{i} lines and to positive abundance corrections in agreement with the previous studies, however, the magnitude of such corrections is smaller compared to the earlier results. Non-LTE corrections do not exceed 0.1~dex for the solar metallicity and mildly metal-deficient stars, and they vary within 0.21~dex and 0.35~dex in the very metal-poor stars HD~84937 and HD~122563, respectively, depending on the assumed efficiency of collisions with hydrogen atoms. Based on the analysis of the \\ion{Fe}{i}/\\ion{Fe}{ii} ionization equilibrium in these two stars, we recommend to apply the Drawin formalism in non-LTE studies of Fe with a scaling factor of 0.1. For the \\ion{Fe}{ii} lines, non-LTE corrections do not exceed 0.01~dex in absolute value. This study reveals two problems. The first one is that $gf-$values available for the \\ion{Fe}{i} and \\ion{Fe}{ii} lines are not accurate enough to pursue high-accuracy absolute stellar abundance determinations. For the Sun, the mean non-LTE abundance obtained from 54 \\ion{Fe}{i} lines is 7.56$\\pm$0.09 and the mean abundance from 18 \\ion{Fe}{ii} lines varies between 7.41$\\pm$0.11 and 7.56$\\pm$0.05 depending on the source of the $gf-$values. The second problem is that lines of \\ion{Fe}{i} give, on average, a 0.1\\,dex lower abundance compared to those of \\ion{Fe}{ii} lines for HD\\,61421 and HD\\,102870, even when applying a differential line-by-line analysis relative to the Sun. A disparity between neutral atoms and first ions points to problems of stellar atmosphere modelling or/and effective temperature determination.} {} ", "introduction": "Iron plays an outstanding role in studies of cool stars thanks to the many lines in the visible spectrum, which are easy to detect even in very metal-poor stars. Iron serves as a reference element for all astronomical research related to stellar nucleosynthesis and the chemical evolution of the Galaxy. Iron lines are used to derive basic stellar parameters, i.e. the effective temperature, $\\Teff$, from the excitation equilibrium of \\ion{Fe}{i} and the surface gravity, $\\logg$, from the ionization equilibrium between \\ion{Fe}{i} and \\ion{Fe}{ii}. In stellar atmospheres with $\\Teff > 4500$~K, neutral iron is a minority species, and its statistical equilibrium (SE) can easily deviate from thermodynamic equilibrium due to deviations of the mean intensity of ionizing radiation from the Planck function. Therefore, since the beginning of the 1970s, a large number of studies attacked the problem of non-local thermodynamic equilibrium (non-LTE) line formation for iron in the atmospheres of the Sun and cool stars. The original model atoms were from \\citet{Tanaka71,Athay1972,Boyarchuketal85,Gigas86,Takeda91,Grattonetal99,Thevenin1999,Gehren2001a,ShTB01}, and \\citet{Colletal05}, and they were widely applied in stellar parameter and abundance analyses \\citep[see][for references]{aspl05}. It was understood that the main non-LTE mechanism for \\ion{Fe}{i} is ultra-violet (UV) overionization of the levels with excitation energy of 1.4 to 4.5~eV. This results in an underpopulation of neutral iron where all \\ion{Fe}{i} lines are weaker than their LTE strengths, and it leads to positive non-LTE abundance corrections. The need for a new analysis was motivated by the following problems uncovered by the previous non-LTE calculations for iron. First, the results obtained for the populations of high-excitation levels of \\ion{Fe}{i} were not always convincing. The highest levels presented in the model atom did not couple thermally to the ground state of \\ion{Fe}{ii} indicating substantial term incompleteness. To force the levels near the continuum into LTE, an upper level thermalization procedure was applied \\citep[more details in][]{Gehren2001a,Korn03,Colletal05}. A second aspect is the treatment of poorly known inelastic collisions with hydrogen atoms. Their role in establishing the statistical equilibrium of atoms in cool stars is debated for decades, from \\citet{gehren75} to \\citet{barklem2010}. Experimental data on \\ion{H}{i} collision cross-sections are only available for the resonance transition in \\ion{Na}{i} \\citep{fleck91,belyaev99}, and detailed quantum mechanical calculations were published for the transitions between first nine levels in \\ion{Li}{i} \\citep{belyaev03,barklem2003} and \\ion{Na}{i} \\citep{belyaev99,belyaev2010,barklem2010}. For all other chemical species, the basic formula used to calculate collisions with \\ion{H}{i} atoms is the one proposed by \\citet{D68,D69}, as described by \\citet{Steenbock1984}, and it suggests that their influence is comparable to electron impacts. The laboratory measurements and quantum mechanical calculations indicate that the Drawin formula overestimates rate coefficients for optically allowed transitions by one to seven orders of magnitude. Therefore, various approaches were employed in the literature to constrain empirically the efficiency of \\ion{H}{i} collisions. The studies of stellar \\ion{Na}{i} lines favor a low efficiency of this type of collisions. For example, \\citet{h_cool_na_o} found that the center-to-limb variation of the solar \\ion{Na}{i} 6160\\AA\\ line is reproduced in the non-LTE calculations with pure electron collisions. \\citet{Grattonetal99} calibrated \\ion{H}{i} collisions with sodium using RR~Lyr variables and concluded that the Drawin rates should be decreased by two orders of magnitude. Based on their solar \\ion{Na}{i} line profile analysis, \\citet{mg_c6} and \\citet{Takeda95} recommended to scale the Drawin rates by a factor \\kH\\ = 0.05 and 0.1, respectively. With a similar value of \\kH\\ = 0.1, the ionization equilibrium between \\ion{Ca}{i} and \\ion{Ca}{ii} in selected metal-poor stars was matched consistently with surface gravities derived from {\\sc Hipparcos} parallaxes \\citep{mash_ca}. On the other hand, spectroscopic studies of different chemical species suggested that the \\ion{H}{i} collision rates might be reasonably well described by Drawin's formula with \\kH\\ $\\ge 1$. For example, empirical estimates by \\citet{Grattonetal99} resulted in \\kH\\ = 3 for \\ion{O}{i} and \\ion{Mg}{i} and \\kH\\ = 30 for \\ion{Fe}{i}. \\citet{h_cool_na_o} and \\citet{Pereira2009} inferred \\kH\\ = 1 from the analysis of the center-to-limb variation of solar \\ion{O}{i} triplet $\\sim7770$\\AA\\ lines. The same value of \\kH\\ = 1 was obtained by \\citet{Takeda95} from solar \\ion{O}{i} line profile fits. For a review of studies constraining empirically the efficiency of \\ion{H}{i} collisions, see \\citet{Lambert1993,Holweger1996}, and \\citet{mash_lund}. As a result of applying incomplete model atoms and a different treatment of collisions with hydrogen atoms, no consensus on the expected magnitude of the non-LTE effects was achieved in the previous studies of iron, and results were in conflict with each other in some cases. For example, \\citet{Korn03} found a negligible discrepancy between the non-LTE spectroscopic and {\\sc Hipparcos} astrometric distances of the halo star HD\\,84937, while the non-LTE calculations of \\citet{Thevenin1999} resulted in a 34\\,\\%\\ smaller spectroscopic distance of that same star. This study aims to construct a fairly complete model atom of iron, to be tested using the Sun and selected cool stars with high-quality observed spectra and reliable stellar parameters. Compared with the previous non-LTE analyses of iron, the model atom of \\ion{Fe}{i} was extended to high-lying levels predicted by the atomic structure calculations of \\citet{Kurucz2009} and this turned out to be crucial for a correct treatment of the SE of iron in cool star atmospheres. With our improved model atom we tried to constrain the scaling factor \\kH\\ empirically. We realize that the real temperature dependence of hydrogen collision rates could be very different from that of the classical Drawin formalism, and we may not always achieve consistent \\kH\\ values from the analysis of different stars. We also realize that the required thermalizing process not involving electrons in the atmospheres of cool metal-poor stars could be very different from inelastic collisions with neutral hydrogen atoms. For example, \\citet{barklem2010} \\citep[see also][]{barklem2003,belyaev03,belyaev2010} uncovered the importance of the ion-pair production and mutual neutralisation process $A(nl) + {\\rm H}(1s) \\rightleftharpoons A^+ + {\\rm H^-}$ for the SE of Li and Na. Since no accurate calculations of either inelastic collisions of iron with neutral hydrogen atoms or other type processes are available, we simulate an additional source of thermalization in the atmospheres of cool stars by parametrized \\ion{H}{i} collisions. Investigating the Sun as a reference star for further stellar differential line-by-line analysis, we also derive the solar iron absolute abundance and check the solar \\ion{Fe}{i}/\\ion{Fe}{ii} ionization equilibrium using an extended list of lines, which can be detected at solar metallicity down to [Fe/H] = $-2.5$. We find it important to inspect the accuracy of atomic data for various subsamples of iron lines in view of comprehensive abundance studies across the Galaxy targeting at stars of very different metallicities. This paper is organized as follows. The model atom of iron and the adopted atomic data are presented in Sect.\\,\\ref{sect:NLTE}. There we also discuss how including the bulk of predicted \\ion{Fe}{i} levels in the model atom affects the SE of iron. In Sect.\\,\\ref{sect:sun}, the solar iron spectrum is studied to provide the basis for further differential analyses of stellar spectra. Section\\,\\ref{sect:stars} describes observations and stellar parameters of our sample of stars, and Sect.\\,\\ref{sect:stellar_iron} investigates which line-formation assumptions lead to consistent element abundances from both ions, \\ion{Fe}{i} and \\ion{Fe}{ii}. Uncertainties in the iron non-LTE abundances are estimated in Sect.\\,\\ref{sect:uncertainty}. Our recommendations and conclusions are given in Sect.\\,\\ref{conclusion}. ", "conclusions": "\\label{conclusion} In this study, a comprehensive model atom for neutral and singly-ionized iron was built up using atomic data for the energy levels and transition probabilities from laboratory measurements and theoretical predictions. With a fairly complete model atom for \\ion{Fe}{i}, the calculated statistical equilibrium of iron changed substantially by achieving close collisional coupling of the \\ion{Fe}{i} levels near the continuum to the ground state of \\ion{Fe}{ii}. There is no need anymore for the enforced upper level thermalization procedure that was applied in the previous non-LTE analyses \\citep{Gehren2001a,Korn03,Colletal05}. Non-LTE line formation for \\ion{Fe}{i} and \\ion{Fe}{ii} lines was considered in 1D model atmospheres of the Sun and five reference stars with reliable stellar parameters, which cover a broad range of effective temperatures between 4600~K and 6500~K, gravities between $\\logg$ = 1.60 and 4.53, and metallicities between [Fe/H] = $-2.5$ and $+0.1$. We found that the departures from LTE are negligible for the \\ion{Fe}{ii} lines over the whole stellar parameter range considered. For \\ion{Fe}{i}, the non-LTE effects on the abundances are expected to be small for stars with solar-type metallicities such as the Sun, Procyon, and $\\beta$~Vir, and for mildly metal-deficient stars such as $\\tau$~Cet: a non-LTE correction is at the level of $+0.1$~dex in non-LTE with pure electronic collisions and of a few hundredths, when inelastic collisions with hydrogen atoms are taken into account in the SE calculations with \\kH\\ $\\ge 0.1$. From a differential line-by-line analysis of stellar spectra we found that the iron ionization equilibrium is not fulfilled in Procyon and $\\beta$~Vir at their fundamental parameters, 6510/3.96 and 6060/4.11, respectively, independent of either LTE or non-LTE. For Procyon, an upward revision of its temperature by 80~K removes the obtained abundance imbalance. In contrast, consistent iron abundances from both ionization stages were obtained in $\\tau$~Cet at its given stellar parameters in both non-LTE with \\kH\\ $\\ge 0.1$ and in LTE. Significant departures from LTE for \\ion{Fe}{i} were found in the two VMP stars of our sample, HD\\,84937 and HD\\,122563. Our results indicate the need for a thermalizing process not involving electrons in their atmospheres. Since there are no accurate theoretical considerations of appropriate processes for iron, we simulate an additional source of thermalization in the atmospheres of cool stars by parametrized \\ion{H}{i} collisions. Close inspection of the ionization equilibrium between \\ion{Fe}{i} and \\ion{Fe}{ii} in HD\\,84937 and HD\\,122563 leads us to choose {\\it a scaling factor of 0.1} to the formula of \\citet{Steenbock1984} for calculating hydrogenic collisions. The uncertainty in the estimated \\kH\\ value is expected to result in abundance error of no more than 0.08~dex in the non-LTE calculations for \\ion{Fe}{i} in the F5 - K0 type stars in the super-solar metallicity down to [Fe/H] = $-2.5$ domain. However, the situation can be significantly worse for extremely and ultra-metal-poor stars. Exactly for such objects, the determination of the surface gravity relies in most cases on the analysis of the \\ion{Fe}{i}/\\ion{Fe}{ii} ionization equilibrium. Theoretical studies are urgently needed to evaluate the cross-sections of inelastic collisions of \\ion{Fe}{i} with \\ion{H}{i} atoms and to search for and evaluate other types of thermalizing processes. For the Sun, the use of \\kH\\ = 0.1 leads to an average \\ion{Fe}{i} non-LTE correction of 0.03~dex and a mean \\ion{Fe}{i} based abundance of 7.56$\\pm$0.09. A mean solar abundance derived from \\ion{Fe}{ii} lines varies between 7.41$\\pm$0.11 and 7.56$\\pm$0.05 depending on the source of $gf-$values. A statistical error of 0.09 -- 0.11~dex is uncomfortably high for the Sun. It is, most probably, due to the uncertainty in $gf-$values and, in part, van der Waals damping constants. The problem of oscillator strengths of the Fe lines and their influence on the derived solar iron abundance has been debated for decades \\citep{Blackwell95,Holweger1995,Kostik1996,Caffau2010}. Stellar astrophysics needs accurate atomic data for an extended list of the iron lines, which could be measured in stars of very different metallicities. We therefore call on laboratory atomic spectroscopists for further efforts to improve $gf-$values of the \\ion{Fe}{i} and \\ion{Fe}{ii} lines used in abundance analysis. Using our carefully calibrated model of iron, cool stars over a broad range of metallicities encountered in the Galaxy can now be analyzed in a homogeneous way to derive their iron abundance and gravity without resorting to trigonometric parallaxes. The dependence of departures from LTE on stellar parameters and an application of the non-LTE technique to the known ultra-metal-poor stars ([Fe/H] $< -4.5$) will be presented in a forthcoming paper." }, "1101/1101.1952_arXiv.txt": { "abstract": "{ We determine the mean velocity dispersion of six Galactic outer halo globular clusters, AM 1, Eridanus, Pal 3, Pal 4, Pal 15, and Arp 2 in the weak acceleration regime to test classical vs. modified Newtonian dynamics (MOND). Owing to the nonlinearity of MOND's Poisson equation, beyond tidal effects, the internal dynamics of clusters is affected by the external field in which they are immersed. For the studied clusters, particle accelerations are much lower than the critical acceleration $a_0$ of MOND, but the motion of stars is neither dominated by internal accelerations ($a_i \\gg a_e$) nor external accelerations ($a_e \\gg a_i$). We use the N-body code N-MODY in our analysis, which is a particle-mesh-based code with a numerical MOND potential solver developed by Ciotti, Londrillo, and Nipoti (2006) to derive the line-of-sight velocity dispersion by adding the external field effect. We show that Newtonian dynamics predicts a low-velocity dispersion for each cluster, while in modified Newtonian dynamics the velocity dispersion is much higher. We calculate the minimum number of measured stars necessary to distinguish between Newtonian gravity and MOND with the Kolmogorov-Smirnov test. We also show that for most clusters it is necessary to measure the velocities of between 30 to 80 stars to distinguish between both cases. Therefore the observational measurement of the line-of-sight velocity dispersion of these clusters will provide a test for MOND. ", "introduction": "Observable matter in galaxies and in clusters of galaxies cannot produce sufficient gravity to explain their dynamics. Cold dark matter (CDM) scenarios or alternative theories of gravitation are therefore invoked to resolve the problem. Nowadays, the CDM hypothesis is the dominant paradigm. This hypothetical matter does not interact with electromagnetic radiation and only shows its presence through its gravitational interaction. Even though the dark matter hypothesis has successfully explained the internal dynamics of galaxy clusters, gravitational lensing, and the standard model of cosmology within the framework of general relativity (GR) (\\cite{spe03}), much experimental effort has failed to yield a detection of dark matter particles. Moreover the results of high-resolution simulations of structure formation do not reproduce some observations on galactic scales, such as the central structures of rotation curves, the prevalence of low bulge-to-disc ratios, and the numbers and spatial distribution of the subhalos (\\cite{kly99,moore99,metz08, kroupa10}). Even the ability of the dark matter theory to account for the Tully-Fisher and Freeman relations is controversial (\\cite{bos00,gov10}). These shortcomings have not led to the rejection of the theory only because on galactic scales baryons are at least non-negligible contributors to the mass density, consequently simulations that include the complex physics of star formation are essential for reliable predictions. Currently such simulations are still at an experimental stage and are usually substituted by ``semi-analytic'' arguments that have weak theoretical underpinnings. One of the alternative theories to CDM is the so-called modified Newtonian dynamics (MOND) theory, which was originally proposed by Milgrom (1983) to explain the flat rotation curves of spiral galaxies at large distances by a modification of Newton's second law of acceleration below a characteristic scale of $a_{0}\\simeq1.2\\times10^{-10} $ms$^{-2}=3.6$ pc/(Myr)$^2$ without invoking dark matter (\\cite{bek84}). In MOND, Newton's second law is modified to $\\mu(a/a_{0}){\\bf a} = {\\bf a}_{N} + \\nabla\\times{\\bf H}$, where $\\rho$ is the mass-density distribution, ${\\bf a}_N$ is the Newtonian acceleration vector, ${\\bf a}$ is the MONDian acceleration vector, $a=|\\textbf{a}|$ is the absolute value of MONDian acceleration, $\\mu$ is an interpolating function for the transition from the Newtonian to the MONDian regime, which runs smoothly from $\\mu(x)=x$ at $x\\ll1$ to $\\mu(x)=1$ at $x\\gg1$ (\\cite{bek84}). Different interpolating functions have been suggested, such as the simple function, $\\mu(x)=x/(1+x)$ (\\cite{fam05}) and the standard interpolation function, $\\mu(x)=x/\\sqrt{1+x^2}$ (\\cite{mil83}). Because the simple function fits galactic rotation curves better than the standard function (\\cite{gent10}), in this paper we use the simple function. The value of the curl field ${\\bf H}$ depends on the boundary conditions and the spatial mass distribution and vanishes only for some special symmetries (\\cite{bek84}). The non-linearity of the MOND field equation leads to difficulties for standard N-body codes and makes the use of the usual Newtonian N-body simulation codes impossible in the MOND regime. It has been shown that on galactic scales MOND can explain many phenomena better than CDM (\\cite{beg89, begm91, san02, san07, hag06, mal09, gen07, mil94, bra00, mil95, wu08, za06, tir07,has10}). MOND has been generalized to a general-relativistic version (\\cite{bek04,san05,zlo07,mil09}), making it possible to test its predictions for gravitational lensing. Dynamics of galaxies in clusters (\\cite{san02}) and the merging of galaxy clusters, where the baryonic mass is clearly separated from the gravitational mass (\\cite{Clo06}) cannot be completely explained by MOND without invoking some kind of hot dark matter, perhaps in the form of massive (active or sterile, 2 to 11 eV) neutrinos (\\cite{ang06,ang10}). In order to decide whether MOND is a comprehensive theory to explain the dynamics of the universe, it is desirable to study MOND for objects in which no dark matter is supposed to exist and where the characteristic acceleration of the stars is less than the MOND critical acceleration parameter $a_0$. Globular clusters (GCs) are a perfect candidate since they are the largest virialized structure that do not contain dark matter (\\cite{moore96}). In the distant halo of our Milky Way there exist several low-mass GCs where both internal and external accelerations of stars are significantly below the critical acceleration parameter $a_0$ of MOND. Because GCs are assumed to be dark-matter-free, if MOND is true, the motions of stars must deviate from the standard Newtonian dynamics. It has been proposed by Baumgardt (2005) that some of these distant Galactic GCs are perfect tools to test gravitational theories in the regime of very weak accelerations. For MOND, the internal velocity dispersion among the stars in these clusters would be significantly higher than in Newtonian dynamics. The mean velocity dispersion of stellar systems for the two extreme cases of internal ($a_i\\gg a_e$) or external ($a_e\\gg a_i$) field domination have been derived analytically by Milgrom (\\cite{mil86,mil94}), assuming that the systems are everywhere in the deep-MOND regime ($a_e, a_i \\ll a_0$). Many systems that can be used to test MOND are not completely either internally or externally dominated. Globular clusters or dwarf galaxies of the Milky Way for example have internal and external accelerations that are of the same order (\\cite{bau05}), consequently one has to determine the velocity dispersion numerically for intermediate cases. Sollima and Nipoti (2009) constructed self-consistent, spherical models for stellar systems in MOND, neglecting the external field effect and presented a dynamical model for six galactic globular clusters. The presence of the external field effect breaks the spherical symmetry and validity of their model. Recently Haghi et al. (2009, hereafter HBK09) investigated the dynamics of star clusters by numerically modeling them in MOND, assuming circular orbits. They performed N-body simulations and presented analytical formulae for the velocity dispersion of stellar systems in the intermediate MOND regime, which are useful for a comparison with observational data of several GCs and dSph galaxies (for details on the numerical calculations see HBK09). In a follow-up paper, Jordi et al. (2009) determined the velocity dispersion (using 17 stars) and mass-function slope of Pal 14 and showed that MOND can hardly explain the low-velocity dispersion of this system. However, Gentile et al. (2010) showed that with the currently available data, the Kolmogorov-Smirnov (KS) test is still unable to exclude MOND with a sufficiently high confidence level. Moreover, the low density of Pal 14 suggests that binary stars may be an important issue for interpreting its measured velocity dispersion (\\cite{Kuepper10}), and the true velocity dispersion of Pal 14 could be much lower than the value reported by Jordi et al. (2009), thereby possibly posing an even larger challenge for MOND, but also for Newtonian gravity and for any understanding of the dynamics of this object as being in equilibrium. In this paper we calculate the prediction of MOND and Newtonian dynamics on the velocity dispersion of six other distant clusters of the MW (Table 1). In order to see the pure MONDian effects, we concentrate on systems in which the tidal radius is much larger than the gravitational radius \\footnote{ The gravitational radius is a measure of the size of the system and is related to the mass and potential energy as given in Eq. 2-132 in Binney and Tremaine (1987). In many stellar systems the gravitational radius can be approximated by the three-dimensional half-mass radius $r_h$ as $r_g=1.25r_h$ if the assumption of virial equilibrium is valid. } and therefore tidal effects are unimportant. In other words, this paper provides the basis for further observational efforts. The measurements of a low- (Newtonian) velocity dispersion would mean that MOND in its present form is in severe trouble and that globular clusters do not possess dark matter. In contrast, a high-velocity dispersion would either favor MOND or could be a hint to the existence of dark matter in globular clusters (\\cite{bau09}). The paper is organized as follows: In Section 2 we give a brief review of the external field effect (EFE) in MOND. The simulation setup is explained in Section 3. The numerical results for six clusters are discussed in Section 4. We present our conclusions in Section 5. ", "conclusions": "We compared global line-of-sight velocity dispersions of six distant low-density globular clusters of the Milky Way in MONDian and Newtonian dynamics and showed that they have a significantly higher velocity dispersion in MOND than the prediction of Newtonian dynamics. Using the N-MODY code, we obtained a large set of dissipationless numerical solutions for globular clusters with the MONDian initial conditions as end-products of the N-body computations by adding the external field. In order to produce different internal acceleration regimes, we changed the mass of the system over a wide range for each cluster. In addition, our results show that the clusters Pal 4, Pal 3, AM 1, and Eridanus are the best cases to test MOND owing to their larger absolute difference between MONDian and Newtonian velocity dispersions. These results will allow us to test MOND more rigorously than was possible so far, and this will enable us to compare MOND with observational data in the future. Recently Sollima and Nipoti (2009) have performed a test for MOND by constructing self-consistent dynamical models for outer galactic clusters. These authors neglected the external field effect of the Milky Way, and hence obtained higher estimates for the velocity dispersion of clusters compared with our results. This difference is reasonable. Indeed, the external field in MOND is leading the system to move from the deep-MONDian regime toward the quasi-Newtonian regime, and thus to reduce the velocity dispersion. In other words, the larger the external field, the smaller the internal acceleration, which implies a lower velocity dispersion for the cluster (\\cite{hag09}). It should be noted that the clusters could be on eccentric orbits, which means that the MOND predictions would be different because of the variation of the external field (in direction and amplitude) along the orbit. Using a KS test, we calculated the minimum number of stars that are sufficient to exclude MOND (under the hypothesis that these globular clusters are on circular orbits) at the 95\\% confidence level. We found that between 30 to 80 stars are necessary for most clusters to distinguish between both cases. This number of stars can be observed with current $8m$ class telescopes. Additional observational efforts to determine the velocity dispersions of these clusters and constraining the mass of the clusters by star counts would be highly important and provide a strict test of MOND. On the other hand, if MOND is the correct theory, these observations could be used to constrain the external field and consequently to put constraints on the potential in which the systems are embedded. According to the anisotropy profile, the simulated systems are isotropic throughout." }, "1101/1101.5380_arXiv.txt": { "abstract": "{ Radio supernovae (RSNe) are weak and rare events. Their typical maximum radio luminosities are of the order of only $10^{27}$\\,erg\\,s$^{-1}$\\,Hz$^{-1}$. There are, however, very few cases of relatively bright (and/or close) RSNe, from which the expansion of the shock and the radio light curves at several frequencies have been monitored covering several years. Applying the standard model of radio emission from supernovae, it is possible to relate the defining parameters of the modelled expansion curve to those of the modelled light curves in a simple algebraic way, by assuming an evolution law for the magnetic field and for the energy density of the population of synchrotron-emitting electrons. However, cooling mechanisms of the electrons may affect considerably this connection between light curves and expansion curve, and lead to wrong conclusions on the details of the electron acceleration and/or on the CSM radial density profile. In this paper, we study how electron cooling modifies the flux-density decay rate of RSNe for a set of plausible/realistic values of the magnetic field and for different expansion regimes. We use these results to estimate the magnetic fields of different RSNe observed to date and compare them to those obtained by assuming energy equipartition between particles and magnetic fields. For some of the best monitored RSNe, for which deceleration measurements, optically thin spectral index, and power-law time decay have been observed (SN\\,1979C, SN\\,1986J, SN\\,1993J, and SN\\,2008iz), we find self-consistent solutions for the index of the power-law circumstellar density profile ($s=2$ for all cases), the index of the power-law relativistic electron population (rather steep values, $ p = 2.3 - 3.0$) and the initial magnetic field (ranging from $\\sim 20$ to $> 100$\\,G).} ", "introduction": "\\label{I} Radio supernovae (RSNe), which are the radio counterparts of core-collapse supernovae (SNe), are weak and rare events. Only about $10-20$\\% of the observed SNe are detected in radio (e.g., Weiler et al. \\cite{Weiler2002}). Moreover, their typical maximum radio luminosities are of the order of $10^{27}$\\,erg\\,s$^{-1}$\\,Hz$^{-1}$ (flux densities of the order of 1\\,mJy for extragalactic distances, close to the sensitivity limits of present detectors). There are, however, very few cases of relatively bright RSNe, from which the expansion curve of the shock, using Very Long Baseline Interferometry (VLBI) observations, and radio light curves at several frequencies were obtained covering, in some cases, several years, e.g.: SN\\,1979C, SN\\,1986J, SN\\,1993J, and SN\\,2008iz. Although there are only a handful of objects, their detailed study allowed to check and refine the current theoretical models of radio emission in supernovae. This small number of well-observed RSNe may also dramatically increase in the near future, thanks to the forthcoming ultra-sensitive interferometers with a high spatial resolution, like the Square Kilometre Array (SKA). Using the standard model of radio emission from supernovae (Chevalier \\cite{Chevalier1982a},\\cite{Chevalier1982b}), it is possible to relate the defining parameters of the modelled expansion curve to those of the modelled light curves in a simple algebraic way, by assuming an evolution law for the magnetic field (and for the density of the population of synchrotron-emitting electrons) and a radial density profile for the circumstellar medium, CSM, (see, e.g. Weiler et al. \\cite{Weiler2002}). The decay in the radio-light curves according to this model is related to the time decay in the magnetic field and the radial decay of CSM density. However, the continuous energy loss by the relativistic electrons (electron cooling), mainly due to synchrotron radiation (i.e., radiative cooling), but also to adiabatic expansion and inverse-Compton scattering, are not considered in the derivation of this relationship between light curves and expansion curve. Electron cooling may affect considerably the shape of the light curves for a given expansion curve. For instance, Mart\\'i-Vidal et al. (\\cite{MartiVidalII}) succesfully modelled the exponential-like decay of the SN\\,1993J radio light curves at late epochs, reported in Weiler et al. (\\cite{Weiler2007}), using {\\em only} radiative-cooling effects, and assuming that the density of the CSM was negligibly small at large distances to the progenitor star\\footnote{An additional effect due to the escaping of the electrons from the emitting region might also be necessary to model the light curves of SN\\,1993J, were the density of the CSM not negligible at those large distances to the progenitor.}. In any case, it seems clear that if electron cooling is not considered in the modelling of the radio light curves of a supernova, it could result into wrong estimates of the model parameters. In this paper, we study how electron cooling modifies the flux-density decay rate of RSNe for several values of the magnetic field and for different expansion regimes. These results can be used to estimate the magnetic fields of observed RSNe. In the next section, we outline the standard model of radio emission from supernovae. In Sect. \\ref{III} we study the effect of electron cooling in the population of emitting electrons and in the flux-density decay rate. In Sect. \\ref{IV} we present the results of several simulations of the expansion and radio light curves of RSNe. In Sect. \\ref{V}, we explain how these results can be used in real cases to estimate physical quantities in RSNe and estimate the magnetic fields for several observed RSNe, comparing these estimates to those obtained by assuming particle-field energy equipartition. In Sect. \\ref{VI} we summarize our conclusions. ", "conclusions": "\\label{VI} We have shown the impact of energy losses of relativistic electrons in RSNe, and how they affect the flux-density decay rate of the light curves in the optically-thin regime for different values of the magnetic fields and for different expansion curves. If the magnetic-field energy density and the acceleration efficiency of the shock scale with the shock energy density, which is very likely the case for RSNe, we find that there is a tight relation between expansion index, $m$, spectral index, $\\alpha$, and (optically thin) flux-density decay index $\\beta$. This connection between expansion and flux-density evolution in RSNe can be used to estimate the magnetic field of observed RSNe ($B_0$ at a reference epoch) as well as its evolution with time for an assumed CSM radial density profile and energy index, $p$, of the relativistic electrons. For a number of well observed RSNe (e.g., SN 1993J in M81), self-consistent solutions have been found for $B_0$, $m$, $s$, and $p$. A standard CSM density profile (i.e., $s=2$) can explain all observations, although evidences of non-standard values of $s$ are found for SN\\,1986J and SN\\,1979C. The index of the relativistic electron population takes rather high values ($p = 2.3 - 3.0$) and the range of magnetic fields between all cases is large ($B_0 \\sim 20-100$\\,G). These large magnetic fields imply effective amplification mechanisms in the radio-emitting region, possibly related to plasma turbulence (see, e.g., Gull \\cite{Gull} or Jun \\& Norman \\cite{Jun}, and references therein). Previous analyses of the radio light curves and expansion curves of these RSNe did not take into account the correct coupling between $m$, $\\beta$, and $\\alpha$ for different magnetic fields. Some of the results previously reported for these supernovae could, therefore, be internally inconsistent. The magnetic fields obtained with our approach are in similar to the equipartition magnetic fields. For SN\\,1979C and SN\\,1986J, we obtain a range of self-consistent magnetic fields similar to those derived from equipartition with a lower acceleration efficiency for ions (i.e., low-to-intermediate values of $k$ in Eq. \\ref{EquipEq}). Additionally, for SN\\,1986J there is evidence of $s < 2$, provided the magnetic field is small. For SN\\,2008iz, either a very low magnetic field (with $s\\sim 2$) {\\em or} an extremely large magnetic field (with $s > 2$) are necessary to model the light curve, given the large flux-density decay rate ($\\beta =-1.43$). For SN\\,1993J, we obtain a magnetic field similar to that reported in Fransson \\& Bj\\\"ornsson (\\cite{Fransson1998}) and Mart\\'i-Vidal et al. (\\cite{MartiVidalII}), although we use in our approach a subset of flux-density observations (and not the whole data set), to avoid possible biasing effects coming from the ejecta opacity (Mart\\'i-Vidal et al. \\cite{MartiVidalII}). For the RSNe that will be detected in the future (the large sensitivity of the forthcoming radio observatories, like ALMA and SKA, will allow the detection and monitoring of many other RSNe), it will be necessary, in light of the results here reported, to study the connection between their expansion and flux-density evolution, in order to obtain self-consistent results for the CSM profile, the electron energy index, and the magnetic field, based on the observed spectral index, expansion curve, and flux-density decay index." }, "1101/1101.4453_arXiv.txt": { "abstract": "We report new mid-infrared observations of the remarkable object \\irasx\\ using the space telescopes \\akari\\ and \\spitzer, which demonstrate the presence of prominent crystalline silicate emission in this bright source. \\irasx\\ has a complex morphology with a bright central compact source (IRS1) surrounded by knots, spurs, and several extended ($\\sim 4'$) arc-like filaments. The source is seen only at $\\ge 10$~\\um. The \\spitzer\\ mid-infrared (MIR) spectrum of IRS1 shows prominent emission features from Mg-rich crystalline silicates, strong [Ne II] 12.81 \\um\\ and several other faint ionic lines. We model the MIR spectrum as thermal emission from dust and compare with the Herbig Be star HD 100546 and the luminous blue variable R71, which show very similar MIR spectra. Molecular line observations reveal two molecular clouds around the source, but no associated dense molecular cores. We suggest that IRS1 is heated by UV radiation from the adjacent O star Muzzio 10 and that its crystalline silicates most likely originated in a mass outflow from the progenitor of the supernova remnant (SNR) \\msh. IRS1, which is embedded in the SNR, could have been shielded from the SN blast wave if the progenitor was in a close binary system with Muzzio 10. If \\msh\\ is a remnant of Type Ib/c supernova (SN Ib/c), as has been previously proposed, this would confirm the binary model for SN Ib/c. IRS1 and the associated structures may be the relics of massive star death, as shaped by the supernova explosion, the pulsar wind and the intense ionizing radiation of the embedded O star. ", "introduction": "Space infrared observations have revealed distinct spectral features due to crystalline silicates in diverse objects such as Asymptotic Giant Branch (AGB) stars, planetary nebulae, Herbig Ae/Be stars, comets and ultraluminous infrared galaxies \\citep{hen10}. Since silicate dust grains in the interstellar medium are essentially amorphous \\citep{kemper04}, crystalline silicates must form in circumstellar disks and/or outflows of evolved stars or young stellar objects (YSOs). However, neither the formation process of the crystalline dust nor its relation to the central stellar source is understood. In this paper, we report the discovery of prominent crystalline silicate emission in \\irasx, a bright, mid-infrared compact source previously detected by {\\em Infrared Astronomical Satellite} (\\iras) and {\\em Midcourse Space Experiment} (\\msx). The source is located close to the pulsar B1509$-$58 in the supernova remnant (SNR) \\msh\\ (G320.4$-$1.2), a young SNR of complex morphology at a distance of $5.2\\pm 1.4$~kpc \\citep[][see \\S~5]{arendt91, gaensler99}. \\cite{arendt91} concluded that \\irasx\\ is heated either by hot plasma associated with the SNR or by the nearby O star Muzzio 10 \\citep{muzzio79}, which is at a distance $\\sim 4$~kpc \\citep{bessell11}. We show that the latter explanation is most likely correct and that the source is probably associated with the progenitor of the SNR. ", "conclusions": "" }, "1101/1101.1089_arXiv.txt": { "abstract": "{ In the past few years, the extinction law has been measured in the infrared wavelengths for various molecular clouds and different laws have been obtained. } % { In this paper we seek variations of the extinction law within the \\object{Trifid} nebula region. Such variations would demonstrate a local dust evolution linked to variation of the environment parameters such as the density or the interstellar radiation field. } { The extinction values, $A_{\\lambda}/A_V$, are obtained using the 2MASS, UKIDSS and Spitzer/GLIMPSE surveys. The technique is to inter-calibrate color-excess maps from different wavelengths to derive the extinction law and to map the extinction in the Trifid region. } { We measured the extinction law at 3.6, 4.5, and 5.8 $\\mu$m and we found a transition at $A_V\\approx20$~mag. Below this threshold the extinction law is as expected from models for $R_V=5.5$ whereas above 20~mag of visual extinction, it is flatter.Using these results the color-excess maps are converted into a composite extinction map of the Trifid nebula at a spatial resolution of 1~arcmin. A tridimensional analysis along the line-of-sight allowed us to estimate a distance of $2.7\\pm0.5$~kpc for the Trifid. The comparison of the extinction with the 1.25~mm emission suggests the millimeter emissivity is enhanced in the dense condensations of the cloud. } { Our results suggest a dust transition at large extinction which has not been reported so far and dust emissivity variations. } ", "introduction": "During the last three decades, dust evolution in the interstellar medium has been studied mostly through the grain emission in the far-infrared and then the sub-millimeter wavelengths. The analysis of the surface brightnesses at 60 and 100~$\\mu$m observed by IRAS lead to distinguish a small and a big grain component where the small grains are the major contributor at 60~$\\mu$m \\citep{BBD88,LCP91}. The transition from small to big grains occurs at a typical visual extinction of 1~mag. More recently a new transition was discovered in the dust composition with the formation of fluffy grains. The analysis of PRONAOS sub-millimeter observations by \\citet{BAR+99} yielded to the discovery of a population of enhanced-emissivity dust grains which have a lower equilibrium temperature. This emissivity enhancement in the far-infrared by a factor of about 3 \\citep{CBLS01} requires the existence of composite grains formed by the coagulation of small grains. A large scale study on the whole Galactic anticenter hemisphere showed this transition also occurs at low extinction, $A_V\\approx 1$~mag \\citep{CJB05}. It might be closely related to the recent discovery made by \\citet{MSB+10a} with Herschel data where this same density threshold seems to indicate the transition where diffuse clouds start to collapse into filaments. With the wealth of available deep data in the near and mid-infrared it is now possible to investigate the interstellar grain properties at shorter wavelengths. In this spectral domain a lot of progresses has been made in the past few years through the study of the extinction law. Dense clouds can be probed on large scales and for various astrophysical environments. Recent studies with Spitzer data \\citep{IMB+05,FPM+07,RLMA07, CMLE09} focused on different clouds and obtained different extinction laws. We propose in this work to look for possible variations of the extinction curve within a single cloud. We choose an active star-forming region which includes the famous HII region M20, the Trifid nebula. Its location in the Galactic plane at only 7~degrees from the Galactic center offers a high stellar density which is ideal to probe the extinction distribution. The distance of this young nebula ($10^5$ years) is still debated. The observation of the central O7 star responsible for the photoionization, \\object{HD 164492}, yields a distance from 1.67~kpc \\citep{LCO85} to 2.8~kpc \\citep{KML99}. This active region has its first generation of massive stars interacting with the interstellar medium and is very likely triggering a second generation of star formation \\citep{CLC+98}. The state-of-the-art for the Trifid star-forming region is presented in a review paper by \\citet{RLRC08}. The extinction law is generally estimated in molecular clouds as a whole and the discrepancies between the different results are interpreted as variations of the dust optical properties from cloud-to-cloud due to the environment. Obviously variations should also be observed within molecular clouds. As the density increases from the envelope to the core of a cloud, the dust properties, hence the extinction law, should be affected. An attempt to detect such a variation was proposed by \\citet{RLMA07} in their study of a dark cloud core, B59, up to $A_V\\approx59$~mag but they did not succeed. \\citet{OO10} used five different methods to derive the extinction law in B335. The scatter in their result was consistent with the uncertainties so they could not conclude on any evolution of the extinction law within the cloud. \\citet{CMLE09} managed to find variations by averaging their result from several clouds and they showed evidence for grain growth at $A_{K_s}>1$~mag ($A_V>9$). A global galactocentric variation of the extinction law has also been reported by \\citet{ZMI+09} in a large scale study that excluded molecular clouds. The present study focuses on the variations of the extinction law within the molecular cloud associated with the Trifid nebula. Our analysis relies on the 2MASS \\citep{SCS+06}, UKIDSS \\citep{LWA+07} and GLIMPSE II \\citep{CBM+09} surveys. Special attention is given to the data filtering to limit the possible biases as explained in Sect.~\\ref{s.data}. Our original method to derive the extinction law from 3.6 to 5.8~$\\mu$m is based on the color mapping rather than the classical analysis of the source list distribution in a color-color diagram. The details and the advantages of this technique are presented in Sect.~\\ref{s.extlaw} while Sect.~\\ref{s.extmap} describes the final extinction map making and a tridimensional analysis of the Trifid line-of-sight which allows us to disentangle the complexity of this direction. In addition we studied the dust emissivity at 1.25~mm by comparing the masses derived from the dust continuum emission and the dust absorption. We conclude in Sect.~\\ref{s.conclusion}. ", "conclusions": "We performed a detailed analysis of the extinction in the molecular cloud associated with the Trifid nebula. The direction and the distance of this cloud is a major difficulty compared to the other similar investigations published in the literature. To overcome this specificity we proposed an original method which combines the color-excess mapping with the color-color diagram analysis. The gain in sensitivity with this method allowed us to measure to extinction law at 3.6, 4.5 and 5.8~$\\mu$m and to unambiguously detect a variation at large absorption through an abrupt change of slope in the color-color plots. We interpreted this result as an evidence for a rapid dust evolution at high density like a new dust transition phase in the dense interstellar medium. The extinction law parameters, $A_{\\lambda}/A_{K_s}$, are found to be larger in the very dense cores of the cloud, in agreement with dust models. Our values for $A_V>20$~mag are however larger than predicted by WD01 for dense regions ($R_V=5.5$) and larger than previously reported in the literature at 3.6~$\\mu$m. For $A_V<15$~mag, which does still refer to dense material, our values remarkably match those predicted by the model for $R_V=5.5$, as expected. Our study is not sensitive to the diffuse medium for which $R_V=3.1$ with a typical visual extinction as low as 1 mag. Moreover, \\citet{ZMI+09} found a correlation of the extinction law for the diffuse interstellar medium with the galactocentric radius. They concluded the extinction curve follows the dust model with $R_V=5.5$ for a galactocentric radius consistent with the Trifid distance. It suggests the so-called diffuse medium already contains big grains in this part of the Galaxy. If the diffuse medium at the Trifid distance behaves like the dense medium in the solar neighborhood we can imagine the Trifid dense cores are also affected by the distance with a flatter extinction law. Using our varying extinction law we built a composite extinction map of the Trifid nebula region. The maximum visual extinction is 80~mag at 1~arcmin resolution. We have also investigated the matter distribution along the line-of-sight since several small clouds could have mimicked the presence of a single giant massive dark cloud. We found no significant contribution of other clouds in this direction and we were able to estimate the distance of the Trifid nebula at $2.7\\pm0.5$~kpc. The Trifid molecular cloud is about twice the mass of the Orion molecular cloud. The comparison of the dust extinction and 1.25~mm continuum emission in the densest cores suggests the dust emissivity is probably enhanced by a factor of 2-3 in these regions. This cloud appears as a young and massive version of the Orion or Rosette regions. The Trifid represents an earlier evolutionary status in the star-formation process with a first generation of OB stars but no significant embedded cluster yet. Several protostars in TC3/4 are likely the precursors of such clusters. Longer wavelengths observations from Herschel and Planck are needed to better understand the dust properties and the interaction between the interstellar medium and the OB stars." }, "1101/1101.4948_arXiv.txt": { "abstract": "{We elaborate on a minimal inflation scenario based entirely on the general properties of supersymmetry breaking in supergravity models. We identify the inflaton as the scalar component of the Goldstino superfield. We write plausible candidates for the effective action describing this chiral superfield. In particular the theory depends (apart from parameters of $O(1)$) on a single free parameter: the scale of supersymmetry breaking. This can be fixed using the amplitude of CMB cosmological perturbations and we therefore obtain the scale of supersymmetry breaking to be $10^{12-14}$ GeV. The model also incorporates explicit R-symmetry breaking in order to satisfy the slow roll conditions. In our model the $\\eta-$problem is solved without extra fine-tuning. We try to obtain as much information as possible in a model independent way using general symmetry properties of the theory's effective action, this leads to a new proposal on how to exit the inflationary phase and reheat the Universe.} \\begin{document} ", "introduction": "The inflationary paradigm provides a robust framework to explain the size, flatness, homogeneity of the universe and its perturbations \\cite{guth:1981,mukhanov:1981,sato:1981, albrecht.steinhardt:1982,guth/pi:1982,hawking:1982,linde:1982,starobinsky:1982, bardeen/steinhardt/turner:1983}. It postulates that the universe went through a quasi de-Sitter phase in the past leading to an exponential expansion of space-time. If this quasi de-Sitter phase is associated to a scalar field, then it can be shown, rather generally, that vacuum fluctuations of the field are stretched by inflation to produce a nearly scale invariant power spectrum and with sufficient strength to produce the currently observed large scale structure in the universe \\cite{mukhanov:1981,bardeen/steinhardt/turner:1983}. In spite of its successes, we still face several questions to be answered in inflation, namely: what is the inflaton? how do we bring the universe out of the inflationary phase? how can we stop inflation and transit to the decelerating/accelerating universe we live in today? what sets the energy scale of inflation? At a more fundamental level we have to deal with the problem of explaining the rich dynamics of the inflaton without too many free parameters and fine tuning. In a previous letter \\cite{AlvarezGaume:2010rt}, guided by the idea that the inflaton should be found naturally among the fields of any fundamental physics model, we have exploited some generic features of the supersymmetry (supergravity) breaking mechanism to design a model of inflation. In this model we identified the inflaton with the order parameter of supersymmetry breaking and associated the supersymmetry breaking scale with that generating cosmological perturbations. Under these conditions, and imposing explicit R-symmetry breaking, we showed that one can obtain enough number of e-foldings ($> 70$) to explain the observed universe. Supersymmetry and inflation have a long history \\cite{ellisetal, ovrutetal, rossetal, dvalietal, randalletal, lythriotto}, however in our approach we try to avoid making specific and concrete models and try to see to what extent generic features of supersymmetry breaking are enough to provide a good inflationary scenario. Our main motivation to propose to identify the inflaton field with the order parameter of supersymmetry breaking is guided by the fact that, independently of the particular microscopic mechanism driving supersymmetry breaking (in what follows we will restrict ourselves to $F$-breaking ) we can, whenever we have violation of conformal invariance in the UV, define a superfield $X$ whose $ \\theta$ component at large distances becomes the \"Goldstino\" (see \\cite{volkovakulovrocek,seiberg2}). This superfield is the chiral superfield that appears in the divergence of the superfield of currents, the Ferrara-Zumino (FZ) super-multiplet \\cite{ferrarazumino} with universal properties at low energy shared by large classes of models with supersymmetry breaking. In the UV the scalar component $x$ of $X$ is well defined as a fundamental field while in the IR, once supersymmetry is spontaneously broken, this scalar field becomes the superpartner of the Goldstino i.e a two Goldstino state. The realization of $x$ as $GG$ can be implemented by imposing a non linear constraint in the IR for the $X$ field of the type $X^2=0$. In our previous approach to inflation we used one real component of the UV $x$ field as the inflaton. We assumed the existence of a F-breaking effective superpotential for the $X$-superfield and we induced a potential for $x$ from gravitational corrections to the K\\\"ahler potential\\footnote{Like most inflationary theories containing supersymmetry, we present a simple model of multifield inflation (sometimes called hybrid) \\cite{lindehybrid}}. In this paper we refine our model \\cite{AlvarezGaume:2010rt} by exploring a family of K\\\"ahler potentials depending on few parameters and which can lead to a reasonable cosmology. In particular we consider a simple class of models that apart from a few parameters of $O(1)$, depend on the scale of supersymmetry breaking $f$. The K\\\"ahler potential explicitly breaks the R-symmetry as is needed in supergravity in order to have a slow-roll phase where the universe inflates. By fitting these models to cosmological data, we can read a supersymmetry breaking scale of $O(10^{12-13})$ leading to a fairly heavy gravitino. Given the scarcity of parameters in the model presented, it is quite remarkable that many cosmological constraints can be satisfied in such an economical manner. The minimal choice we make has to also provide a graceful exit from inflation without invoking a ``waterfall\" field that will bring the theory out of exponential expansion. In our case, the end of inflation is reached when the $X$-field begins to enter the nonlinear phase ($X^2\\,\\sim\\, 0$), the scalar component of the FZ chiral superfield is converted into a pair of Goldstinos and the state of the universe can in principle be viewed as some weakly interacting Fermi liquid. The Fermi nature of the elementary components of the liquid creates the necessary pressure to exit the inflationary period. This is admittedly a far-off idea, and we are currently exploring its microscopic properties in detail. We expect to report on our results in a future publication \\cite{landauandus}. This detailed physical description is not necessary in order to explore some of the phenomenological properties of our model. This is what we will do in the rest of the paper. Let us insist once more that the key feature of our philosophy is to obtain the basic properties of the inflation scenario (inflation, graceful exit and reheating) out of a single superfield $X$ and the general properties of supersymmetry breaking. ", "conclusions": "" }, "1101/1101.4723_arXiv.txt": { "abstract": "Baryon acoustic oscillations (BAOs) imprinted in the galaxy power spectrum can be used as a standard ruler to determine angular diameter distance and Hubble parameter at high redshift galaxies. Combining redshift distortion effect which apparently distorts the galaxy clustering pattern, we can also constrain the growth rate of large-scale structure formation. Usually, future forecast for constraining these parameters from galaxy redshift surveys has been made with a full 2D power spectrum characterized as function of wavenumber $k$ and directional cosine $\\mu$ between line-of-sight direction and wave vector, i.e., $P(k,\\mu)$. Here, we apply the multipole expansion to the full 2D power spectrum, and discuss how much cosmological information can be extracted from the lower-multipole spectra, taking a proper account of the non-linear effects on gravitational clustering and redshift distortion. The Fisher matrix analysis reveals that compared to the analysis with full 2D spectrum, a partial information from the monopole and quadrupole spectra generally degrades the constraints by a factor of $\\sim1.3$ for each parameter. The additional information from the hexadecapole spectrum helps to improve the constraints, which lead to an almost comparable result expected from the full 2D spectrum. ", "introduction": "Baryon acoustic oscillations (BAOs) imprinted on the clustering of galaxies are now recognized as a powerful cosmological probe to trace the expansion history of the Universe \\cite{Eisenstein:2005su,Percival:2007yw,Percival:2009xn}. In particular, the BAO measurement via a spectroscopic survey can provide a way to simultaneously determine the angular diameter distance $D_A$ and Hubble parameter $H$ at given redshift of galaxies through the cosmological distortion, known as Alcock-Paczynski effect (e.g., \\cite{Alcock_Paczynski:1979,Seo:2003pu,Blake:2003rh,Shoji:2008xn, Padmanabhan:2008ag}). Further, measuring the clustering anisotropies caused by the redshift distortion due to the peculiar velocity of galaxies, we can also probe the growth history of structure formation (e.g., \\cite{Linder:2007nu,Guzzo:2008ac,Yamamoto:2008gr,Song:2008qt}), characterized by the growth-rate parameter $f\\equiv d\\ln D/d\\ln a$, with quantities $D$ and $a$ being linear growth factor and the scale factor of the Universe, respectively. With the increased number of galaxies and large survey volumes, on-going and future spectroscopic galaxy surveys such as Baryon Oscillation Spectroscopic Survey (BOSS) \\cite{Schlegel:2009hj}, Hobby-Eberly Dark Energy Experiment (HETDEX) \\cite{Hill:2008mv}, Subaru Measurement of Imaging and Redshift equipped with Prime Focus Spectrograph (SuMIRe-PFS), and EUCLID/JDEM \\cite{Beaulieu:2010qi,Gehrels:2010fn} aim at precisely measuring the acoustic scale of BAOs as a standard ruler. These surveys will cover the wide redshift ranges, $0.3\\lesssim z\\lesssim 3.5$, and provide a precision data of the redshift-space power spectrum with an accuracy of a percent level over the scales of BAOs. In promoting these gigantic surveys, a crucial task is a quantitative forecast for the size of the statistical errors on the parameters $D_A$, $H$ and $f$ in order to clarify the scientific benefits as well as to explore the optimal survey design. The Fisher matrix formalism is a powerful tool to investigate these issues, and it enables us to quantify the precision and the correlation between multiple parameters (\\cite{Seo:2003pu,Seo:2007ns,White:2008jy,Shoji:2008xn}, especially for measuring $D_A$, $H$ and $f$). So far, most of the works on the parameter forecast study have focused on the potential power of the BAO measurements, and attempt to clarify the achievable level of the precision for the parameter estimation. For this purpose, they sometimes assumed a rather optimistic situation that a full shape of the redshift-space power spectrum, including the clustering anisotropies due to the redshift distortion, is available in both observation and theory. In this paper, we are particularly concerned with the parameter estimation using a partial information of the anisotropic BAOs from a practical point-of-view. In redshift space, the power spectrum obtained from the spectroscopic measurement is generally described in the two dimension, and is characterized as functions of $k$ and $\\mu$, where $k$ is the wavenumber and $\\mu$ is the directional cosine between the line-of-sight direction and $k$ \\footnote{Throughout the paper, we work with the distant-observer approximation, and neglect the angular dependence of the line-of-sight direction, relevant for the high-redshift galaxy surveys.}. While most of the forecast study is concerned with a full 2D power spectrum, the multipole expansion of redshift-space power spectrum has been frequently used in the data analysis to quantify the clustering anisotropies. Denoting the power spectrum by $P(k,\\mu)$, we have \\begin{align} &P(k,\\mu)=\\sum_{\\ell=0}^{\\rm even} P_\\ell(k)\\,\\mathcal{P}_\\ell(\\mu) \\label{eq:multipole_pk} \\end{align} with the function $\\mathcal{P}_\\ell$ being the Legendre polynomials. Although the analysis with full 2D spectrum will definitely play an important role as improving the statistical signal, most of the recent cosmological data analysis has focused on the angle-averaged power spectrum $(\\ell=0)$, i.e., monopole spectrum, and a rigorous analysis with full 2D spectrum is still heavy task due to the time-consuming covariance estimation (e.g., \\cite{Okumura:2007br,Takahashi:2009bq,Cabre:2008sz}). In linear theory, the redshift-space power spectrum is simply written as $P(k,\\mu)=(1+\\beta\\,\\mu^2)^2P_{\\rm gal}(k)$, where $\\beta=f/b$ with $b$ being the linear bias parameter, and $P_{\\rm gal}$ is the galaxy power spectrum in real space \\cite{Kaiser:1987qv,1992ApJ385L5H,Hamilton:1997zq}. Then, the non-vanishing components arises only from the monopole ($\\ell=0$), quadrupole ($\\ell=2$) and hexadecapole spectra ($\\ell=4$). That is, cosmological information contained in the $\\ell=0$, $2$ and $4$ moments is equivalent to the whole information in the full 2D power spectrum. Observationally, however, this is only the case when we a priori know the cosmological distance to the galaxies. The Alcock-Paczynski effect can induce non-trivial clustering anisotropies, which cannot be fully characterized by the lower multipole spectra, in general. Further, in reality, linear theory description cannot be adequate over the scale of the BAOs, and the non-linear effects of the redshift distortion as well as the gravitational clustering must be accounted for a proper comparison with observation. These facts imply that non-vanishing multipole spectra higher than $\\ell>4$ generically appear, and a part of the cosmological information might be leaked into those higher multipole moments. An important question is how much amount of the cosmological information can be robustly extracted from the lower multipole spectra instead of the full 2D spectrum. In the light of this, Ref.~\\cite{Padmanabhan:2008ag} recently examined a non-parametric method to constrain $D_A$ and $H$ from the monopole and quadrupole spectra, and numerically estimate the size of errors (see also Ref.~\\cite{TocchiniValentini:2011mt} for the estimation of growth-rate parameter). Here, as a complementary and comprehensive approach, we will investigate this issue based on the Fisher matrix formalism, and derive the useful formulae for parameter forecast using the multipole power spectra. We then explore the potential power of the lower multipole spectra on the cosmological constraints, particularly focusing on the parameters $D_A$, $H$ and $f$. To do so, we consider the Figure-of-Merit (FoM) and Figure-of-Bias (FoB) for these parameters, and investigate their dependence on the assumptions for the number density of galaxies, the amplitude of clustering bias, the maximum wavenumber used for the parameter estimation. In Sec.~\\ref{sec:Fisher_formalism}, we present the Fisher matrix formalism for cosmological parameter estimation from the multipole power spectra. Sec.~\\ref{sec:model_assumption} deals with the model of redshift-space power spectrum and the assumptions used in the Fisher matrix analysis. Then, in Sec.~\\ref{sec:results}, the results for FoM and FoB are shown, and the sensitivity of the results to the assumptions and choice of the parameters is discussed in greater details. Finally, Sec.~\\ref{sec:summary} briefly summarize our present work. Throughout the paper, we assume a flat Lambda cold dark matter (CDM) model, and the fiducial model parameters are chosen based on the five-year WMAP results \\cite{Komatsu:2008hk}: $\\Omega_{\\rm m}=0.279$, $\\Omega_{\\Lambda}=0.721$, $\\Omega_{\\rm b}=0.0461$, $h=0.701$, $n_s=0.96$, $A_s=2.19\\times10^{-9}$. ", "conclusions": "\\label{sec:summary} In this paper, we have studied the cosmological constraints from the anisotropic BAOs based on the multipole expansion of redshift-space power spectrum. We have derived the several formulae for the Fisher analysis using the multipole power spectra; Eqs.~(\\ref{eq:formula_F_ij}) and (\\ref{eq:cov_formula}) for the Fisher matrix, and Eqs.~(\\ref{eq:systematic_bias}) and (\\ref{eq:vector_s}) for the estimation of systematic biases. We then consider the hypothetical galaxy survey of $V_s=4h^{-3}$Gpc$^3$ and $z=1$, and discuss the potential power of the lower multipole spectra on the cosmological constraints, particularly focusing on the parameters $D_A$, $H$ and $f$. Compared to the analysis with full 2D power spectrum, a partial information from the monopole and quadrupole power spectra generally degrades the constraints on $D_A$, $H$, and $f$. Typically, the constraint is degraded by a factor of $\\sim1.3$ for each parameter. The interesting finding is that adding the information from hexadecapole spectra ($P_4$) to that from the monopole and quadrupole spectra greatly improves the constraints, and the resultant constraints would become almost comparable to those expected from the full 2D power spectrum (see Fig.~\\ref{fig:FoM}). Note also that the situation would be relatively improved depending on the properties of galaxy samples, and for highly biased galaxy samples with $b\\sim4$, the total power of the constraints defined by the Figure-of-Merit [FoM, Eq.~(\\ref{eq:def_FoM})] can reach $\\sim80\\%$ of the one expected from the full 2D power spectrum. We have also investigated the impacts of systematic biases on the best-fit values of $D_A$, $H$ and $f$. The incorrect model of redshift distortion tends to produce a large systematic bias in the growth-rate parameter, and the size of biases would be rather significant for the analysis with full 2D spectrum. An interesting suggestion is that the situation would be greatly relaxed if we only use the combination of monopole and quadrupole spectra, and the estimated value of Figure-of-Bias defined by Eq.~(\\ref{eq:FoB}) is mostly below the critical value for stage-III class surveys (Fig.~\\ref{fig:FoB}). In this respect, the analysis with partial information from monopole and quadrupole may be still helpful in cross-checking the results derived from the full 2D power spectrum. On the other hand, wrong prior assumption of cosmological parameters in computing the template power spectrum severely affects the acoustic-scale determination, and a percent-level precision is required for the prior information in order to avoid a large systematic biases on $D_A$ and $H$ (Fig.~\\ref{fig:FoB_cosmoparams}). This is true irrespective of the choice of template power spectra used in the analysis. Finally, we note that the assumptions and situations considered in the paper are somewhat optimistic or too simplistic, and a more careful study is needed for a quantitative parameter forecast. One critical aspect is the modeling of the galaxy power spectrum. In reality, the assumption of linear and deterministic galaxy biasing is idealistic, and the scale-dependence or non-linearity/stochasticity of the galaxy biasing should be consistently incorporated into the theoretical template of redshift-space power spectrum. Although this is tiny effect for the scale of our interest, the distance information, $D_A$ and $H$, is rather sensitive to a slight modification of the acoustic structure in the power spectrum, and results in this paper might be somehow changed. A more elaborate modeling for power spectrum is thus quite essential." }, "1101/1101.3335_arXiv.txt": { "abstract": "We present new mid-infrared (MIR) imaging data for three Type-1 Seyfert galaxies obtained with T-ReCS on the Gemini-South Telescope at subarcsecond resolution. Our aim is to enlarge the sample studied in a previous work to compare the properties of Type-1 and Type-2 Seyfert tori using clumpy torus models and a Bayesian approach to fit the infrared nuclear spectral energy distributions (SEDs). Thus, the sample considered here comprises 7 Type-1, 11 Type-2, and 3 intermediate-type Seyferts. The unresolved IR emission of the Seyfert 1 galaxies can be reproduced by a combination of dust heated by the central engine and direct AGN emission, while for the Seyfert 2 nuclei only dust emission is considered. These dusty tori have physical sizes smaller than 6 pc radius, as derived from our fits. Unification schemes of AGN account for a variety of observational differences in terms of viewing geometry. However, we find evidence that strong unification may not hold, and that the immediate dusty surroundings of Type-1 and Type-2 Seyfert nuclei are intrinsically different. The Type-2 tori studied here are broader, have more clumps, and these clumps have lower optical depths than those of Type-1 tori. The larger the covering factor of the torus, the smaller the probability of having direct view of the AGN, and vice-versa. In our sample, Seyfert 2 tori have larger covering factors (C$_T$=0.95$\\pm$0.02) and smaller escape probabilities (P$_{esc}$=0.05$\\pm^{0.08}_{0.03}$ \\%) than those of Seyfert 1 (C$_T$=0.5$\\pm$0.1; P$_{esc}$=18$\\pm$3 \\%). All the previous differences are significant according to the Kullback-Leibler divergence. Thus, on the basis of the results presented here, the classification of a Seyfert galaxy as a Type-1 or Type-2 depends more on the intrinsic properties of the torus rather than on its mere inclination towards us, in contradiction with the simplest unification model. ", "introduction": "\\label{intro} Observational evidence in the X-rays and the MIR indicates that the strong AGN continuum source must be absorbed by obscuring material over a wide solid angle (see e.g., \\citealt{Antonucci85,Maiolino98,Risaliti02}). According to observed spectra of different AGN types, the obscuring structure has to block the emission of the subparsec-scale Broad-Line Region (BLR) where the broad lines are produced, but not that of the kiloparsec-scale Narrow-Line Region (NLR). The unified model for active galaxies \\citep{Antonucci93,Urry95} is based on the existence of a dusty toroidal structure surrounding the central region of AGN. This toroidal geometry explains the biconical shapes observed in Hubble Space Telescope (HST) imaging of several AGNs \\citep{Tadhunter89,Malkan98,Tadhunter99} and also the polarimetric observations \\citep{Antonucci85,Packham97}. Thus, considering this geometry of the obscuring material, the central engines of Type-1 AGN can be seen directly, resulting in typical spectra with both narrow and broad emission lines, whereas in Type-2 AGN the BLR is obscured. Pioneering work in modelling the dusty torus \\citep{Pier92,Pier93,Granato94,Efstathiou95,Granato97,Siebenmorgen04} assumed a uniform dust density distribution to simplify the modelling, although from the start, \\citet{Krolik88} realized that smooth dust distributions cannot survive within the AGN vicinity. They proposed instead that the material in the torus must be distributed in a clumpy structure, in order to prevent the dust grains from being destroyed by the hot surrounding gas. The IR range (and particularly the MIR) is key to set constraints on the torus models, since the reprocessed radiation from the dust in the torus is re-emitted in this range. However, in comparing the predictions of any torus model with observations, its small-scale emission must be isolated. High angular resolution is then essential to separate torus emission from stellar emission and star-heated dust in the near-IR (NIR) and MIR, respectively. Indeed, starlight dominates the nuclear NIR emission of Seyfert 2 galaxies when using large aperture data (see e.g., \\citealt{Alonso96}) and still has a significant contribution for Seyfert 1 galaxies \\citep{Kotilainen92}. Similar contamination problems can be present in the MIR with the star-heated dust and dust in the ionization cones \\citep{Alonso06,Mason06}. Another controversial issue about the torus structure is its typical dimensions. \\citet{Pier93} and \\citet{Granato94} reproduced the infrared observations of nearby Seyfert galaxies with $\\sim$100 pc scale tori. However, hard X-ray observations showed that about half of nearby Type-2 Seyferts are Compton-thick (i.e., they are obscured by a column density higher than 10$^{24}~cm^{-2}$; \\citealt{Risaliti99}). For these highly obscured sources the torus dimensions are expected to be of a few parsecs, because otherwise the dynamical mass of the obscuring material would be too large to be realistic \\citep{Risaliti99}. In addition, recent ground-based MIR observations of nearby Seyferts reveal that the torus size is likely restricted to a few parsecs. \\citet{Packham05} and \\citet{Radomski08} established upper limits of 2 and 1.6 pc for the outer radii of the Circinus galaxy and Centaurus A tori, respectively. Besides, interferometric observations obtained with the MIR Interferometric Instrument (MIDI) at the Very Large Telescope Interferometer (VLTI) of Circinus, NGC 1068, and Centaurus A suggest a scenario where the torus emission would only extend out to R = 1 pc \\citep{Tristram07}, R = 1.7 - 2 pc \\citep{Jaffe04,Raban09}, and R = 0.3 pc \\citep{Meisenheimer07}, respectively. In order to solve the discrepancies between observations and previous models, an intensive search for an alternative torus geometry has been carried out in the last decade. The first results of radiative transfer calculations of a clumpy rather than a smooth medium were reported by \\citet{Nenkova02} and \\citet{Elitzur06}, and further work was done by \\citet{Dullemond05}. The clumpy dusty torus models \\citep{Nenkova02,Nenkova08a,Nenkova08b,Honig06,Schartmann08} propose that the dust is distributed in clumps, instead of homogeneously filling the torus volume. These models are making significant progress in accounting for the MIR emission of AGNs (\\citealt{Mason06,Mason09,Mor09,Horst08,Horst09,Nikutta09}; Ramos Almeida et al.~2009a; \\citealt{Honig10}). In our previous work (Ramos Almeida et al.~2009a; hereafter \\citealt{Ramos09a}), we constructed subarcsecond resolution IR SEDs for eighteen Seyfert galaxies, mostly Type-2 Seyferts. From the comparison between our high angular resolution MIR fluxes and large aperture data, such as those from ISO, IRAS, or Spitzer, we confirmed that the former provide a spectral shape that is substantially different from that of the large aperture data \\citep{Rodriguez97}. Since our nuclear measurements allowed us to better characterize the torus emission, we modelled our SEDs with clumpy torus models. In general, we found that Type-2 views are more inclined than those of Type-1s, and more importantly, we derive larger covering factors for the Type-2 tori (i.e., more clumps and wider torus angular distributions). This would imply that the observed differences between Type-1 and Type-2 AGN would not be due to orientation effects only, but to intrinsic differences in their tori. However, due to the limited size of the sample analyzed by \\citealt{Ramos09a}, and in particular of Type-1 Seyferts compared with Type-2s, our aim is to enlarge the sample studied in the previous work with new Seyfert 1 infrared data to compare the properties of Type-1 and Type-2 Seyfert tori. In this work, we report new subarcsecond MIR imaging data for the 3 nearby Type-1 Seyfert galaxies NGC 7469, NGC 6221, and NGC 6814, for which we estimate unresolved nuclear MIR fluxes. We enlarge the sample by including the galaxies NGC 1097, NGC 1566, NGC 3227, and NGC 4151, which have similar MIR data, and we compile NIR nuclear fluxes from the literature of similar resolution to construct nuclear SEDs for all the galaxies. We fit these SEDs with clumpy torus models which we interpolate from the \\citet{Nenkova08a,Nenkova08b} database, and compare them with the larger sample studied by \\citealt{Ramos09a}. Table \\ref{sources} summarizes key observational properties of the sources in the sample. Section \\ref{observations} describes the observations, data reduction, and compilation of NIR and MIR fluxes. Sections \\ref{extended} and \\ref{sed} present the main observational results, and in \\S \\ref{modelling} we report the modelling results. We discuss differences between Type-1 and Type-2 Seyferts and draw conclusions about the clumpy torus models and AGN obscuration in general in \\S \\ref{discussion}. Finally, Section \\ref{final} summarizes the main conclusions of this work. \\begin{deluxetable*}{lcccccc} \\tablewidth{0pt} \\tablecaption{Basic Galaxy Data} \\tablehead{ \\colhead{Galaxy} & \\colhead{Seyfert Type} & Ref. & \\colhead{$z$} & \\colhead{Distance} & \\colhead{Scale} & Ref. \\\\ & & & & \\colhead{(Mpc)} & \\colhead{(pc~arcsec$^{-1}$)} & } \\startdata NGC 1097 & Sy1\\tablenotemark{\\dag}& A1& 0.0042 & 19 & 92 & B1 \\\\ NGC 1566 & Sy1\t\t & A2 & 0.0050 & 20 & 97 & B2 \\\\ NGC 6221 & Sy1\t\t & A3 & 0.0050 & 18 & 87 & B3 \\\\ NGC 6814 & HII/Sy1.5\t & A4 & 0.0052 & 21 & 102 & B4 \\\\ NGC 7469 & Sy1 \t\t & A5 & 0.0163 & 65 & 315 & B5 \\\\ \\hline NGC 3227 & Sy1.5 & A6 & 0.0039 & 17 & 82 & B6 \\\\ NGC 4151 & Sy1.5 & A7 & 0.0033 & 13 & 64 & B7 \\\\ \\enddata \\tablecomments{\\footnotesize{Classification and distance are taken from the literature (references below) and spectroscopic redshift from the NASA/IPAC Extragalactic Database (NED).}} \\tablenotetext{\\dag}{\\footnotesize{Originally classified as a LINER by \\citet{Keel83} and \\citet{Phillips84}.}} \\tablerefs{\\footnotesize{(A1) \\citet{Storchi97}; (A2) \\citet{Kriss90}; (A3) \\citet{Levenson01}; (A4) \\citet{Veron06}; (A5) \\citet{Osterbrock93}; (A6) \\citet{Rubin68}; (A7) \\citet{Ayani91}; (B1) \\citet{Willick97}; (B2) \\citet{Sandage94}; (B3) \\citet{Koribalski04}; (B4) \\citet{Liszt95}; (B5) \\citet{Heckman86}; (B6) \\citet{Garcia93}; (B7) \\citet{Radomski03}.}} \\label{sources} \\end{deluxetable*} ", "conclusions": "" }, "1101/1101.3429_arXiv.txt": { "abstract": "It is generally believed that turbulence has a significant impact on the dynamics and evolution of molecular clouds and the star formation which occurs within them. Non-ideal magnetohydrodynamic effects are known to influence the nature of this turbulence. We present the results of a suite of $512^3$ resolution simulations of the decay of initially super-Alfv\\'enic and supersonic fully multifluid MHD turbulence. We find that ambipolar diffusion increases the rate of decay of the turbulence while the Hall effect has virtually no impact. The decay of the kinetic energy can be fitted as a power-law in time and the exponent is found to be $-1.34$ for fully multifluid MHD turbulence. The power spectra of density, velocity and magnetic field are all steepened significantly by the inclusion of non-ideal terms. The dominant reason for this steepening is ambipolar diffusion with the Hall effect again playing a minimal role except at short length scales where it creates extra structure in the magnetic field. Interestingly we find that, at least at these resolutions, the majority of the physics of multifluid turbulence can be captured by simply introducing fixed (in time and space) resistive terms into the induction equation without the need for a full multifluid MHD treatment. The velocity dispersion is also examined and, in common with previously published results, it is found not to be power-law in nature. ", "introduction": "\\label{sec:intro} Turbulence is recognized as a possible source of support against gravitational collapse for molecular clouds. The precise role and source of the observed motions interpreted as evidence of turbulence in these clouds has been studied extensively by many researchers (see the reviews of \\citealt{mac04,elm04}). Clearly, if turbulence can support molecular clouds then it can influence star formation in terms of rate, efficiency and initial mass function \\citep{elm93,kle03}. Many studies of turbulence in molecular clouds have focused on ideal magnetohydrodynamics (MHD) as an approximation of the physics governing this system \\citep{maclow98, maclow99, ost01, ves03, gus06, glover07, lem08, lem09}. The assumption of ideal MHD, while desirable for technical reasons, is perhaps risky in the context of turbulence. The reason for this is that while ideal MHD is valid in molecular clouds on fairly large length scales, on shorter length scales non-ideal effects are thought to become significant \\citep{wardle04, ois06}. Given that turbulence in 3 dimensions involves the transfer of energy from large scales to ever smaller scales, the assumption of ideal MHD will be invalid below some critical spatial scale and the correct nature of the energy cascade may not be observed at this range. The most important of the non-ideal effects for molecular cloud dynamics is ambipolar diffusion. Some authors \\citep{ois06, li08, kud08} have studied driven MHD turbulence in the presence of ambipolar diffusion. All these authors find that ambipolar diffusion produces significant differences in the properties of the turbulence. While most likely of lesser significance, it has been suggested that although the Hall resistivity is generally at least an order of magnitude lower than the ambipolar resistivity in molecular clouds \\citep{wardle04}, its effect should not be ignored. Although relatively weak, it is capable of inducing topological changes in the magnetic field which are quite distinct to any influence caused by ambipolar diffusion. In support of this assertion, we note that researchers working on reconnection and the solar wind have studied the Hall effect in the context of turbulence and found that, although the overall decay rate appears not to be affected, the usual coincidence of the magnetic and velocity fields seen in MHD does not occur at small scales \\citep{mat03, min06, ser07}. Almost no work has been done on comparing the influences of this effect coupled with that of ambipolar diffusion on turbulence with the exception of \\citet[hereafter Paper I]{dos09}. In Paper I a series of simulations of decaying supersonic non-ideal MHD turbulence incorporating both ambipolar diffusion and the Hall effect were performed. These simulations, however, were constrained in that the resistivities associated with each of ambipolar diffusion, the Hall effect and the Pederson resistivity were kept fixed in both space and time. The authors found that, at length scales of 0.2\\,pc, ambipolar diffusion has a significant impact on the decay of the turbulence. The Hall effect was less significant in this respect but does have an influence on the magnetic field at short length scales. Here we present simulations in which the resistivities are self-consistently calculated from the evolution of both the magnetic field and the densities of all of the component species of the fluid. Using these dynamically evolving resistivities we study the decay of fully multifluid MHD turbulence. This is the first such study presented in the literature, with the exception of the low resolution simulations presented by \\cite{dos08}. The aim of this work is to examine in detail the differences between the decay of ideal MHD turbulence and that of multifluid MHD turbulence with a full tensor resistivity incorporating the effects of ambipolar diffusion, the Hall effect and Ohmic resistivity. We will use the results of Paper I in our discussion of these differences as it represents an intermediate stage between the calculations presented here and those of ideal MHD. This work is new in two respects: notwithstanding Paper I, no previous work has focused on {\\em decaying} (i.e.\\ un-driven) {\\em multifluid} MHD turbulence and, in addition, no previous work has addressed the issue of turbulence in the presence of both ambipolar diffusion and the Hall effect simultaneously. In section \\ref{sec:num-method} we outline the numerical techniques used in this work, as well as the initial conditions and general set-up for the simulations while in section \\ref{sec:analysis} we describe the methods used to analyze the simulation data. In section \\ref{sec:results} we present and discuss the results of our simulations of turbulent decay. Finally, section \\ref{sec:conclusions} contains a summary of our results. ", "conclusions": "\\label{sec:conclusions} We have presented results from a suite of $512^3$ resolution simulations of fully multifluid MHD decaying turbulence. The effects incorporated include the Hall effect and ambipolar diffusion. We have performed a resolution study to ensure that the energy decay rate, being the main result presented here, is reliable. We have confirmed the results of the simplified calculations in Paper I that the Hall effect has little impact on the nature and behavior of turbulence in molecular clouds under the well motivated physical parameters assumed in this work. Further, the presence of ambipolar diffusion increases the rate of energy decay at length scales of 0.2\\,pc and less. The same conclusion is drawn for the behavior of the energy in the magnetic field. The power spectra for these simulations again suggest that the Hall effect has little impact on the flows with the exception of the spectrum of magnetic field variations. We must keep in mind that the maximum resolution used here ($512^3$) is only enough to resolve about half a decade in $k$-space and it is therefore difficult to be confident of the details of the power spectra. Notwithstanding this consideration it does appear clear that ambipolar diffusion steepens the power spectra of the neutral velocity, density and the magnetic field. As noted in Paper I, it appears that at a resolution of $512^3$ and an assumed length scale of 0.2\\,pc we have resolved the length at which ambipolar diffusion begins to influence the turbulent cascade. In Paper I only constant resistivities were implemented and hence it was unclear whether this latter result would survive the inclusion of more realistic fully multifluid MHD in which the resistivities vary strongly in space and time. The results presented here imply that it does. The power spectra of the neutral velocity and the magnetic field differ qualitatively from that of the density with breaks occurring in the former which are not seen in the latter. This suggests a decoupling between these fields. The velocity dispersion as a function of length does not behave as a power law. This is not unexpected as the nature of MHD turbulence implies a wide range of applicable signal speeds which can, when combined, remove the power-law behavior which might be expected if only one signal speed were relevant." }, "1101/1101.2808_arXiv.txt": { "abstract": "The determination of the stellar parameters of M dwarfs is of prime importance in the fields of galactic, stellar and planetary astronomy. M stars are the least studied galactic component regarding their fundamental parameters. Yet, they are the most numerous stars in the galaxy and contribute to most of its total (baryonic) mass. In particular, we are interested in their metallicity in order to study the star-planet connection and to refine the planetary parameters. As a preliminary result we present a test of the metallicity calibrations of \\cite{B05}, \\cite{JA09}, and \\cite{SL10} using a new sample of 17 binaries with precise V band photometry. ", "introduction": " ", "conclusions": "" }, "1101/1101.2349_arXiv.txt": { "abstract": "The multi-faceted contributions of Dr. Peter Scheglov (1932-2002) in the area of site testing are briefly reviewed. He discovered and studied astronomical sites in the Central Asia, developed new site-testing instruments, promoted new methods and techniques among his colleagues and teached new generation of observational astronomers. ", "introduction": "\\begin{figure}[h] \\begin{tabular}{cc} \\parbox[b]{9.5cm}{ % \\begin{itemize}\\itemsep=-3pt \\item[1932 --] On September 4, 1932, Petr (Peter) Vladimirovich Scheglov is born in Tashkent (Uzbekistan) \\item[1954 --] P.S. graduated from the Moscow University, chair of Astrophysics \\item[1957 --] PhD in astronomy (adviser -- I.S.~Shklovsky). P.S. starts working at the Sternberg Astronomical Institute (Moscow). \\item[1966 --] P.S. meets with Jurgen Stock at the IAU General Assembly in Prague, sparkling his interest in site-testing \\item[1967-70 --] Development of the double-beam instrument (DBI) and first site-testing missions to Maidanak, Sanglok, Alma-Ata and Crimea. \\item[1970 --] Dr.Sci dissertation \\item[1974 --] Sternberg Institude decides to build its observatory at Maidanak. \\item[1975 --] Another mission, choice of the observatory location on the West summit of Maidanak. \\item[1976-80 --] Development of the photoelectric seeing monitor (FEP) \\item[1980s --] Seeing measurements at various sites. Evaluation of the ground-layer turbulence with micro-thermal sensors and Sodar. \\item[1980 --] P.S. publishes his book ``Problems of optical astronomy'' \\item[1990 --] Comprehensive study of turbulence at Maidanak \\item[2002 --] On December 2, 2002 Petr Vladimirovich Scheglov passed away \\end{itemize} } & \\includegraphics[height=10cm]{fig1.eps}\\\\ \\end{tabular} \\end{figure} ", "conclusions": "" }, "1101/1101.1456_arXiv.txt": { "abstract": "The detection and flux estimation of point sources in cosmic microwave background (CMB) maps is a very important task in order to clean the maps and also to obtain relevant astrophysical information. In this paper we propose a maximum a posteriori (MAP) approach detection method in a Bayesian scheme which incorporates prior information about the source flux distribution, the locations and the number of sources. We apply this method to CMB simulations with the characteristics of the Planck satellite channels at 30, 44, 70 and 100 GHz. With a similar level of spurious sources, our method yields more complete catalogues than the matched filter with a $5\\sigma$ threshold. Besides, the new technique allows us to fix the number of detected sources in a non-arbitrary way. ", "introduction": "\\label{sec:intro} The detection and estimation of the intensity of compact objects embedded in a background plus instrumental noise is a problem of interest in many different areas of science and engineering. A classic example is the detection of point-like extragalactic objects --i.e. galaxies-- in sub-millimetric Astronomy. Regarding this particular field of interest, different techniques have proven useful in the literature. Some of the existing techniques are: the standard matched filter \\citep[MF,][]{MF_radio92}, the matched multifilter \\citep{herr02a,lanz10} or the recently developed matched matrix filters \\citep{herranz08a}. Other methods include continuous wavelets like the standard Mexican Hat \\citep{wsphere} and other members of its family \\citep{MHW206}. All these filters have been applied to real data of the Cosmic Microwave Background (CMB), like those obtained by the WMAP satellite \\citep{NEWPS07} and CMB simulated data \\citep{challenge08} for the experiment on board the \\emph{Planck} satellite \\citep{planck_tauber05}. Besides, Bayesian methods have also been recently developed \\citep{hob03,psnakesI}. A more detailed review on point source detection techniques in microwave and sub-mm Astronomy, with a more complete list of references, can be found in \\cite{review2010}. When a MF or a wavelet is applied to a CMB map in the blind detection case, i.e. when it is assumed that the number of point sources, their positions and fluxes are unknown, the most common method for detection is based on the well-known idea of thresholding: the maxima of the filtered map above a given threshold are selected and considered as the positions of the sources, so that the number of detected sources is the number of maxima above that threshold. The fluxes are estimated then by using the corresponding estimation formulas with the MF or the wavelet. The value of this threshold remains arbitrary, though a $5\\sigma$ cut is often applied, since it guarantees that under reasonable conditions a few detected sources are spurious. Apart from the arbitrariness of this procedure, the prior knowledge regarding the average number of sources in the surveyed patch, the flux distribution of these sources or other properties are not used, so that useful information is being neglected. Bayesian detection techniques provide a natural way to take into account all the available information about the statistical distribution of both the sources and the noise. Unfortunately, up to this date only a few works have addressed the problem of detecting extragalactic point sources in CMB data \\citep{hob03,psnakesI}. The reason for this is twofold: on the one hand, the statistical properties of extragalactic sources at sub-mm frequencies are still very poorly known. On the other hand, mapping the full posterior probability density of the sources is often very difficult and computationally expensive. These two problems explain, at least partially, the predominance of frequentist over Bayesian methods in the literature. Let us consider the previous two problems separately: The microwave and sub-mm region has been until very recently one of the last uncharted areas in astronomy. Concerning extragalactic sources, this region of the electromagnetic spectrum is where the total number of counts passes from being dominated by radio-loud galaxies to being dominated by dusty galaxies. Although a minimum of the emission coming from extragalactic sources is expected to occur around 100--300 GHz, they are still considered as the main contaminant of the CMB at small angular scales at these frequencies \\citep{tof98,zotti05}. The uncertainties about the number counts at intermediate and low flux, redshift distribution, evolution and clustering properties of this mixed population of objects are large. In most cases this has motivated the use of noninformative priors, which avoid to make adventurous assumptions about the sources but on the other hand miss part of the power of the priors that are based on observations and physical intuition. But, in spite of what has been said above, our knowledge about the statistical properties of point sources is growing day by day thanks to the new generation of surveys and experiments. In the high-frequency radio regime, WMAP observations are in agreement with the de Zotti model \\citep{zotti05,gnuevo08}. Priors for the number density and flux distributions in the range of frequencies $ > 5 $ GHz are more and more reliable thanks to the information provided by recent surveys such as CRATES at 8.4 GHz \\citep{CRATES}, the Ryle-Telescope 9C at 15.2 GHz \\citep{taylor01,waldram03} or the AT20G survey at 20 GHz \\citep{ricci04,ATCA_BSS,ATCA10}. For a recent review on radio and millimiter surveys and their astrophysical implications, see \\cite{review_dezotti10}. The situation is worse in the far-infrared part of the spectrum, where relatively large uncertainties remain in the statistical properties, the evolution and, above all, the clustering properties of dusty galaxies. Most of the existing dusty galaxy surveys have been carried out in the near and medium infrared with IRAS, ISO and Spitzer, but the wave-band from 60 to 500 $\\mu$m is still virtually \\emph{terra incognita}. The only survey of a large area of the extragalactic sky at a wavelength above 200 $\\mu$m is the one recently carried out by the Herschel pathfinder experiment, the Balloon Large Area Survey Telescope \\citep[BLAST,][]{BLAST}. In the next few months, however, the luminosity function and the dust-mass function of dusty galaxies in the nearby Universe will be much better understood thanks to the Herschel-ATLAS Survey \\citep{ATLAS}, which covers the wavelength range between 110 and 500 $\\mu$m and has already produced interesting results during the Herschel Science Demonstration Phase \\citep{ATLAScounts}. Thanks to these and the previously mentioned observations, the sub-mm gap is narrowing and our knowledge of galaxy populations in this wave band, albeit far from perfect, is quickly improving. Apart from the uncertainties on the priors, the other complication that has traditionally deterred microwave astronomers from attempting Bayesian point source detection is computational and algorithmic complexity. Depending on the choice of priors and the likelihood function, the full posterior distribution of the parameters of the sources may be very complex and in most cases it is impossible to obtain maximum a posteriori (MAP) values of the parameters and their associated errors via analytical equations. Numerical sampling techniques such as Monte Carlo Markov Chain (MCMC) methods are required in order to solve the inference problem, but these methods are computationally intensive. It is thus necessary to apply computing techniques specifically tailored for accelerating the convergence and improving the efficiency of the sampling \\citep{feroz08} and/or to find smart approximations of the posterior near its local maxima \\citep{psnakesI}. But these enhancements have the cost of increasing dramatically the algorithmic complexity of the detection software, introducing new layers of intricacy in the form not only of additional assumptions and routines, but also of regularization 'constants', hidden variables, hyperparameters and selection thresholds that in many cases must be fine-tuned manually in order to be adapted to the specific circumstances of a given data set. The complexity of the algorithms can rise to almost baroque levels, having a negative effect on the portability of the codes and on the reproducibility of the results. We propose in this paper a simple strategy based on Bayesian methodology which incorporates sensible prior information about the source locations, the source fluxes and the source number distribution. With these priors and assuming a Gaussian likelihood, we can obtain an explicit form of the negative log-posterior of the number of sources and their fluxes and positions. Assuming a MAP methodology, we introduce a straightforward top-to-bottom detection algorithm that allows us to determine the number, fluxes and positions of the sources. We give a simple proof that the positions of the sources \\emph{must} be located in the local maxima of the matched-filtered image if there is not a significant overlap between sources. The main computational requirement of our algorithm is the solution of a system of non-linear equations. Our method differs from the one presented by \\cite{psnakesI} in five main points: \\begin{enumerate} \\item We use a more realistic set of priors for the source number, intensity and location distributions. In particular, our choice of the prior on the locations is also flat but depends on the number of sources $n$, which later proves to be decisive for the log-posterior. \\item We obtain an explicit form of the negative log-posterior and an explicit solution of the MAP estimate of the source intensities as the solution of a non-linear system of equations. \\item We prove that, for non-overlapping sources and a Gaussian likelihood, the MAP estimation of the positions of the sources is given by the location of the local maxima of the matched filtered images. \\item Since we are interested only in point sources, we fix the size parameter of the objects to be detected. \\item We can also find the MAP solution for the number of sources present in the images with a simple top-to-bottom search strategy. We do not need to resort to costly evaluations of the Bayesian evidence. \\end{enumerate} The layout of the paper is as follows: in section~\\ref{sec:methodology} we present the method and derive the corresponding posterior which includes the data likelihood and the priors. In section~\\ref{sec:simulations} we apply the method to CMB simulations with the characteristics of the radio Planck channels (from 30 to 100 GHz, where the number count priors are most reliable) and compare it with the standard procedure of using a MF with a $5\\sigma$ threshold. The main results are also presented in section~\\ref{sec:simulations}. The conclusions are given in the final section. ", "conclusions": "\\label{sec:conclusions} In this paper we propose a new strategy based on Bayesian methodology (BM), that can be applied to the blind detection of point sources in CMB maps. The method incorporates three prior distributions: a uniform distribution (\\ref{eq:prior_locations}) on the source locations, an extended power law on the source fluxes (\\ref{eq:normalized_gcauchy}) and a Poisson distribution on the number of point sources per patch (\\ref{eq:prior_number}). Together with a Gaussian likelihood, these priors produce the negative log-posterior (\\ref{eq:log_posterior}). We minimize this negative log-posterior with respect to the source fluxes in order to estimate them. At the same time, we show that the detected sources must be in the peaks of the matched-filtered maps. Finally, we choose the number of point sources which minimizes (\\ref{eq:log_posterior}) for the estimated fluxes. In this way, we give a non-arbitrary method to select the number of point sources. Finally, to check the performance of this technique, we carry out flat CMB simulations for the Planck channels from 30 to 100 GHz. For simplicity, we have excluded the foregrounds in our simulations, assuming that we are considering zones of the sky which have been cleaned by the application of component separation methods. However, we have included the confusion noise due to unresolved point sources in our simulations. \\begin{figure} \\includegraphics[width=\\columnwidth]{figflujos100GHz.eps} \\caption{\\label{fig8}{ Estimated flux against real flux for the BM (100 GHz). We have plotted the straight line $y=x$ for comparison. }} \\end{figure} We compare our Bayesian strategy with the application of a matched filter with a standard $5\\sigma$ threshold. We calculate the contamination, the completeness and the relative error for both methods. Though the percentage of spurious sources is a little higher for the BM at low fluxes $ \\simeq 0.2-0.3$ Jy, the completeness is much better, allowing us to obtain catalogues with a $99\\%$ completeness and no spurious sources from $0.7$ Jy (30 GHz), $0.8$ Jy (44 GHz), $0.55$ Jy (70 GHz) and $0.3$ Jy (100 GHz) on. The reconstruction errors in the estimated fluxes are similarly low for both methods. \\begin{figure} \\includegraphics[width=\\columnwidth]{fig9ref.eps} \\caption{\\label{fig9}{ Expectation value of the flux against estimated flux. The $95\\%$ confidence intervals are also plotted (100 GHz). }} \\end{figure}" }, "1101/1101.3179_arXiv.txt": { "abstract": "Introducing a surface layer of matter on the edge of a neutron star in slow rigid rotation, we analyze, from an intrinsic point of view, the junction conditions that must be satisfied between the interior and exterior solutions of the Einstein equations. In our model the core-\\textit{crust} transition pressure arise as an essential parameter in the description of a configuration. As an application of this formalism, we describe giant \\textit{glitches} of the Vela pulsar as a result of variations in the transition pressure, finding that these small changes are compatible with the expected temperature variations of the inner crust during \\textit{glitch} time ", "introduction": "Many pulsars show sudden spin jumps, \\textit{glitches}, superimposed to the gradual spin down due to the continued loss of angular momentum suffered by the star. The study of the properties of the star during the \\textit{glitch} time is essential to understand the structure of the neutron star that models the pulsar \\cite{Lorenz:1993,Link:1999}. In Vela pulsar, giant \\textit{glitches} with relative period variation of the order of $10^{-6}$ have been observed. Many mechanisms have been proposed to explain \\textit{glitches}. The models that deal with \\textit{glitches} are associated with the layer structure of neutron stars. Two regions of the star can be differentiated: the core and the \\textit{crust}. It is thought that the \\textit{crust} region has a solid crystalline structure similar to a metal \\cite{Haensel:2004nu}. Some theories propose that the \\textit{glitch} is triggered by the rupture of the \\textit{crust} as a consequence of the tensions on the \\textit{crust} that try to adequate the ellipticity of the \\textit{crust} to the changing angular velocity of the star. Because of the large density gradient close to the transition between the core and the \\textit{crust}, most of the crustal matter resides in the shells in the inner part of the \\textit{crust} \\cite{Pethick:1995}. Furthermore, it is noteworthy that the dynamical properties of the neutron star depend strongly on the transition pressure between the star's core and the \\textit{crust} as have been pointed up by \\cite{Lattimer:2001} and \\cite{Cheng:2002}. Because most of the matter of the \\textit{crust} is found in the shells near the transition region between the core of the star and the inner \\textit{crust}, this is the region where the properties of the \\textit{crust} are relevant \\cite{Pethick:1995}. Hence, we will treat the \\textit{crust} as a surface layer that envelops the star's core. In the next section we will describe how our model of slowly rotating neutron star with a surface layer \\textit{crust} is constructed. ", "conclusions": "" }, "1101/1101.3515_arXiv.txt": { "abstract": "The Magellanic Clouds (MCs) offer a unique opportunity to study the stellar evolution and nucleosynthesis of massive Asymptotic Giant Branch (AGB) stars in low metallicity environments where distances are known. Rubidium is a key element to distinguish between high mass AGB stars and low mass AGBs or other type of astronomical objects such as massive red supergiant stars. Theoretically, high mass AGBs are predicted to produce a lot of Rb. We present the discovery of massive Rb-rich AGB stars in the MCs, confirming for the first time that these stars also exist in other galaxies. Our findings show that these stars are generally brighter than the standard adopted luminosity limit (M$_{bol}$$\\sim$$-$7.1) for AGB stars. The observations of massive MC AGBs are qualitatively predicted by the present theoretical models. However, these theoretical models are far from matching the extremely high Rb overabundances observed. This might be related with an incomplete present understanding of the atmospheres of these stars. ", "introduction": "The Magellanic Clouds (hereafter MCs) provide a unique opportunity to study the evolution and nucleosynthesis of low- and intermediate-mass stars (0.8 $<$ M $<$ 8 M$_{\\odot}$) in low metallicity environments where distances - and hence luminosity - are known. Low- and intermediate-mass stars experience thermal pulses and strong mass loss on the Asymptotic Giant Branch \\citep[AGB; e.g.,][]{herw05}. Repeated thermal pulses and ``3$^{rd}$ dredge-up\" events can convert the originally O-rich AGB star into a C-rich one. However, in the case of the more massive AGB stars (M$>$4 M$_{\\odot}$), Hot Bottom Burning \\citep[HBB; e.g.,][]{mazz99} prevents the C-star formation and these stars remain O-rich despite the dredge-up. The activation of HBB in massive AGB stars is supported by previous studies on visually bright MC AGB stars \\citep[e.g.,][]{plez93} and on heavily obscured O-rich AGBs - the so-called ``OH/IR\" stars - of our Galaxy \\citep{gh06,gh07}. AGB stars also produce heavy neutron-rich elements ($s$-process elements) such as Rb, Zr, Sr, Nd, Ba, etc., which can be dredged-up to the stellar surface \\citep[e.g.,][]{buss99}. In the more massive AGB stars, free neutrons are predicted to be mainly released by the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction, while the $^{13}$C($\\alpha$,n)$^{16}$O reaction seems to be the dominant neutron source in lower mass AGB stars. Rb is a key element to distinguish between the operation of the $^{13}$C versus the $^{22}$Ne neutron source in AGB stars and, as such, is a good indicator of the progenitor stellar mass\\footnote{Note that other astronomical sources such as massive red supergiant stars are not expected to overproduce Rb.}. This is because the relative abundance of Rb to other $s$-process elements such as Zr (i.e., the Rb/Zr ratio) is sensitive to the neutron density owing to branchings in the s-process path at $^{85}$Kr and $^{86}$Rb \\citep[e.g.,][]{raai08b}. Interestingly, we discovered strong Rb overabundances (up to 10$-$100 times solar) with apparently only mild Zr enhancements in massive galactic O-rich AGB stars \\citep{gh06}. This work provided the first observational evidence that $^{22}$Ne is the dominant neutron source in the more massive AGB stars. Surprisingly, Rb was not found to be overabundant in the few unobscured O-rich massive AGBs previously studied in the SMC \\citep{plez93}. Here, we present the first detections of massive Rb-rich AGB stars in the MCs. ", "conclusions": "" }, "1101/1101.1510_arXiv.txt": { "abstract": "{} % {We seek to understand the way massive stars form. The case of a luminous YSO IRAS~17527-2439 is studied in the infrared.} {Imaging observations of IRAS~17527-2439 are obtained in the near-IR $JHK$ photometric bands and in a narrow-band filter centred at the wavelength of the H$_2$ 1-0S(1) line. The continuum-subtracted H$_2$ image is used to identify outflows. The data obtained in this study are used in conjunction with {\\it Spitzer}, AKARI, and IRAS data. The YSO driving the outflow is identified in the {{\\it Spitzer}} images. The spectral energy distribution (SED) of the YSO is studied using available radiative transfer models.} {A parsec-scale bipolar outflow is discovered in our H$_2$ line image, which is supported by the detection in the archival {{\\it Spitzer}} images. The H$_2$ image exhibits signs of precession of the main jet and shows tentative evidence for a second outflow. These suggest the possibility of a companion to the outflow source. There is a strong component of continuum emission in the direction of the outflow, which supports the idea that the outflow cavity provides a path for radiation to escape, thereby reducing the radiation pressure on the accreted matter. The bulk of the emission observed close to the outflow in the WFCAM and {{\\it Spitzer}} bands is rotated counter clockwise with respect to the outflow traced in H$_2$, which may be due to precession. A model fit to the SED of the central source tells us that the YSO has a mass of 12.23\\,M$_{\\odot}$ and that it is in an early stage of evolution.} {} ", "introduction": "Low- and intermediate-mass stars are known to form by gravitational collapse and subsequent accretion of their parent molecular clouds, and driving collimated outflows. However, the main mechanism leading to the formation of massive stars is debated as to whether it is either disk accretion similar to that for lower mass stars (e.g. Yorke \\& Sonnhalter \\cite{yorke02}) or a merger of lower-mass stars (e.g. Bonnell, Bate \\& Zinnecker \\cite{bonnell98}). Many of the recent CO line surveys show outflows from massive YSOs (e.g. Zhang et al. \\cite {zhang05}; Beuther et al. \\cite{beuther02}). Several massive YSO outflows have been observed in the near-IR, where the spatial resolution is better than that in single-dish CO line observations. A recent near-IR imaging survey by Varricatt et al. (\\cite{varricatt10}) shows that massive stars up to at least late-O spectral types form primarily by disk accretion. The case of a luminous YSO taking birth by accretion is presented in this paper. IRAS~17527-2439 (hereafter IRAS~17527) is a luminous YSO located in a dark cloud situated close to the Galactic plane ($l$ = 4.8273$^{\\circ}$, $b$ = 0.2297$^{\\circ}$) in the Ophiuchus region. It is associated with emission from dense gas and dust typical of massive YSOs. Molinari et al. (\\cite{molinari96}) detected NH$_3$ emission lines from this region. From the radial velocity of the NH$_3$ lines (V$_{LSR}$=13.3 km~s$^{-1}$), they estimated a kinematic distance of 3.23\\,kpc, with the distance ambiguity resolved. Based on the IRAS colours they classified IRAS~17527 as a ``high'' source, which is possibly in a UCH{\\sc{ii}} phase. (Note that at 12\\,$\\mu$m, the IRAS catalogue gives only an upper limit flux for this source). In their 97.981-GHz CS(2-1) survey of UCH{\\sc{ii}} candidates, Bronfman, Nyman \\& May (\\cite{bronfman96}) detected IRAS~17527 at V$_{LSR}$=13.5 km~s$^{-1}$, similar to the velocity at which Molinari et al. (\\cite{molinari96}) detected NH$_3$ emission. Massive star formation is often associated with H$_2$O, CH$_3$OH and OH maser emission. Palla et al. (\\cite{palla91}) detected H$_2$O maser emission from IRAS~17527. From the radial velocity of the maser (V$_{peak}$=-1.81~km~s$^{-1}$) they estimated a kinematic distance of 17.9\\,kpc, which is quite different from the distance estimated from radial velocity of the dense gas tracers NH$_3$ and CS. In massive YSOs, H$_2$O maser is often considered to be excited in jets (e.g. Felli, Palagi \\& Tofani \\cite{felli92}; Goddi et al. \\cite{goddi05}). The blueshift of the wavelength of the maser emission in this region with respect to the velocity of the dense gas tracers indicates that H$_2$O maser near IRAS~17527 also may be excited by a jet. Hence we adopt the distance estimate of 3.23\\,kpc (Molinari et al. \\cite{molinari96}) for the calculations in this paper. Surveys by van der Walt, Gaylard \\& MacLeod (\\cite{vanderwalt95}), Walsh et al. (\\cite{walsh97}) and Slysh et al. (\\cite{slysh99}) did not detect any 6.7\\,GHz Class-II methanol maser from this region. Edris, Fuller \\& Cohen (2007) detected faint (0.4\\,Jy) OH maser emission at V$_{peak}$=11.53\\,km~s$^{-1}$, located $\\sim$6$\\arcmin$ NW of IRAS~17527. Nevertheless, this offset is less than their beam size. Therefore it remains to be investigated at better spatial resolution. The faintness of the detected OH maser is consistent with their observation that the OH masers associated with younger sources are weaker than the ones associated with H{\\sc{ii}} or UCH{\\sc{ii}} regions. The VLA survey of Hughes \\& MacLeod (\\cite{hughes94}) detected 6-cm emission from IRAS~17527 with a peak flux of 5.3\\,mJy. The source was however rejected as a UCH{\\sc{ii}} candidate because the radio emission was diffuse. ", "conclusions": "\\begin{enumerate} \\item IRAS~17527 appears to be a luminous YSO taking birth through disk accretion. A well-collimated parsec-scale outflow is discovered in IRAS~17527 in H$_2$ line emission. \\item The outflow is seen to be bent in an `S'-shaped fashion, suggesting a precession of the jet. There is a tentative detection of a second outflow in the H$_2$ image. The possibility of more than one YSO in IRAS~17527 needs to be explored. \\item The H$_2$ image exhibits a strong continuum component in the emission along the outflow, which may be caused by radiation escaping through the outflow cavity. This, in turn, reduces the radiation pressure on the accreted matter and aids growth of the central source through accretion. \\item The bulk of the emission in the direction of the outflow, observed in $K$ and the {\\it Spitzer} bands, is rotated counter-clockwise with respect to the direction of the outflow traced by the H$_2$ line emission knots. It is probably a result of the precession of the jet. Spectroscopy in the the near-IR through {\\it Spitzer}-IRAC bands is required for a proper understanding of this. \\item The model fit to the SED shows that the central source is probably a Class-I protostar; this is supported by its location in the {\\it Spitzer}-IRAC colour-colour diagram. The YSO has a mass of $\\sim$12.23\\,M$_{\\odot}$ and a total luminosity of $\\sim$6.17$\\times$10$^3$\\,L$_{\\odot}$. \\item The disk parameters and the disk accretion ratio are poorly determined by the SED fitting. This may be caused by the contribution to the source magnitudes from emission from the outflow and from a possible companion. Observations at high angular resolution are required in the 2-20\\,$\\mu$m wavelength range for a more accurate determination of the source magnitudes at these wavelengths. \\end{enumerate}" }, "1101/1101.1504_arXiv.txt": { "abstract": "Probably No! As an example, using soft EOSs consistent with existing terrestrial nuclear laboratory experiments for hybrid neutron stars containing a quark core described with MIT bag model using reasonable parameters, we show that the recently discovered new holder of neutron star maximum mass PSR J1614-2230 of $1.97\\pm0.04M_{\\odot}$ can be well described by incorporating a Yukawa gravitational correction that is consistent with existing constraints from neutron-proton and neutron-lead scatterings as well as the spectroscopy of antiproton atoms. ", "introduction": "What is gravity? Are there additional spacetime dimensions? These are among the Eleven Science Questions for the New Century identified by the Committee on the Physics of the Universe, US National Research Council \\citep{11questions}. Interestingly, despite the fact that gravity is the first force discovered in nature, the quest to unify it with other fundamental forces remains elusive because of its apparent weakness at short-distance, see, e.g., refs. \\citep{Ark98,Pea01,Hoy03,Long03,Jean03,Boehm04a,Boe04,Dec05}. In developing grand unification theories, the conventional inverse-square-law (ISL) of Newtonian gravitational force has to be modified due to either the geometrical effect of the extra spacetime dimensions predicted by string theories and/or the exchange of weakly interacting bosons, such as the neutral spin-1 vector $U$-boson \\citep{Fayet}, proposed in the super-symmetric extension of the Standard Model, see, e.g., refs. \\citep{Adel03,Adel09,Fis99,New09,Uzan03,Rey05} for recent reviews. The modified gravity has also been proposed as an explanation for the present period of cosmological acceleration, see, e.g., ref.\\ \\citep{Ded08}. The search for evidence of modified gravity is at the forefront of research in several sub-fields of natural sciences including geophysics, nuclear and particle physics, as well as astrophysics and cosmology, see, e.g., refs.\\ \\citep{Fujii71,Pea01,Hoy03,Ark98,Long03,Adel09,Kap07,Nes08,Kam08,Aza08,Ger10,Luc10}. Various upper limits on the deviation from the ISL has been put forward down to femtometer range. Since the composition of neutron stars are determined mainly by the weak and electromagnetic forces through the $\\beta$ equilibrium and charge neutrality conditions while their stability is maintained by the balance of strong and gravitational forces, neutron stars are thus a natural testing ground of grand unification theories of fundamental forces. Moreover, neutron stars are among the densest objects with the strongest gravity in the Universe, making them ideal places to test strong-field predictions of General Relativity (GR) \\citep{Psa08}. The masses and radii of neutron stars are solely determined by both the strong-field behavior of gravity and the Equation of State (EOS) of dense stellar matter. However, there is no fundamental reason to choose Einstein's equations over other alternatives and it is known that the GR theory itself may break down at the limit of very strong gravitational fields, see, e.g., ref. \\citep{Psa08} for a comprehensive review. In fact, effects of modified gravity on properties of neutron stars have been under intense investigation. As expected, results of these studies are strongly model dependent, see, e.g., refs. \\citep{Ger01,Wis02,Aza08,Kri09,Wen0911}. Nevertheless, it is very interesting to note that alternative gravity theories that have all passed low-field tests but diverge from GR in the strong-field regime predict neutron stars with significantly different properties than their GR counterparts \\citep{Ded03}. Moreover, the deviations for neutron star properties from the GR predictions for these theories are larger than the uncertainty due to the poorly known EOS of dense matter in neutron stars. It was also clearly shown that the neutron star maximum mass alone can not distinguish gravity theories \\citep{Ded03}. Furthermore, in the endeavor of testing GR theory of gravity using properties of neutron stars, it is known that there is a degeneracy between the matter content and gravity. This degeneracy is tied to the fundamental Strong Equivalence Principle and can only be broken by using at least two independent observables \\citep{Yun10}. Recently, using the general relativistic Shapiro delay the mass of PSR J1614-2230 was precisely measured to be $1.97\\pm0.04M_{\\odot}$\\citep{Demo10}, making it the new holder of the maximum mass of neutron stars. Comparing with mass-radius relations predicted from solving the TOV equation using various EOSs within GR theory of gravity, it was shown that the mass of PSR J1614-2230 can rule out almost all soft EOSs especially those associated with hyperon or boson condensation. While conventional quark stars with soft EOSs are also ruled out by this observation, neutron stars with strongly interacting quark cores are allowed \\citep{Demo10,Ozel10,RXu10}. It was further shown that a transition to quark matter in neutron star cores can occur at densities comparable to the nuclear saturation density $\\rho_0$ only if the quarks are strongly interacting and are color superconducting \\citep{Ozel10}. The mass of PSR J1614-2230 was then used to constrain the interacting parameters of quarks. It was also shown that neutron stars with interacting quark clusters in their cores or solid quark stars can be very massive. Using the Lennard-Jones potential for interactions between quark clusters, the mass of the PSR J1614-2230 was used to constrain the number of quarks inside individual quark clusters \\citep{RXu10}. In this work, using soft nuclear EOSs for hybrid stars containing a quark core described by the MIT bag model with reasonable parameters, we show that the mass of PSR J1614-2230 is readily obtained by incorporating a Yukawa gravitational correction that is consistent with existing constraints from terrestrial nuclear laboratory experiments. ", "conclusions": "Among all fundamental forces, gravity remains the most uncertain one despite being the first discovered in nature. Neutron stars are natural testing grounds of grand unification theories of fundamental forces. In particular, they are ideal places to test GR predictions at the strong-field limit. Interpretations of observed properties of neutron stars require a comprehensive understanding of both gravity and the EOS of dense stellar matter. Before strong-field gravity is well understood, it is unlikely that the maximum mass of neutron stars alone can rule out any EOS. As an example, using soft nuclear EOSs consistent with existing terrestrial experiments for hybrid stars containing a quark core described by the MIT bag model with reasonable parameters, the maximum mass of PSR J1614-2230 is readily obtained by incorporating the Yukawa gravitational correction that is consistent with existing constraints from terrestrial nuclear laboratory experiments. We thank W.Z. Jiang, W. G. Newton, A.W. Steiner and Y. Zhang for useful discussions. D.H. Wen is supported in part by the National Natural Science Foundation of China under Grant No.10947023 and the Fundamental Research Funds for the Central University, China under Grant No.2009ZM0193. B.A. Li is supported in part by the US National Science Foundation under grant PHY-0757839, the National Aeronautics and Space Administration under grant NNX11AC41G issued through the Science Mission Directorate and the Texas Coordinating Board of Higher Education under grant No. 003565-0004-2007. L.W. Chen is supported in part by the National Natural Science Foundation of China under Grant Nos. 10675082 and 10975097, MOE of China under project NCET-05-0392, Shanghai Rising-Star Program under Grant No. 06QA14024, the SRF for ROCS, SEM of China, the National Basic Research Program of China (973 Program) under Contract No. 2007CB815004." }, "1101/1101.1732_arXiv.txt": { "abstract": "{% We examine the nature and role of mass loss via an equatorial decretion disk in massive stars with near-critical rotation induced by evolution of the stellar interior. In contrast to the usual stellar wind mass loss set by exterior driving from the stellar luminosity, such decretion-disk mass loss stems from the angular momentum loss needed to keep the star near and below critical rotation, given the interior evolution and decline in the star's moment of inertia. Because the specific angular momentum in a Keplerian disk increases with the square root of the radius, the decretion mass loss associated with a required level of angular momentum loss depends crucially on the outer radius for viscous coupling of the disk, and can be significantly less than the spherical, wind-like mass loss commonly assumed in evolutionary calculations. We discuss the physical processes that affect the outer disk radius, including thermal disk outflow, and ablation of the disk material via a line-driven wind induced by the star's radiation. We present parameterized scaling laws for taking account of decretion-disk mass loss in stellar evolution codes, including how these are affected by metallicity, or by presence within a close binary and/or a dense cluster. Effects similar to those discussed here should also be present in accretion disks during star formation, and may play an important role in shaping the distribution of rotation speeds on the ZAMS.} \\keywords {stars: mass-loss -- stars: evolution -- stars: rotation -- hydrodynamics} ", "introduction": "Classical models of stellar evolution focus on the dominant role of various stages of nuclear burning in the stellar core. But in recent years it has become clear that stellar evolution, particularly for more massive stars, can also be profoundly influenced by the loss of mass and angular momentum from the stellar envelope and surface \\citep{mm08}. In cool, low-mass stars like the sun, mass loss through thermal expansion of a coronal wind occurs at too-low a rate to have a direct effect on its mass evolution; nonetheless the moment arm provided by the coronal magnetic field means the associated wind angular momentum loss can substantially spin down the star's rotation as it ages through its multi-Gyr life on the main sequence. Except in close binary systems, the rotation speeds of cool, low-mass stars are thus found to decline with age, from up to $\\sim$100~km s$^{-1}$ near the ZAMS to a few km s$^{-1}$ for middle-age stars like the sun. By contrast, in hotter, more massive stars the role and nature of mass and angular momentum loss can be much more direct and profound, even over their much shorter, multi-Myr lifetimes. While some specific high-mass stars appear to have been spun down by strongly magnetized stellar winds (e.g HD 191612, \\citealt{donetal06}, or HD 37776, \\citealt{mikbra}), most massive stars are comparitively rapid rotators, with typical speeds more than 100~km s$^{-1}$, and in many stars, e.g.\\ the Be stars, even approaching the {\\em critical} rotation rate, at which the centrifugal acceleration at the equatorial surface balances Newtonian gravity \\citep{how04, toh, how07}. Indeed, models of the MS evolution of rotating massive stars show that, at the surface, the velocity approaches the critical velocity. This results from the transport of angular momentum from the contracting, faster rotating inner convective core to the expanding, slowed down radiative envelope \\citep{mee}. In stars with moderately rapid initial rotation, and with only moderate angular momentum loss from a stellar wind, this spinup from internal evolution can even bring the star to critical rotation \\citep{memb07}. Since any further increase in rotation rate is not dynamically allowed, the further contraction of the interior must then be balanced by a net loss of angular momentum through an induced mass loss. In previous evolutionary models, the required level of mass loss has typically been estimated by assuming its removal occurs from spherical shells at the stellar surface \\citep{mee}. This paper examines the physically more plausible scenario that such mass loss occurs through an {\\em equatorial, viscous decretion disk} \\citep{los91}. Such decretion disk models have been extensively applied to analyzing the rapidly (and possibly near-critically) rotating Be stars, which show characteristic Balmer emission thought to originate in geometrically thin, warm, gaseous disks in Keplerian orbit near the equatorial plane of the parent star \\citep{porriv, carbjo}. But until now there hasn't been much consideration of the role such viscous decretion disks might play in the rotational and mass loss {\\em evolution} of massive stars in general. As detailed below, a key point of the analysis here is that, per unit mass, the angular momentum loss from such a decretion disk can greatly exceed that from a stellar wind outflow. Whereas the angular momentum loss of a nonmagnetized wind is fixed around the transonic point very near the stellar surface, the viscous coupling in a decretion disk can transport angular momentum outward to some outer disk radius $R_\\text{out}$, where the specific angular momentum is a factor $\\sqrt{R_\\text{out}/R_\\text{eq}}$ higher than at the equatorial surface. For disks with an extended outer radius $R_\\text{out} \\gg R_\\text{eq}$, the angular momentum loss required by the interior evolution can then be achieved with a much lower net mass loss than in the wind-like, spherical ejection assumed in previous evolution models. For a given angular momentum shedding mandated by interior evolution, quantifying the associated disk mass loss thus requires determining the disk outer radius. For example, in binary systems, this would likely be limited by tidal interactions with the companion, and so scale with the binary separation \\citep{ok02}. But in single stars, the processes limiting this outer radius are less apparent. Here we explore two specific mechanisms that can limit the angular momentum loss and/or outer radius of a disk, namely thermal expansion into supersonic flow at some outer radius, and radiative ablation of the inner disk from the bright central star. For each case, we derive simple scaling rules for the required disk mass loss as a function of assumed stellar and wind parameters, given the level of interior-mandated angular momentum loss. The organization for the remainder of this paper is as follows: Sect.~\\ref{secana} presents simple analytical relations for how the presence of a disk affects the mass loss at the critical limit. Sect.~\\ref{secnu} develops set of equations governing structure and kinematics of the disk, while Sect.~\\ref{kaprenum} solves these to derive simple scaling for how thermal expansion affects the outer disk radius and disk mass loss. Sect.~\\ref{secabla} discusses the effects of inner-disk ablation by a line-driven disk wind induced from the illumination of an optically thick disk by the centeral star, deriving the associated ablated mass loss and its effect on the net disk angular momentum and mass loss. Sect.~\\ref{secevol} gives a synthesis of the different cases discussed here and offers a specific recipe for incorporating disk mass loss rates into stellar evolution codes. Sect.~\\ref{secot} discusses some complementary points (e.g.\\ viscous decoupling, tidal effects of nearby stars, reduced metallicity, etc.), while Sect.~\\ref{seccon} concludes with a brief summary of the main results obtained in this work. ", "conclusions": "\\label{seccon} We examine the mechanism of the mass and angular momentum loss via decretion disks associated with near-critical rotation. The disk mass loss is set by the angular momentum needed to keep the stellar rotation at or below the critical rate. We study the potentially important role of viscous coupling in outward angular momentum transport in the decretion disk, emphasizing that the specific angular momentum at the outer edge of the disk can be much larger than at the stellar surface. For a given stellar interior angular momentum excess, the mass loss required from a decretion disk can be significantly less than invoked in previous models assuming a direct, near-surface release. The efficiency of the angular momentum loss via disk depends on the radius at which the viscous coupling ceases the transport the angular momentum to the outflowing material. When the radiative force is negligible, we argue that this likely happens close to the disk sonic (critical) point setting the most efficient angular momentum loss. In the opposite case, when the radiative force is nonnegligible, there is not a single point beyond which the viscous coupling disappears. The disk is continuously ablated below the sonic point, and the ablated material ceases to be viscously coupled, decreasing the efficiency of angular momentum loss. We describe the method to include these processes into evolutionary calculations. The procedure provided enables to calculate the mass-loss rate necessary for a required angular momentum loss just from the stellar and line force parameters. We can distinguish three different physical circumstances: \\addtolength\\leftmargini{8mm} \\begin{itemize} \\item[case A:] When the disk wind is able to remove the whole excess of angular momentum (the disk is completely ablated by the wind, see Eq.~\\eqref{defab}) then the outer disk radius is given by Eq.~\\eqref{defab}, and the required mass loss is given by Eq.~\\eqref{dmdtjed}. The limiting case $R_\\text{out}\\approx R_\\text{eq}$ would then correspond to the near surface release of the matter without any disk. Note that in a rare case when the analysis leads to $R_\\text{out}>R_\\text{crit}$ the radius $R_\\text{crit}$ should be used as the outer disk radius (case B). The expressions presented in the paper are given in the hypothesis of an optically thick disk and should be appropriately modified for optically thin disks. \\item[case B:] If the raditiave force is not able to remove sufficient angular momentum (the disk is not completely ablated) then part of the excess angular momentum must be carried away by the disk (Eq.~\\eqref{partj}). In this case Eqs.~\\eqref{partj}, \\eqref{partm} can be used to estimate the mass-loss rate. The outer disk edge could be identified with the critical point. \\item[case C:] If the effects of the radiative force are negligible, then the whole excess of angular momentum is carried away by the disk and the the outer disk edge is approximately given by $R_\\text{crit}$ and the required mass-loss rate could be derived from Eq.~\\eqref{mamol}. \\end{itemize} \\addtolength\\leftmargini{-8mm} Finally, we note that, in absence of strong magnetic field, many of the features discussed here may also be applicable to the case of star-formation accretion disks." }, "1101/1101.4098_arXiv.txt": { "abstract": "A new method to constrain the distance of blazars with unknown redshift using combined observations in the GeV and TeV regimes will be presented. The underlying assumption is that the Very High Energy (VHE) spectrum corrected for the absorption of TeV photons by the Extragalactic Background Light (EBL) via photon-photon interaction should still be softer than the extrapolation of the gamma-ray spectrum observed by Fermi/LAT. Starting from the observed spectral data at VHE, the EBL-corrected spectra are derived as a function of the redshift z and fitted with power laws. Comparing the redshift dependent VHE slopes with the power law fits to the LAT data an upper limit to the source redshift can be derived. The method is applied to all TeV blazars detected by LAT with known distance and an empirical law describing the relation between the upper limits and the true redshifts is derived. This law can be used to estimate the distance of unknown redshift blazars: as an example, the distance of PKS 1424+240 is inferred. ", "introduction": "\\begin{center} \\begin{table} {\\small \\centering \\begin{tabular}{llcccccc} \\hline Source Name & $z$[real] &{\\it Fermi}/LAT slope & TeV slope & $z^*$ \\\\%^ & $z$[rec] \\\\ \\hline Mkn 421 & 0.030 & 1.78 $\\pm$ 0.03 & 2.3 $\\pm$ 0.1 & 0.08 $\\pm$ 0.02 \\\\ Mkn 501 & 0.034 & 1.73 $\\pm$ 0.06 & 2.3 $\\pm$ 0.1 & 0.10 $\\pm$ 0.02 \\\\ 1ES 2344$+$514 & 0.044 & 1.76 $\\pm$ 0.27 & 2.9 $\\pm$ 0.1 & 0.20 $\\pm$ 0.06\\\\ Mkn 180 & 0.045 & 1.91 $\\pm$ 0.18 & 3.3 $\\pm$ 0.7 & 0.20 $\\pm$ 0.12 \\\\ 1ES 1959$+$650 & 0.047 & 1.99 $\\pm$ 0.09 & 2.6 $\\pm$ 0.2 & 0.09 $\\pm$ 0.04 \\\\ BL Lacertae & 0.069 & 2.43 $\\pm$ 0.10 & 3.6 $\\pm$ 0.5 & 0.23 $\\pm$ 0.12 \\\\ PKS 2005$-$489 & 0.071 & 1.91 $\\pm$ 0.09 & 3.2 $\\pm$ 0.2 & 0.19 $\\pm$ 0.04 \\\\ W Comae & 0.102 & 2.02 $\\pm$ 0.06 & 3.7 $\\pm$ 0.2 & 0.23 $\\pm$ 0.05 \\\\ PKS 2155$-$304 & 0.116 & 1.87 $\\pm$ 0.03 & 3.4 $\\pm$ 0.1 & 0.22 $\\pm$ 0.01 \\\\ 1ES 0806$+$524 & 0.138 & 2.04 $\\pm$ 0.14 & 3.6 $\\pm$ 1.0 & 0.23 $\\pm$ 0.15 \\\\ 1ES 1218$+$304 & 0.182 & 1.63 $\\pm$ 0.12 & 3.1 $\\pm$ 0.3 & 0.21 $\\pm$ 0.08 \\\\ 1ES 1011$+$496 & 0.212 & 1.82 $\\pm$ 0.05 & 4.0 $\\pm$ 0.5 & 0.49 $\\pm$ 0.12 \\\\ S5 0716$+$714 & 0.310$^{a,b}$ & 2.16 $\\pm$ 0.04 & 3.4 $\\pm$ 0.5 & 0.21 $\\pm$ 0.09 \\\\ PG 1553+113 & 0.400$^c$ & 1.69 $\\pm$ 0.04 & 4.1 $\\pm$ 0.2 & 0.57 $\\pm$ 0.05 \\\\ 3C66A & 0.444$^a$ & 1.93 $\\pm$ 0.04 & 4.1 $\\pm$ 0.4 & 0.34 $\\pm$ 0.05 \\\\ 3C279 & 0.536 & 2.34 $\\pm$ 0.03 & 4.1 $\\pm$ 0.7 & 0.75 $\\pm$ 0.72\\\\ \\hline \\end{tabular} \\caption {TeV blazars used in this study. The sources used in this study are listed in the first column, their redshift (second column), their {\\it Fermi}/LAT slope (third column), the VHE slope of the observed differential energy spectrum fit (fourth column) and the value $z^*$ (last column). $^a$Uncertain; $^b$from \\cite{nilsson08}; $^c$from \\cite{danforth10}. Detailed references can be found in \\cite{prandini10}.}\\label{table_values} } \\end{table} \\end{center} The extragalactic TeV sky catalogue ($E>100$ GeV), counts nowadays 45 objects\\footnote{For an updated list see: http://www.mppmu.mpg.de/$\\sim$rwagner/sources/}. Many of these sources have been recently detected also at GeV energies by the {\\it Fermi} satellite~\\cite{abdo09}, allowing for the first time a quasi-continuous coverage of the spectral shape of extragalactic VHE emitters over more than 4 decades of energy. The large majority of extragalactic TeV emitting objects are blazars, radio-loud active galactic nuclei with a relativistic jet closely oriented toward the Earth, as described in \\cite{urry}. Here, we discuss a method, recently published in \\cite{prandini10}, to derive an upper limit on the redshift of a blazar, based on the comparison between the spectral index at GeV energies as measured by LAT (unaffected by the cosmological absorption up to redshifts far beyond those of interest here) and the TeV spectrum corrected for the absorption. Starting from the derived limits, we find a simple law relating these values to real redshift, which can be used to guess the distance of unknown redshift blazars. We assume a cosmological scenario with $h=0.72$, $\\Omega_M=0.3$ and $\\Omega_\\Lambda=0.7$. ", "conclusions": "" }, "1101/1101.0358_arXiv.txt": { "abstract": "We report the discovery ($20\\sigma$) of kilohertz quasi-periodic oscillations (kHz QPOs) at $\\sim 690$ Hz from the transient neutron star low-mass X-ray binary EXO 1745--248. We find that this is a lower kHz QPO, and systematically study the time variation of its properties using smaller data segments with and without the shift-and-add technique. The quality (Q) factor occasionally significantly varies within short ranges of frequency and time. A high Q-factor ($264.5\\pm38.5$) of the QPO is found for a 200 s time segment, which might be the largest value reported in the literature. We argue that an effective way to rule out kHz QPO models is to observationally find such high Q-factors, even for a short duration, as many models cannot explain a high coherence. However, as we demonstrate, the shift-and-add technique cannot find a very high Q-factor which appears for a short period of time. This shows that the coherences of kHz QPOs can be higher than the already high values reported using this technique, implying further constraints on models. We also discuss the energy dependence of fractional rms amplitude and Q-factor of the kHz QPO. ", "introduction": "Many neutron star low mass X-ray binary (LMXB) systems show high frequency ($\\sim 200-1200$ Hz) and somewhat coherent intensity variations \\citep{vanderKlis2006, Bhattacharyya2010}. Such variations are known as kilohertz quasi-periodic oscillations (kHz QPOs). Sometimes these QPOs appear in a pair; the higher frequency one is called the upper kHz QPO, while the lower frequency one is known as the lower kHz QPO. Soon after their discovery \\citep{vanderKlisetal1996, Strohmayeretal1996}, it was realized that they originate from within a few Schwarzschild radii of the neutron star, and hence could be useful to probe the strong gravity regime, as well as the supranuclear degenerate matter of the stellar core. Despite this potential, so far kHz QPOs could not be used as a reliable tool, because their correct theoretical model has not been identified yet. Many proposed models of this timing feature primarily attempt to explain the frequency (e.g., \\citet{Milleretal1998, StellaVietri1998, StellaVietri1999, LambMiller2003, KluzniakAbramowicz2001, AbramowiczKluzniak2001, Wijnandsetal2003, Leeetal2004, Zhang2004, Mukhopadhyay2009}). Some of these models involve various general relativistic frequencies at preferred radii and the neutron star spin frequency. Such models can have good predictive power \\citep{Bhattacharyya2010}. However, it is not enough to explain only the frequency in order to understand the kHz QPOs. The modulation mechanism (how the intensity actually varies) and the decoherence mechanism (why the QPOs are not very narrow) are also required to be understood. \\citet{Mendez2006} proposed that, although the kHz QPO frequencies are plausibly determined by the characteristic disk frequencies, the modulation mechanism is likely associated to the high energy spectral component (e.g., accretion disk corona, boundary layer, etc.). This is because the disk alone cannot explain the large observed amplitudes, especially at hard X-rays where the contribution of the disk is small. Moreover, the kHz QPO fractional amplitude increases with energy, and plausibly reaches a saturation value (e.g., \\citet{Gilfanovetal2003}). Therefore, an accurate measurement of the amplitude vs. energy curve might be useful to understand the modulation mechanism. The lack of coherence (aperiodicity) of kHz QPOs could be intrinsic (e.g., shot-noise models; \\citet{vanderKlis2006} and references therein), or might be because of a decoherence mechanism or frequency drift. The decoherence could be because of a damped harmonic oscillator or a finite lifetime of a clump of matter. A timing feature originating from a finite-width disk annulus could also be broadened because of the superposition of a range of frequencies. A frequency drift, for example due to a varied preferred disk radius, might also broaden the feature. Elimination of the effects of frequency superposition and drift is required to measure the coherence of the underlying signal, which is essential to understand the kHz QPOs. The shift-and-add technique of \\citet{Mendezetal1998}, which was originally used to discover an upper kHz QPO from 4U 1608--52, has been used by several authors \\citep{Barretetal2005a, Barretetal2005b, Barretetal2006, Barretetal2010} to track the frequency drift. Subsequent alignment of kHz QPOs in smaller time segments make the signal much more prominent, and allows one to measure its quality (Q) factor or coherence. The Q-factor is the ratio of the QPO centroid frequency to the full-width half maximum (FWHM) of the QPO profile. \\citet{Barretetal2005a, Barretetal2005b, Barretetal2006} noticed that, while the Q-factor of lower kHz QPOs can be up to $\\sim 200$, that of upper kHz QPOs is always less than 50. These authors have also found that the Q-factor of the lower kHz QPO first increases with frequency, and then after a frequency characteristic of a source, the Q-factor decreases. They tentatively interpreted this as a signature of the innermost stable circular orbit (ISCO; but see \\citet{Mendez2006}), whose existence is a key prediction of strong field general relativity. Note, however, although the shift-and-add technique tracks the frequency drift correctly, it gives an average Q-factor value over a long period. Consequently, this technique would miss a plausible very high Q-factor value, if such a value appears for a short period of time. In this Letter, we report the discovery of a very strong lower kHz QPO from the neutron star LMXB EXO 1745-248, and a very high coherence of this QPO for a short duration. ", "conclusions": "\\label{Discussion} In this Letter, we report the first detection and analysis of a kHz QPO from the neutron star LMXB EXO 1745--248. In order to investigate if its apparent large width and structure are caused by the frequency drift, we have divided the data segment into smaller parts (see \\S~\\ref{DataAnalysisandResults}). Note that initially we have not used the shift-and-add technique, so that the Q-factor of the underlying signal does not get averaged, and we can study its time variation within a short duration. The measured $Q > 50$ values for all the smaller parts except one (Fig.~\\ref{fig_Qfactor-final.ps}) show that this is a lower kHz QPO (\\S~\\ref{Introduction}). We have not found an upper kHz QPO even by the shift-and-add technique (\\S~\\ref{DataAnalysisandResults}). The frequency of the QPO roughly monotonically drifts by $\\sim 30$ Hz in about an hour (Fig.~\\ref{HF-Powspec2}). The variations of fractional rms amplitude ($8-10 \\%$) and Q-factor, however, are not correlated with the frequency change in our limited time and frequency ranges. The fact that the Q-factor is generally larger for smaller time segments, and the frequency shifts from one segment to another shows that the frequency drift plays a major role in broadening the QPO. At the level of one-sixteenth segment (200 s; \\S~\\ref{DataAnalysisandResults}), the width of the QPO may still be caused by the combined effect of a small frequency drift, the superposition of a small frequency range, and a decoherence mechanism. Since we cannot divide the time segment any further (see \\S~\\ref{DataAnalysisandResults}), our estimated Q-factor is only a lower limit. The corresponding lower limit of the coherence time $\\tau$ can be calculated from $\\tau = 1/(\\pi\\times{\\rm FWHM})$, where the QPO profile is fitted with a Lorentzian, and we assume that the signal consists of exponentially damped sinusoidal oscillator \\citep{Barretetal2005a, vanderKlis2006}. For a one-sixteenth segment we have measured a Q-factor of $264.5 \\pm 38.5$, which might be the largest value reported in the literature (\\citet{Barretetal2005b} reported $Q = 222\\pm24$ for 4U 1636--536). Moreover, this value appears at $\\approx 694$ Hz. From \\citet{Barretetal2006} we find that Q-factor near this frequency is found to be less than 150. The minimum $\\tau$ corresponding to a Q-factor of 264.5 is 0.12 s. Many models will find it difficult to explain such a large $\\tau$ (see \\citet{Barretetal2005a} for a discussion). For example, kHz QPO models based on accretion flow inhomogeneities in the form of clumps should have $\\tau \\lsim 0.01$ s \\citep{Barretetal2005a}. Therefore, an effective way to constrain and rule out kHz QPO models is to observationally find high Q-factors, even for a short duration. If the Q-factor considerably varies within short ranges of time and frequency (see Fig.~\\ref{fig_Qfactor-final.ps}), the shift-and-add technique (\\S~\\ref{Introduction}), which gives an average Q-factor and QPO structure, cannot find such high Q-factors. We have demonstrated this in Fig.~\\ref{shiftandaddvs6of16_kHzQPO-final-coma.ps} (\\S~\\ref{DataAnalysisandResults}). Therefore, while the Q-factor vs. frequency trend reported by Barret and coauthors should be useful to probe fundamental physics, the Q-factor of lower kHz QPOs can be much larger at times than the reported peak values of these authors. Finally, the observed increase of fractional rms amplitude with energy is common for kHz QPOs (e.g., \\citet{Gilfanovetal2003}), and may provide crucial information about the hard X-ray modulation mechanism by identifying a modulation site (see \\S~\\ref{Introduction}). Moreover, a plausible increase of Q-factor with energy indicates that oscillations from harder X-ray components are more coherent than those from softer X-ray components. {\\it Astrosat} with its sufficient time resolution and an area larger than {\\it RXTE} PCA at hard X-rays might be ideal to probe these trends." }, "1101/1101.0996_arXiv.txt": { "abstract": "In an extreme mass-ratio binary black hole system, a non-equatorial orbit will list (i.e. increase its angle of inclination, ${\\iota}$) as it evolves in Kerr spacetime. \\ The abutment, a set of evolving, near-polar, retrograde orbits, for which the instantaneous Carter constant (${Q)}$ is at its maximum value (${Q}_{X}$) for given values of latus rectum (${\\tilde{l}}$) and eccentricity (${e}$), has been introduced as a laboratory in which the consistency of $dQ/dt$ with corresponding evolution equations for $d\\tilde{l}/dt$ and $de/dt$ might be tested independently of a specific radiation back-reaction model. To demonstrate the use of the abutment as such a laboratory, a derivation of $dQ/dt$, based only on published formulae for $d\\tilde{l}/dt$ and $de/dt$, was performed for elliptical orbits on the abutment. \\ The resulting expression for $dQ/dt$ matched the published result to the second order in $e$. \\ We believe the abutment is a potentially useful tool for improving the accuracy of evolution equations to higher orders of $e$ and $\\tilde{l}^{-1}$. ", "introduction": "An extreme mass-ratio binary black hole system (EMRI) is composed of a primary object, which can be a Kerr black hole of mass $M\\sim10^{6}-10^{7}$ solar masses with a spin \\footnote{Our use of ${\\tilde{S}}$ for black hole spin arose during our initial studies of the work of Barack and Cutler \\cite{Barack:2004uq}.} ${\\tilde{S}}=\\left\\vert {{\\mathbf{J}}}\\right\\vert /{M}^{2}$ (where ${{\\mathbf{J}}}$ is the spin angular momentum), and an orbiting secondary object of mass $m\\sim1-10$ solar masses. Theoretical models to describe the orbital evolution of the secondary object in various situations have been derived and presented in the literature: circular orbits in the equatorial plane of the primary object \\cite {1973blho.conf..215B,1993PhRvD..47.1497P,1993PhRvD..47.1511C,Ryan:1995ix,2000PhRvD..62l4022O,2011arXiv1105.2959P}% , elliptical orbits in the equatorial plane \\cite {PhysRev.131.435,PhysRev.136.B1224,1991STIA...9240399B,1992MNRAS.254..146J,PhysRevD.50.3816,2002PhRvD..66d4002G,2005PhRvL..94v1101H}% , and an extensive body of research on circular or elliptical orbits inclined with respect to the equatorial plane \\cite {1973ApJ...185..635T,1973ApJ...185..649P,1995PhLA..202..347O,Ryan:1996ly,1996PhRvD..53.4319K,PhysRevD.55.3444,1999CQGra..16.2929D,PhysRevD.61.084004,2001PhRvD..64f4004H,2002PhRvD..66f4005G,Barack:2004uq,2005PThPh.114..509S,2006PThPh.115..873S,2006PhRvD..73f4037G,2007PhRvD..76d4007B,2007PThPh.117.1041G,Flanigan:2007kx,2007PhRvD..75b4005B}% . Such models are used to generate hypothetical gravitational waveforms (GW), which provide templates for use in the detection of gravitation wave signals by pattern recognition (Punturo et al. \\cite{Punturo:2010zza}). The detection of GW radiation by the Earth-based Laser Interferometer Gravitational Wave Observatory\\ (LIGO) or the Laser Interferometer Space Antenna (LISA) depends fundamentally on the availability of correct templates \\cite{2005PhRvL..94v1101H,Cutler:2007mi,Normandin:2008:GWS:1460936.1461188}. Performing direct observations of relativistic effects is an important challenge. The Solar System affords one the opportunity to observe and model the motions of natural and artificial bodies in Kerr spacetime in the weak-field, slow motion limit \\cite{Iorio2011AstroSpaceSci,EverittPRL2011}; and recent measurements of artificial-satellite orbits have produced estimates of the Lense-Thirring precession to an accuracy of 10\\% \\cite{Iorio2011AstroSpaceSci}. Further, the discovery of Sagittarius A*, a massive black hole (MBH) of $\\sim4.0\\times10^{6}$ solar masses, at the centre of our galaxy (see \\cite{IorioPRD2011,SadeghianCQG2011,IorioMNRAS2011} and references therein), offers a new opportunity to study Kerr spacetime by the observation of various stars in inclined, highly elliptical orbits, and by the analysis of their orbital dynamics \\cite{IorioPRD2011,SadeghianCQG2011,IorioMNRAS2011,WillAJL2008}. Relativistic effects are difficult to discern since the orbital periods of the stars are in the tens of years \\cite{WillAJL2008}, and for orbits that come close to the MBH, tidal disruption is a concern (\\cite{WillAJL2008} and see Appendix B in \\cite{SadeghianCQG2011}); yet, observation has great potential to aid in the study of Kerr spacetime. In the case of an EMRI, unfortunately, the part played by a theoretician is a fiduciary one; thus, the introduction of tools with which the evolution equations can be tested for consistency is most beneficial: the abutment is one such tool, but it is not intended to replace existing methods. The concept of the abutment, a boundary that defines a set of near-polar retrograde orbits, was developed and introduced by P. G. Komorowski in his Doctoral thesis \\cite{Komorowski:2011fk} and in a previous work \\cite{Komorowski:2010we} (we shall review the abutment in detail in section \\ref{sub:The-abutment.}); two uses of the abutment had emerged: first, it suggested a means of testing the consistency of the evolution of the Carter constant of circular orbits ($d{Q}/d{t}$) with respect to that of the latus rectum ($d{\\tilde{l}}/dt$); and second, it permitted a numerical analysis of the rate of change of the orbital angle of inclination, $\\iota$, with respect to ${\\tilde{l}}$ ($\\left( \\partial{\\iota}/\\partial{\\tilde{l}}\\right) _{\\min }$) for circular orbits constrained to evolve along the abutment. In this work we shall extend these uses to orbits of non-zero eccentricity ($0\\leq e\\leq1$) by testing the consistency of expressions for $d{Q}/d{t}$ with expressions for $d{\\tilde{l}}/dt$ and $d{e}/dt$, and we shall perform an analytical treatment of $\\iota$ and the list rate of the same. Further, a physically realistic orbital evolution follows the abutment (${Q}_{X}$) in only one case, the evolution of an orbit in a Schwarzschild black hole (SBH) system (${\\tilde{S}% }=0$). We shall now consider the general case of an evolving orbit that intersects the abutment, ${Q}_{X}$, tangentially at a single point (contact of the first order (see 99 in \\cite{J.W.:1955fk})) as it follows a path defined by ${Q}_{path}$. Further, by performing our analysis for elliptical orbits, the abutment becomes a two dimensional surface that defines the maximum value of ${Q}$ for given values of $e$ and latus rectum, ${\\tilde{l}=l/M}$. Therefore one must view the abutment as a set of contiguous points rather than a path to be followed by an evolving orbit; and it is at these points that the derivatives, $\\partial{Q}_{X}/\\partial{\\tilde{l}}$ and $\\partial{Q}% _{X}/\\partial{e}$, fix the corresponding slopes of ${Q}_{path}$. But as reported in \\cite{Komorowski:2010we}, the second-order effect\\footnote{When we refer to the second-order effect at the abutment, we refer to the second derivative of $Q_{X}$ with respect to $\\tilde{l}$, not to $\\tilde{S}^{2}$. \\ See \\cite{Komorowski:2010we} for background discussion.} must be included when working with $\\iota$ at the abutment. In section \\ref{sec:An-Analytical-Formula} we shall analytically derive the formula for ${\\iota}$ for elliptical orbits\\ on the abutment, and thus confirm the result for $\\left( \\partial{\\iota}/\\partial{\\tilde{l}}\\right) _{\\min}$ \\cite{Komorowski:2010we}, which was derived numerically for circular orbits. In addition, we shall analytically derive $\\partial{\\iota}/\\partial{e}$ for elliptical orbits that evolve on the abutment. In section \\ref{sec:Correction-of-didl-for} we shall include the effect of the second derivative of ${Q}_{path}$ (i.e. the second-order effect) by introducing reductive ans\\\"{a}tze for circular and elliptical orbits, and thus create a more physically realistic model for an evolving orbit at the abutment. Because our abutment model is independent of any specific radiation back-reaction model, we now have a laboratory that allows us to perform tests of established listing formulae. In section \\ref{Treatment-of-Qdot-and-idot-at-the-abutment}, we shall demonstrate the usefulness of the abutment in testing the consistency of $d{Q}/d{t}$ equations with respect to $d\\tilde{l}/dt$ and $de/dt$ evolution equations, and\\ in calculating $d\\iota/dt$ for elliptical orbits of small eccentricity (i.e. near-circular). In section \\ref{Conclusions} we shall conclude our work and recommend directions that warrant further study. We define $\\iota$ to be the maximum polar angle reached by the secondary object in its orbit (see equation (42) in \\cite{Komorowski:2010we}). This definition differs from that used by others (Gair and Glampedakis \\cite{2006PhRvD..73f4037G} and Glampedakis, Hughes, and Kennefick \\cite{2002PhRvD..66f4005G}); but when performing our analysis to the leading order in $\\tilde{S}$, there is no significant difference. ", "conclusions": "Conclusions} For inclined test-particle orbits around a black hole, two solutions for ${X}^{2}$ (where $X=\\tilde{L}_{z}-\\tilde{S}\\tilde{E}$) can be derived: ${X}_{-}^{2}$ and ${X}_{+}^{2}$. Given a Schwarzschild black hole (SBH), ${X}_{+}^{2}% ={X}_{-}^{2}$ on any polar orbit, where ${X}_{-}^{2}$ corresponds to prograde orbits and ${X}_{+}^{2}$ corresponds to retrograde orbits. For a Kerr black hole (KBH) the orbits on which ${X}_{+}^{2}={X}_{-}^{2}$ are not polar, but near-polar and retrograde. Such orbits comprise the abutment at which the value of the Carter constant ($Q$) is a maximum for given values of latus rectum ($\\tilde{l}$) and eccentricity ($e$). In this work we derived an analytical formula for the value of orbital inclination, $\\iota$, of an elliptical orbit on the abutment. By performing the partial differentiation of $\\iota$ with respect to $\\tilde{l}$, we were able to confirm the numerical result for $\\partial\\iota/\\partial\\tilde{l}$ reported in Komorowski et al. \\cite{Komorowski:2010we} \\ for circular orbits, and we were able to extend the formula to include $\\partial\\iota /\\partial\\tilde{l}$ for elliptical orbits. A result for $\\partial \\iota/\\partial e$ was also obtained for elliptical orbits. Further, it allowed one to redefine, in terms of $e$, $\\tilde{l}$, and $\\tilde{S}$, any trigonometric function that might be found in an evolution equation to be tested at the abutment. Evolving orbits in Kerr spacetime are not constrained to follow the abutment. Instead, the value of $Q$ will follow $Q_{path}$, which intersects the abutment tangentially at an arbitrary point of contact of the first order. This behaviour is assured because the value of $Q_{path}$ cannot exceed that of $Q$ on the abutment; to do so would make ${X}_{\\pm}^{2}$ complex and thus unphysical. For circular orbits, we modelled the second-order behaviour reported in \\cite{Komorowski:2010we} by introducing a bounded function $f\\left( e,\\tilde{l}\\right) $ (also in terms of $\\tilde{S}$) to reduce the value of $\\partial^{2}Q_{path}/\\partial\\tilde{l}^{2}$ while leaving $Q_{path}$ and $\\partial Q_{path}/\\partial\\tilde{l}$ equal to their corresponding values ($Q_{X}$ and $\\partial Q_{X}/\\partial\\tilde{l}$) on the abutment. This approach was then applied to elliptical orbits, and a new bounded function $g\\left( e,\\tilde{l}\\right) $, which depends upon $de/d\\tilde{l}$, was used to reduce the value of $\\partial^{2}Q_{X}/\\partial\\tilde{l}^{2}$ to $\\partial^{2}Q_{path}/\\partial\\tilde{l}^{2}$. It was discovered that the value of $\\partial^{2}Q_{path}/\\partial e^{2}$ remained unchanged by the reductive ansatz elliptical. The consistency of published evolution equations, $dQ/dt$, $d\\tilde{l}/dt$, and $de/dt$, was tested by using $d\\tilde{l}/dt$ and $de/dt$ to generate an expression for $dQ/dt$ at the abutment. In general, the calculation of $dQ/dt$ is more difficult to perform than that of $d\\tilde{l}/dt$ and $de/dt$ \\cite{PhysRevD.55.3444}; hence, the abutment provides a useful mechanism for testing the validity of radiation back-reaction models. Indeed, the evolution equations reported by Ganz et al. \\cite{2007PThPh.117.1041G} were confirmed to their 2.5PN order. The abutment provides a consistency condition that is limited to near-polar retrograde orbits; and yet, some back-reaction models have been found to exhibit pathological behaviour for polar orbits (Gair and Glampedakis \\cite{2006PhRvD..73f4037G}). Another consistency condition is already known for orbits of any $\\iota$ (see equations (13) and (14) in \\cite{2006PhRvD..73f4037G}), but it applies to circular orbits (i.e. in the limit $e\\rightarrow0$); the abutment is valid for orbits of arbitrary eccentricity, depending on the accuracy of the back-reaction model used. This method promises to be a useful tool for confirming the accuracy of evolution equations to greater order in $e$ and $\\tilde{l}^{-1}$. Further work might entail the development of a more precise mathematical treatment of the ans\\\"{a}tze in relation to the underlying physical concepts of the radiation back-reaction process and its effect on the listing behaviour of orbits near the abutment. It would also be intriguing to investigate the consistency condition reported by Gair and Glampedakis \\cite{2006PhRvD..73f4037G}, on the abutment." }, "1101/1101.5866_arXiv.txt": { "abstract": "Spectropolarimetric observations of the pre-main sequence early-G star HD 141943 were obtained at three observing epochs (2007, 2009 and 2010). The observations were obtained using the 3.9-m Anglo-Australian telescope with the UCLES echelle spectrograph and the SEMPOL spectropolarimeter visitor instrument. The brightness and surface magnetic field topologies (given in Paper I) were used to determine the star's surface differential rotation and reconstruct the coronal magnetic field of the star. The coronal magnetic field at the 3 epochs shows on the largest scales that the field structure is dominated by the dipole component with possible evidence for the tilt of the dipole axis shifting between observations. We find very high levels of differential rotation on HD 141943 ($\\sim$8 times the solar value for the magnetic features and $\\sim$5 times solar for the brightness features) similar to that evidenced by another young early-G star, HD 171488. These results indicate that a significant increase in the level of differential rotation occurs for young stars around a spectral type of early-G. Also we find for the 2010 observations that there is a large difference in the differential rotation measured from the brightness and magnetic features, similar to that seen on early-K stars, but with the difference being much larger. We find only tentative evidence for temporal evolution in the differential rotation of HD 141943. ", "introduction": "\\label{Sec_int} One of the key drivers of the solar dynamo is differential rotation. In the Sun strong shears occur in the interface layer between the differentially rotating outer convective zone and the inner radiative zone, which rotates as a solid body. It is these shears that help convert the large-scale solar poloidal field into a strong toroidal component. However, for young, rapidly-rotating, solar-type stars a fundamentally different dynamo may be in operation. Spectropolarimetric observations of young solar-type stars \\citep[i.e.][]{DonatiJF:1997a, DonatiJF:1999a, DonatiJF:1999b, DonatiJF:2003a, MarsdenSC:2006, DunstoneNJ:2008, JeffersSV:2008} have shown that their reconstructed magnetic topologies have large regions of near-surface azimuthal field. These are interpreted as the toroidal component of the large-scale dynamo field in the stars. The presence of these regions near the stellar surface has led to the belief that the dynamo operating in such stars is in fact distributed throughout the stellar convective zone, rather than being restricted to the interface layer as in the solar case. Differential rotation is still thought to play a role in the generation of magnetic fields in these stars, however how the dynamos in these stars operates is not well understood. Most theoretical models are based on our knowledge of the solar dynamo. The models of \\citet{KitchatinovLL:1999} predict that the level of differential rotation on an early-G star should be greater than that on a mid-K star and that the level of surface differential rotation should decrease for stars with shorter rotational periods. While the models of \\citet{KukerM:2005} also show that the level of differential rotation on a star should be dependent upon its effective temperature (with hotter stars having higher levels of differential rotation) but only weakly dependent on its rotation rate. These are predictions that we can now observationally test. There are several techniques we can use to observationally measure the level of differential rotation on a star. Direct starspot tracking from multiple Doppler images of the surface features \\citep*[i.e.][]{CameronAC:2002}, cross-correlating two independent Doppler images \\citep[i.e.][]{DonatiJF:1997a}, incorporating a differential rotation law into the Doppler imaging process \\citep*[i.e.][]{DonatiJF:2000,PetitP:2002}, or by Fourier analysis of stellar line profiles \\citep[i.e.][]{ReinersA:2002, ReinersA:2003}. Combining differential rotation measurements of a number of young solar-type stars, found using the Doppler imaging method, \\citet{BarnesJR:2005} found that the level of differential rotation increases with stellar temperature with early-G stars having higher levels of differential rotation than lower-mass stars, in agreement with the findings of \\citet{KitchatinovLL:1999} and \\citet{KukerM:2005}. The results also showed only a weak (if any) correlation with stellar rotation rate, again in agreement with the findings of \\citet{KukerM:2005}. However, more recent measurements of the young early-G star HD 171488 \\citep{MarsdenSC:2006, JeffersSV:2008, JeffersSV:2010} have shown this star to have an even higher level of differential rotation than that indicated by \\citet{BarnesJR:2005} with differential rotation measurements up to 10 times the current solar value. This supports the findings of \\citet{ReinersA:2006} which show high levels of differential rotation on a number of F stars. Work by \\citet*{DonatiJF:2003b} has shown that for some early-K stars the level of differential rotation measured from surface brightness features is consistently lower than that measured from the magnetic features (using the same dataset). This has led them to surmise that this is a result of the magnetic and brightness features being anchored at different depths within the stellar convective zone and that the convective zone has a radially varying differential rotation structure, unlike the Sun. The work also showed that these early-K stars evidence temporal variation in their levels of differential rotation that the authors attribute to a feed-back mechanism in the stellar dynamo periodically converting magnetic energy into kinetic energy and vice-versa. As mentioned there is currently only one early-G star for which there are differential rotation measures from both brightness and magnetic features obtained at multiple epochs, HD 171488 \\citep{MarsdenSC:2006, JeffersSV:2008, JeffersSV:2010}. This work has shown very little difference in the level of differential rotation measured from the magnetic and brightness features, although the measurement errors are larger than that found on the early-K stars \\citep{DonatiJF:2003b}. Additionally, there appears to be little evidence of temporal evolution in the level of differential rotation on HD 171488. This has been speculated to be caused by the thinner convective zone of HD 171488 compared to that of the early-K stars previously studied. In order to expand the number of young early-G stars studied with spectropolarimetry this paper along with the first paper in the series \\citet[][Paper I]{MarsdenSC:2010} and \\citet{WaiteIA:2010}, presents differential rotation measurements, and magnetic maps, of two young early-G pre-main sequence (PMS) stars. Paper I deals with the reconstruction of the brightness and magnetic topologies of the star HD 141943. This paper (Paper II) presents the coronal magnetic field reconstructions, H$\\alpha$ activity and differential rotation measurements of HD 141943. The third paper in the series \\citep{WaiteIA:2010} deals with observations of HD 106506. As detailed in Paper I, HD 141943 is a young, active and rapidly-rotating PMS star. The stellar parameters have been determined in Paper I through the Doppler imaging process and are reproduced here in Table~\\ref{Tab_par}. \\begin{table} \\caption{Fundamental parameters of HD 141943 from Paper I.} \\label{Tab_par} \\centering \\begin{tabular}{lc} \\hline\\hline Parameter & value\\\\ \\hline Age & $\\sim$17 Myrs\\\\ Mass & $\\sim$1.3 M\\subs{\\odot}\\\\ Photospheric temperature & 5850 $\\pm$ 100 K$^{a}$\\\\ Spot temperature & $\\sim$3950 K\\\\ Unspotted luminosity & 2.8 $\\pm$ 0.1 L$^{a}_{\\odot}$\\\\ Stellar radius & 1.6 $\\pm$ 0.15 R\\subs{\\odot}\\\\ \\vsinis & 35.0 $\\pm$ 0.5 \\kmss\\\\ Radial velocity ($v_{\\rm rad}$) & $\\sim$0.1 \\kmss\\\\ Inclination angle ($i$) & 70\\sups{\\circ} $\\pm$ 10\\sups{\\circ}\\\\ Equatorial rotation period ($P_{\\rm eq}$) & $\\sim$2.182 days\\\\ \\hline $^{a}$assumed errors (see Paper I). \\end{tabular} \\end{table} The reconstructed brightness images show that it has a weak polar spot and a significant amount of low-latitude spot features at all epochs. The magnetic reconstructions show that its has a predominately non-axisymmetric radial field while its azimuthal field is predominately axisymmetric with ring of azimuthal field seen at the pole, similar to that of other active stars. ", "conclusions": "\\label{Sec_con} We have extrapolated the coronal magnetic field of HD 141943 from surface magnetograms acquired at 3 epochs. These show that the large-scale field structure of the corona is dominated by a dipole component with the axis of the dipole shifting between the 3 epochs. The small scale structure shows an increase in the modelled rotational X-ray modulation in 2009, compared to the other epochs (2007 \\& 2010). The surface differential rotation of HD 141943 has been measured at 2 epochs (2007 \\& 2010) from both the surface brightness features (2010) and the surface magnetic features (2007 \\& 2010). The differential rotation measured from the magnetic features is one of the highest values of differential rotation measured on a young solar-type star and is similar in level to the other young early-G star HD 171488, but higher than the more swollen star HD 106506. We thus conclude that the depth of the stellar convective zone plays a strong role in the level of surface differential rotation seen on solar-type stars, with a large increase in differential rotation seen for star's with convective zone depths shallower than $\\sim$0.2 R\\subs{\\star}. The 2010 dataset for HD 141943 shows a large increase in the level of differential rotation measured from magnetic features to that measured from brightness features. This is similar to that seen on early-K stars but with a much greater difference and is in contrast to the results from other early-G stars which show little or no difference between the differential rotation measured from brightness and magnetic features. Our results only find tentative evidence for temporal evolution in the differential rotation of HD 141943. These results when combined with those from the early-G star HD 171488 (which shows no evidence of temporal evolution in differential rotation) imply that early-G stars do not undergo large-scale evolution in their differential rotation. However, the errors in our measurements are too large to rule out small scale evolution in differential rotation similar to that seen on early-K stars. HD 141943 and stars of similar spectral type warrant further observations to determine what effect a shallow convective zone has on the differential rotation levels of such stars and indeed if they do show temporal evolution of their differential rotation as seen on early-K stars." }, "1101/1101.3321_arXiv.txt": { "abstract": "{Of the many ways of detecting high redshift galaxies, the selection of objects due to their redshifted Ly$\\alpha$ emission has become one of the most successful. But what types of galaxies are selected in this way? Until recently, Ly$\\alpha$ emitters were understood to be small star-forming galaxies, possible building-blocks of larger galaxies. But with increased number of observations of Ly$\\alpha$ emitters at lower redshifts, a new picture emerges. Ly$\\alpha$ emitters display strong evolution in their properties from higher to lower redshift. It has previously been shown that the fraction of ultra-luminous infrared galaxies (ULIRGs) among the Ly$\\alpha$ emitters increases dramatically between redshift three and two. Here, the fraction of AGN among the LAEs is shown to follow a similar evolutionary path. We argue that Ly$\\alpha$ emitters are not a homogeneous class of objects, and that the objects selected with this method reflect the general star forming and active galaxy populations at that redshift. Ly$\\alpha$ emitters should hence be excellent tracers of galaxy evolution in future simulations and modeling. } ", "introduction": "Several decades of studies of high redshift galaxies have shown that the star formation rate density of the Universe peaked around redshift $z \\sim 2$ (e.g. Hogg et al.~1998, Hopkins 2004, Hopkins \\& Beacom 2006). At the peak of the star formation history, nearly ten times more stars were formed than in our local Universe, whereas at higher redshifts, the star formation density was equally low as it is now. Similarly, a trend in the volume density of AGN has been found with a peak at $z = 1.5 - 2$ (e.g. Miyaji et al.~2000, Wolf et al.~2003, Bongiorno et al.~2007). The coincidence that the two density functions, for star formation and numbers of AGN, peak at similar redshifts has been proposed as being due to both of these properties being linked to the hierarchical build-up of galaxies and mergers of dark matter haloes (e.g. Kauffmann \\& Haehnelt~2000, Bower et al.~2006). As for the high redshift Universe, one of the strongest emission lines observable is the Lyman-$\\alpha$ (Ly$\\alpha$) line. By now, several hundreds of Ly$\\alpha$ emitters (LAEs) have been detected through narrow-band imaging at $z = 0.3 - 7.7$ (e.g. M{\\o}ller \\& Warren 1993, Fynbo et al.~2003, Gronwall et al.~2007, Venemans et al.~2007, Nilsson et al.~2007, Finkelstein et al.~2007, Ouchi et al.~2008, Grove et al.~2009, Hibon et al.~2010). Ly$\\alpha$ emission may be generated by three main mechanisms; the ionising flux of O and B stars, indicative of star formation, the ionising flux of an energetic UV source, e.g. an active galactic nucleus (AGN), or due to infall of gas on a massive dark matter halo (c.f. Dijkstra et al.~2006a,b, Nilsson et al.~2006). The volume density of sources where the Ly$\\alpha$ emission is dominated by the latter is expected to be very low compared to those where the Ly$\\alpha$ emission comes from star formation or AGN sources, hence, the volume density of Ly$\\alpha$ emitting objects found in the Universe is expected to follow the general evolutionary occurrences of the star formation history, and the AGN history, with redshift. In this \\emph{Letter} we discuss the fractions of ULIRGs and AGN among Ly$\\alpha$ emitters at different redshifts. In a previous publication, the ULIRG fraction among the LAEs was shown to exhibit a sharp transition from very few to a larger sub-sample at a redshift around 2.5 (Nilsson \\& M{\\o}ller 2009). Here, the fraction of AGN among LAEs is shown to follow a very similar relation, indicating that the underlying galaxy population is transitioning rapidly from $z > 3$ to $z \\sim 2$, (see also a similar result in Bongiovanni et al.~2010). We here ask the question how these results relate to the general galaxy evolution in the Universe. \\vskip 5mm Throughout this paper, we assume a cosmology with $H_0=72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega _{\\rm m}=0.3$ and $\\Omega _\\Lambda=0.7$. ", "conclusions": "Based on the apparent evolution of the AGN and ULIRG fractions of LAEs, fractions that are expected to follow the evolution of star forming and active galaxy populations throughout the Universe, some conclusions can be drawn regarding the general galaxy evolution scenario. From the evolutions of the fractions and volume densities in Fig.~\\ref{fig:agnevol} and \\ref{fig:laesurfdens} it is seen that at first, in the very young ($z > 4$) Universe, only stars formed, regardless of AGN or ULIRGs. Later, at an age of approximately 3 Gyrs or redshift $\\sim 2.5$, a secondary process started, suddenly increasing the number of both AGN and ULIRGs in the galaxy population. After this abrupt increase in dusty and active galaxies, an equilibrium was reached that has lasted to this day. A possible explanation for this very stable equilibrium may be due to feed-back effects controlling the SFR in the galaxies. It is clear that this remains speculation until further data is collected, but proves the importance of studying objects emitting Ly$\\alpha$ emission in the Universe." }, "1101/1101.4995_arXiv.txt": { "abstract": "In this paper, we explore an idea of having Newton's constant change its value depending on the curvature scale involved. Such modification leads to a particular scalar-tensor gravity theory, with the Lagrangian derived from renormalization group (RG) flow arguments. Several of the well-known $f(R)$ modified gravity models have remarkably simple description in terms of the infrared renormalization group, but not the ``designer'' types in general. We find that de Sitter-like accelerated expansion can be generated even in the absence of cosmological constant term, entirely due to running of the Newton's constant. In hopes of tackling the problem of cosmological constant's smallness, we explore the flows which are capable of generating exponential hierarchy between infrared and ultraviolet scales, and investigate cosmological evolution in the models thus derived. ", "introduction": "It has long been hoped that quantum theory of gravity, at least in some limit, allows description in terms of an effective field theory \\cite{Weinberg:2009bg, Weinberg:2009wa}. The usual Einstein-Hilbert action is merely the first two terms of an effective action \\begin{equation}\\label{eq:eft} S[\\lambda] = \\int\\left\\{\\sum\\limits_{n=0}^{\\infty} \\lambda^{4-2n} g^{(n)}(\\lambda)\\, {\\cal R}^{(n)} + \\ldots \\right\\}\\sqrt{-g}\\,d^4x, \\end{equation} expanded in local $n$-th order curvature invariants ${\\cal R}^{(n)}$, which explicitly depends on the ultraviolet cutoff scale $\\lambda$. Should the effective theory of gravity be asymptotically safe, this would allow a sensible ultraviolet-complete description \\cite{Niedermaier:2006ns, Percacci:2007sz}. However, recent observations of accelerated expansion of the Universe indicate the presence of a tiny but nonvanishing cosmological constant \\cite{Riess:1998cb, Perlmutter:1998np, Riess:2004nr, Komatsu:2010fb}, which presents technical hierarchy problem between (infrared) cosmological acceleration scale and (ultraviolet) Planck scale, which seems irreconcilable despite the many efforts put forward \\cite{Weinberg:1988cp, Peebles:2002gy}. Given enormous separation of cosmological and Planck scales, it would seem highly dubious that quantum gravity is somehow responsible for cosmological acceleration. Corrections and corresponding beta functions reviewed in \\cite{Niedermaier:2006ns, Percacci:2007sz} originate from ultraviolet degrees of freedom; they are Planck-suppressed and irrelevant for cosmology. But the language of effective field theory is universal, and should be applicable to infrared phenomena on cosmological scales as well. To give cosmological acceleration a dynamical origin without invoking a new dark energy matter component, one would need a ``modified'' gravity, with new infrared degrees of freedom becoming active at cosmological scales, while remaining hidden in solar system and local tests. A lot has been said on the subject, with $f(R)$ gravity models in particular receiving a lot of attention as of late \\cite{Sotiriou:2008rp}, yet the question of large hierarchy ever remains in these attempts. The main aim of this paper is to apply the idea of renormalization group, which has been very fruitful in high energy physics, to description of these new (and still unknown) infrared degrees of freedom. This obviously entails the shift of attention from ultraviolet cutoff in the effective action (\\ref{eq:eft}) to the lower limit of integration -- the {\\em infrared cutoff}, and on how things scale when it is varied. Of particular interest is the question whether it is even in principle possible to generate anomalously low scale of cosmological constant by specific running, perhaps in a way similar to dimensional transmutation phenomenon, which is what we will discuss here. Analyzing renormalization flows dependent on curvature scale, we find that de Sitter-like accelerated expansion can be generated even in the absence of cosmological constant term, entirely due to running of the Newton's constant. This presents a novel way to view the hierarchy problem, in contrast with previous studies \\cite{Machado:2007ea, Shapiro:2009dh} involving running of $\\Lambda$. Cosmological constant should be set to zero and protected by symmetry, while de Sitter asymptote is attributable to flow of Newton's constant, which we describe by renormalization group equations. Several historically important modified gravity models of $f(R)$ type \\cite{Starobinsky:1980te, Capozziello:2003tk, Carroll:2003wy} turn out to have a very simple description in this terms, but not the ``designer'' types, in general. \\begin{figure*} \\begin{center} \\begin{tabular}{c@{\\hspace{24pt}}c} \\begin{tikzpicture}[scale=1.8] \\filldraw[color=red!20!white] (-1,-1) -- (1,1) -- (-1,1) -- cycle; \\node[color=red!40!black] at (-0.1,0.1) {\\rotatebox{45}{forbidden region (ghosts)}}; \\node[color=black] at (2.6,-2.4) {\\rotatebox{-45}{effective $\\Lambda$}}; \\draw[->] (-1.1,0) -- (3,0) node[right] {$\\alpha$}; \\draw[->] (0,-3) -- (0,1.1) node[above] {$\\beta$}; \\clip (-1,1) rectangle (2.8,-2.8); \\draw[very thick, color=blue] (0,0) -- (3,-3); \\draw[very thick, color=red] (-1,-1) -- (1,1); \\node[color=red] at (-0.4,-0.6) {\\rotatebox{45}{$\\alpha-\\beta=0$}}; \\node[color=blue] at (0.4,-0.6) {\\rotatebox{-45}{$\\alpha+\\beta=0$}}; \\draw[color=green!40!black] (0,0) parabola bend (0.5,0.25) (3,-6); \\draw[color=blue!40!black] (0,0) parabola bend (0.5,-0.25) (3,6); \\draw[color=red!40!black] (0,0) parabola bend (0.5,0.50) (3,-12); \\draw (1,0) circle (0.05) node[below right] {~~~GR fixed point}; \\draw[color=green!40!black] (2,-2) circle (0.05) node[above right] {dS in IR}; \\draw[color=red!40!black] (1.5,-1.5) circle (0.05) node[left] {unstable dS~~}; \\node[above,color=green!40!black] at (1.7,-1.5) {\\rotatebox{-66}{\\scriptsize$f(R)=R-2\\Lambda$}}; \\node[below,color=blue!40!black] at (1.5,1.12) {\\rotatebox{60}{\\scriptsize$f(R) = R +\\! \\frac{R^2}{M^2}$}}; \\node[below,color=red!40!black] at (1.55,-1.65) {\\rotatebox{-80}{\\scriptsize$f(R) = R - \\!\\frac{\\mu^4}{R}$}}; \\end{tikzpicture} & \\begin{tikzpicture}[scale=1.8] \\filldraw[color=red!20!white] (-1,-1) -- (1,1) -- (-1,1) -- cycle; \\node[color=red!40!black] at (-0.1,0.1) {\\rotatebox{45}{forbidden region (ghosts)}}; \\node[color=black] at (2.6,-2.4) {\\rotatebox{-45}{effective $\\Lambda$}}; \\draw[->] (-1.1,0) -- (3,0) node[right] {$\\alpha$}; \\draw[->] (0,-3) -- (0,1.1) node[above] {$\\beta$}; \\clip (-1,1) rectangle (2.8,-2.8); \\draw[very thick, color=blue] (0,0) -- (3,-3); \\draw[very thick, color=red] (-1,-1) -- (1,1); \\node[color=red] at (-0.4,-0.6) {\\rotatebox{45}{$\\alpha-\\beta=0$}}; \\node[color=blue] at (0.4,-0.6) {\\rotatebox{-45}{$\\alpha+\\beta=0$}}; \\draw (0,0) parabola (3,-9); \\draw (0.3,0) parabola (3.3,-9); \\draw[color=black] (1.541619849,-1.541619849) circle (0.05); \\draw[color=green!40!black] (0,0) parabola bend (0.5,0.25) (3,-6); \\draw (1,0) circle (0.05) node[above right] {GR fixed point}; \\draw[color=green!40!black] (2,-2) circle (0.05) node[above right] {dS in IR}; \\node[above,color=green!40!black] at (1.7,-1.5) {\\rotatebox{-66}{\\scriptsize$f(R)=R-2\\Lambda$}}; \\draw[color=black] (1,-1) circle (0.05) node[below left] {stable dS}; \\node[below] at (1.35,-1.8) {\\rotatebox{-72}{$f(R)$ flow}}; \\end{tikzpicture} \\\\ (a) & (b) \\\\ \\end{tabular} \\end{center} \\caption{Renormalization group flows causing cosmological acceleration: several known $f(R)$ models that can be generated by renormalization group (a), and possible flows generating exponential hierarchy between infrared and ultraviolet scales (b).} \\label{fig:flow} \\end{figure*} We argue that a ``soft'' running, in which Newton's constant flows with quadratic beta function approaching a high-curvature fixed point, generates exponential hierarchy between infrared and ultraviolet scales, which could be exploited for cosmological model-building. The simplest model of this type, with beta function analogous to the one for QCD coupling constant, is somewhat problematic observationally (due to Newton's constant running to zero in high curvature limit), but is easy to analyze analytically. We present it here due to its simplicity, and then discuss more realistic models with exponential hierarchy derived by this method, which did not appear in the literature before. ", "conclusions": "The main result of this paper is that running Newton's constant can cause accelerated expansion of the universe. The good news is that the simplest model of this type (\\ref{eq:beta:IR},\\ref{eq:f}) is predictive and has essentially the same number of parameters as standard $\\Lambda$CDM, with exponential hierarchy between effective cosmological constant (\\ref{eq:lambda}) and UV scales naturally generated by renormalization group flow (\\ref{eq:beta}). The bad news is that while the recent cosmological expansion history is plausibly reproduced, the Newton's constant changes substantially (\\ref{eq:eq}) between cosmological, galactic, and near-Earth environments in the simplest realization discussed here. Thus, although the scalar degree of freedom is not light inside matter, this model might have difficulty with gravity tests that probe absolute value of Newton's constant, for example constraints on expansion during nucleosynthesis epoch. The difficulties with radically changing Newton's constant can be avoided by having a quadratic beta function flow to a finite UV fixed point instead (\\ref{eq:beta:2}), which leads to models of the type (\\ref{eq:f:2}). These models share similar features with respect to running and hierarchy as the simplest realization we discussed in detail above, but Newton's constant asymptotes to a finite value in the regions of high curvature, as in Einstein's gravity. Taking, for definiteness, $\\alpha_* = 1/4$ and $\\kappa = 4/3$, leading to \\begin{equation}\\label{eq:example} f(R) = \\frac{1 + \\ln\\frac{R~}{R_0}}{1 + \\frac{1}{4}\\ln\\frac{R~}{R_0}}\\, R, \\end{equation} one can repeat the analysis of the previous section. While algebraically more complicated, the story goes pretty much along the same lines. Effective potential $V(\\phi)$ still has a single minimum and an infinite potential wall at high curvature, but develops a turning point at the low curvature limit. For sensible cosmology, this turning point must be below asymptotic de Sitter curvature in value, which places certain condition on the flow (\\ref{eq:beta:2}). The constants in (\\ref{eq:example}) are chosen so that this is the case, otherwise no effort was made to tune them to ``special'' values or to fit $\\Lambda$CDM fiducial model. With these values, the tracker solution is parametrized by $\\Omega_{\\cal Q} = 0.722$, and is shown in the right panel of Fig.~\\ref{fig:modulus}. The deviation of distance modulus from fiducial $\\Lambda$CDM model is about 1/10 of magnitude out to redshift $z=2$, and $\\Omega$ parameters plot looks much more conventional now. The tell-tale sign of this model is much slower dilution of $\\Omega_{\\cal Q}$ at high redshift than usual, with $\\Omega_M$ never quite reaching $1$ during matter domination in this example. This will surely affect the structure formation, and should give a way to test or rule out the model. While parameter space is larger and still remains to be completely explored, this ``soft'' approach to Einstein's gravity is due to logarithmic dependence and is characteristic for all the models we presented here. Severity of bounds placed by solar system tests \\cite{Chiba:2006jp, Hu:2007nk} and large scale structure growth \\cite{Pogosian:2007sw, Oyaizu:2008tb, Schmidt:2008tn} will also need to be investigated." }, "1101/1101.1262_arXiv.txt": { "abstract": "{The question whether or not the initial mass function is universal, i.e.\\ the same in all kinds of environments, is of critical importance for the theory of star formation and still intensely debated. A top-heavy initial mass function may be the result of star formation out of dense molecular clouds with a temperature of $\\sim 100$\\,K. Such a molecular gas phase is not commonly found in the Galactic plane, but may be present in active environments like cores of starburst galaxies or AGN. Unfortunately, the kinetic temperature of the molecular gas in external galaxies is often not well constrained. Having proven the diagnostic power of selected formaldehyde lines as tracers of the properties of the molecular gas in external galaxies, we have engaged in observing these diagnostic lines in a number of starburst galaxies or near AGN. This contribution presents the latest results of these studies.} \\FullConference{10th European VLBI Network Symposium and EVN Users Meeting: VLBI and the new generation of radio arrays \\\\ September 20-24, 2010\\\\ Manchester, UK} \\newcommand{\\htco}{H$_2$CO} \\newcommand{\\kms}{km\\,s$^{-1}$} \\begin{document} \\begin{figure}[h] \\begin{center} \\includegraphics[width=.6\\textwidth]{mangum.eps} \\end{center} \\caption{The lowest energy levels of ortho- and paraformaldehyde, adapted from [7]. The para-H$_2$CO transitions at 218\\,GHz are marked in red, the diagnostic para-H$_2$CO transition at 146\\,GHz in blue.} \\label{fig1} \\end{figure} ", "introduction": " ", "conclusions": "" }, "1101/1101.2782_arXiv.txt": { "abstract": "{Results of processing of data of a VLBI experiment titled RAPL01 are presented. These VLBI observations were made on 4th February, 2010 at 6.28 cm between the 100-m antenna of the Max Planck Institute (Effelsberg, Germany), Puschino 22-m antenna (Astro Space Center (ASC), Russia), and two 32-m antennas of the Istituto di Radioastronomia di Bologna (Bologna, Italy) in Noto and Medicina. 2 well-known sources, 3C84 (0316+413), and BL Lac (2200+420) were included in the schedule of observations. Each of them was observed during 1 hour at all the stations. The Mark-5A registration system was used at 3 European antennae. The alternative registration system known as RDR (RADIOASTRON Data Recorder) was used in Puschino. The Puschino data were recorded in format RDF (RADIOASTRON Data Format). Two standard recording modes designed as 128-4-1 (one bit), and 256-4-2 (two bit) were used in the experiment. All the Mark-5A data from European antennae were successfully converted into the RDF format. Then, the correlation function was estimated at the ASC software correlator. A similar correlation function also was estimated at the Bonn correlator. The Bonn correlator reads Mark5A data, the RDF format was converted into Mark5B format before correlation. The goal of the experiment was to check the functioning and data analysis of the ground based radio telescopes for the RADIOASTRON SVLBI mission} \\FullConference{GAMOW-2010\\\\ August 23 - 27 2010\\\\ Odessa, Ukraine} \\begin{document} ", "introduction": " ", "conclusions": "Thus, we could make the following conclusions : \\begin{enumerate} \\item The experiment RAPL01 demonstrates the possibility to convert the Mark-5A data into RDF data. Antennae with different registration systems could be successfully used for the RADIOASTRON mission\\\\ \\item The integration time value is restricted by 1 second due to a high rate offset at Puschino antenna. The successful estimation of the correlation function demonstrates the possibilities of the ASC software correlator to compensate correctly the abnormally high values of residual delays and fringe rates\\\\ \\item The data at the end of the ASC software correlator are relevant for the secondary processing\\\\ \\item The calibration procedures of the software known as Astro Space Locator allow reconstructing the visibility function\\\\ \\item The (u, v)-plane coverage for the 3C84 is not sufficient to perform the source imaging. The value of estimated source angular size is 11 mas. This value is consisted with the 3C84 properties available in literature\\\\ \\end{enumerate} All the results presented in this paper are preliminary.\\\\ Procedures and techologies used during the VLBA data processing also could be very useful for processing of data of future Space VLBI mission titled RADIOASTRON." }, "1101/1101.0322_arXiv.txt": { "abstract": "We summarize the contribution of the HATNet project to extrasolar planet science, highlighting published planets (HAT-P-1b through HAT-P-26b). We also briefly discuss the operations, data analysis, candidate selection and confirmation procedures, and we summarize what HATNet provides to the exoplanet community with each discovery. ", "introduction": "The Hungarian-made Automated Telescope Network (HATNet; Bakos et al.~2004) survey, has been one of the main contributors to the discovery of transiting exoplanets (TEPs), being responsible for approximately a quarter of the $\\sim 100$ confirmed TEPs discovered to date (Fig.~1). It is a wide-field transit survey, similar to other projects such as Super-WASP (Pollaco et al.~2006), XO (McCullough et al.~2005), and TrES (Alonso et al.~2004). The TEPs discovered by these surveys orbit relatively {\\em bright} stars ($V < 13$) which allows for precise parameter determination (e.g.~mass, radius and eccentricity) and enables follow-up studies to characterize the planets in detail (e.g.~studies of planetary atmospheres, or measurements of the sky-projected angle between the orbital axis of the planet and the spin axis of its host star). Since 2006, HATNet has announced twenty-six TEPs\\footnote{Meaning that the scientific results were submitted to peer reviewed journals and posted to arXiv, the planet host stars have been uniquely identified, and all discovery data have been made public.}. Below we highlight some of the exceptional properties of these planets (Section~2), we then describe the procedures which we followed to discover them (Section~3), and we conclude by summarizing what HATNet provides to the TEP community with each discovery (Section~4). \\begin{figure}[ht] \\begin{center} \\epsfig{width=10cm,file=tepstat_perproject.eps} \\caption{ Mass--radius relation of TEPs, highlighting the findings from different surveys. The dotted lines are lines of constant density from 0.4\\,g\\,cm$^{-3}$ to 11.9\\,g\\,cm$^{-3}$. Also overlaid are models from Fortney et al.~2007. } \\end{center} \\end{figure} ", "conclusions": "" }, "1101/1101.5507_arXiv.txt": { "abstract": "We present exact solutions of the incompressible Navier--Stokes equations in a background linear shear flow. The method of construction is based on Kelvin's investigations into linearized disturbances in an unbounded Couette flow. We obtain explicit formulae for all three components of a Kelvin mode in terms of elementary functions. We then prove that Kelvin modes with parallel (though time--dependent) wavevectors can be superposed to construct the most general plane transverse shearing wave. An explicit solution is given, with any specified initial orientation, profile and polarization structure, with either unbounded or shear--periodic boundary conditions. ", "introduction": " ", "conclusions": "" }, "1101/1101.4136_arXiv.txt": { "abstract": "{Cluster faint low surface brightness galaxies (fLSBs) are difficult to observe. Consequently, their origin, physical properties and number density are not well known. After a first search for fLSBs in the highly substructured Coma cluster, we present here a search for fLSBs in the nearly relaxed Abell~496 cluster.} {Abell~496 appears to be a much more relaxed cluster than Coma, but still embedded in a large scale filament of galaxies. Our aim is to compare the properties of fLSBs in these two very different clusters, to search for environmental effects.} {Based on deep CFHT/Megacam images in the $u^*$, $g'$, $r'$ and $i'$ bands, we selected galaxies with $r'>21$ and $\\mu_{\\rm{r'}}> 24$ mag arcsec$^{-2}$. We estimated photometric redshifts for all these galaxies and kept the 142 fLSBs with photo$-z<0.2$.} {In a $g'-i'$ versus $i'$ color-magnitude diagram, we find that a large part of these fLSBs follow the red sequence (RS) of brighter galaxies. The fLSBs within $\\pm 1\\sigma$ of the RS show a homogeneous spatial distribution, while those above the RS appear to be concentrated along the large scale filament of galaxies.} {These properties are interpreted as agreeing with the idea that RS fLSBs are formed in groups prior to cluster assembly. The formation of red fLSBs could be related to infalling galaxies.} ", "introduction": "Faint low surface brightness galaxies (fLSBs hereafter) remain a poorly known class of galaxies, though they are interesting objects for several reasons, as already discussed in detail by Adami et al. (2009a). We define fLSBs as galaxies with a central surface brightness fainter than $\\mu_{\\rm{r'}}$ = 24 mag arcsec$^{-2}$ and a total magnitude $r'>21$, to be consistent with Adami et al. 2006, hereafter ASU06. Briefly: fLSBs could account for part of the missing low luminosity structures predicted by CDM models of hierarchical structure formation (White \\& Rees 1978), in particular since they appear dominated by dark matter (e.g. McGaugh et al. 2001, de Blok et al. 2001). CDM models predict the existence of low luminosity galaxies in all environments, but fLSBs seem to be present in higher numbers in clusters than in the field (see e.g. Sabatini et al. 2005, ASU06, and references therein). Many fLSBs are fainter than the night sky and clearly extend toward fainter brightnesses than predicted by the Freeman law (1970), as shown for example by Bothun et al. (1997). Due to their extreme faintness both in terms of surface brightness and of total magnitude, fLSBs are therefore very difficult to detect, hence their origin, physical properties and number density are not well known in a statistical way over a large number of clusters, despite numerous studies (e.g. Binggeli et al. 1985; Schombert et al. 1992; Bothun et al. 1993; Bernstein et al. 1995; Impey et al. 1996; Sprayberry et al. 1996; Ulmer et al. 1996; Impey \\& Bothun 1997; O'Neil et al. 1997; Kuzio de Naray et al. 2004). In order to increase the number of fLSBs detected in clusters, our team has searched for Coma cluster fLSBs in the total magnitude versus central surface brightness space (ASU06, Adami et al. 2009a) and found for example that these objects tended to be more concentrated in several areas (not always central). Furthermore, based on their position in the (B$-$R) versus R plane, we found that we could identify three distinct types of fLSBs. Those that fall on the color magnitude relation extrapolated from the bright normal galaxy population we called $sequence$ fLSBs. We interpreted $sequence$ fLSBs as galaxies that formed in small groups prior to the cluster assembly. Then we interpreted the reddest fLSBs as faint stripped ellipticals and the blue fLSBs as galaxies made of material stripped from spiral infalling galaxies. However, the Coma cluster is highly substructured (e.g. Adami et al. 2005) and we do not know how substructure could affect the spatial distribution of the fLSB population. We therefore decided to analyze in the same way the distribution and properties of fLSBs in a more relaxed cluster where substructures will not complicate the picture. Abell~496 is one of the rare nearby nearly relaxed clusters (see e.g. Durret et al. 2000). Bou\\'e et al., (2008) reported the detailed analysis of the galaxy luminosity functions of Abell~496, based on deep CFHT Megacam images in four bands which are ideal to search for fLSBs. They confirmed that this cluster appears very relaxed, with no particular structure at the cluster scale, though at larger scale an extended filament of galaxies with redshifts close to that of Abell~496 was found to spread from the north-west to the south-east of the cluster (see Fig.~10 in Bou\\'e et al. 2008). The mean heliocentric velocity of Abell~496 is cz=$9885 \\, \\rm km \\, s^{-1}$, corresponding to a redshift $z = 0.0329$, its distance modulus is 35.69, and the scale is 0.666 kpc arcsec$^{-1}$, assuming H$_0=72$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$ and $\\Omega_{\\Lambda}=0.7$. It has an angular virial radius of $0.77^\\circ$ (1.85~Mpc), obtained by extrapolating the radius of overdensity 500 (Markevitch et al. 1999), measured relative to the critical density of the Universe to the radius of overdensity 100. We will give magnitudes in the AB system. The paper is organized as follows. The data and method to search for fLSBs are described in Section~2. Results concerning the color-magnitude relation, spatial distribution and luminosity function of fLSBs are presented in Section~3 and discussed in Section~4. We give in the Appendix the list of the 142 fLSBs with photo$-z<0.2$ as well as the images in the four bands and the surface brightness profile for one of them. ", "conclusions": "\\label{sec:discu} As described above, we have found 142 fLSBs in the direction of Abell~496 with photo$-z<0.2$, out of which about 80\\% are probably cluster members. Their angular density profile is well fit by a King model with a core radius about twice as large as for normal galaxies. The King distribution of fLSBs in Abell~496 is very different from what was observed in Coma by ASU06, where fLSBs do not follow any King-like distribution. This difference is consistent with the idea that Abell~496 is relaxed while Coma is not. Furthermore, the wider radial distribution of the fLSBs versus normal galaxies in Abell~496 is consistent with the idea that mass segregation has occurred in Abell~496. The detected fLSBs fall reasonably well on the extension of the bright end of the color-magnitude relation established by Bou\\'{e} et al. (2008). The fact that we have found (see section 3.1) the $\\pm 1\\sigma$ interval for the fLSBs around the red sequence similar in Abell~496 and Coma, one a relaxed cluster, the other not, fits with the idea that the relaxation state of the cluster does not influence the position of the fLSBs on the red-sequence. The similar red-sequence width in both clusters could be attributed to $sequence$ fLSBs having evolved in similar groups that fell into the clusters later, as suggested by ASU06. On scales of $\\geq 1$~Mpc, we note that there is a filament in the normal galaxy population with redshifts $< 0.2$ found by Bou\\'{e} et al (2008). The filament extends along a north-west to south-east line. In Fig.~4 we can see that $red$ fLSBs (with redshifts $< 0.2$) seem to have an anisotropic distribution similar to the filament found by Bou\\'{e} et al. However, {\\it blue} fLSBs show no obvious anisotropic distribution, suggesting they had a different evolutionary history. {\\it Blue} fLSBs are perhaps the remnants of tidally disrupted late-type galaxies as hypothesized by ASU06 for Coma. In terms of tidal disruption, we note that the spatial distribution of fLSBs seems to show no holes in the cluster center, which is not the case for Coma (ASU06). For Coma the fLSBs could have been destroyed by tidal disruption due to the massive D galaxies in the Coma core. In contrast, there is only one central galaxy in the center of Abell~496, which could produce much less tidal disruption. It is beyond the scope of this work, though, to carry out numerical simulations to verify or falsify the idea that fLSBs are tidally destroyed in the core of Coma and not in Abell~496." }, "1101/1101.1525_arXiv.txt": { "abstract": " ", "introduction": "% Despite decades of model building, we are still looking for a preferred, compelling, UV-complete model of inflation. Some scenarios, however, appear to be more promising than others. In particular, due to their radiative stability, models where the inflaton is a pseudo-Nambu-Goldstone boson (pNGb~\\cite{Freese:1990rb}) have received a significant amount of interest (see e.g.~\\cite{models}). PNGbs are pseudoscalar particles: as they roll down their potential, they provide a macroscopic source of parity violation. The goal of this paper is to present a simple mechanism able to imprint such parity violation on the Cosmic Microwave Background. A pseudoscalar inflaton $\\phi$ (especially a pNGb, that is characterized by a broken shift symmetry) is generically expected to interact with gauge fields through the coupling $\\phi\\,F_{\\mu\\nu}\\,\\tilde{F}^{\\mu\\nu}/f$. As $\\phi$ rolls down its potential, it provides a time-dependent background for the quantization of the gauge field, amplifying its vacuum fluctuations into classical modes~\\cite{Garretson:1992vt}, that in their turn are a source of gravitational waves. Because of the parity-violating nature of the system, only photons of a given helicity are produced~\\cite{Anber:2006xt}, implying that gravitational waves of different helicities are produced with different amplitude. This way, parity violation in the inflaton sector manifests itself as a different amplitude of the power spectrum of the left-handed modes ${\\cal P}^{t,-}$ of the graviton with respect to that, ${\\cal P}^{t,+}$, of the right-handed ones. A measure of the net handedness of the tensor modes is the parameter $\\Delta\\chi\\equiv \\left({\\cal P}^{t,+}-{\\cal P}^{t,-}\\right)/\\left({\\cal P}^{t,+}+{\\cal P}^{t,-}\\right)$~\\cite{Saito:2007kt,Gluscevic:2010vv}. As we will see, in this scenario the value of $\\Delta\\chi$ will depend only on the value of the Hubble parameter $H$ during inflation and (exponentially) on the combination $\\dot\\phi/(H\\,f)$, and can easily attain values indistinguishable from unity. How do parity-violating tensor modes show up in the CMB? Tensor modes produced during inflation leave a trace on the CMB in the form of B-modes, the divergence-free component of the polarized radiation. While B-modes have not been yet detected, the sensitivity of CMB experiments to these modes will improve substantially in the next years. B-modes are parity-odd, whereas E-modes (the curl-free component of the polarized CMB radiation) as well as the temperature fluctuations are parity-even. As a consequence, a nonvanishing $\\langle B\\,E\\rangle$ or $\\langle B\\,T\\rangle$ correlation will signal parity violation in the CMB~\\cite{Lue:1998mq}. In the best case scenario, values of $\\Delta\\chi$ as small as $.3$ might be detected at the $3\\,\\sigma$ level in a cosmic-variance limited experiment~\\cite{Saito:2007kt,Gluscevic:2010vv}. Of course, in order to be viable, the model has to satisfy all experimental constraints. In this system, the main constraint comes from the requirement that nongaussianities are below the observational limits~\\cite{Barnaby:2010vf}, and originates from the fact that the source of chiral tensor modes is also a source of scalar modes. Since these modes arise from a second order effect, they are intrinsically nongaussian. We discuss two scenarios where this constraint does not apply. The first option is that most of the density perturbations are generated by a second field (a curvaton~\\cite{curvaton}) with gaussian perturbations. A second possibility is that the system contains several gauge fields. In both cases nongaussianities can be reduced to a safely small value while the tensor modes maintain an amplitude large enough to allow detectability. The possibility of parity violation in the CMB was already considered in the past (see e.g.~\\cite{Lue:1998mq,chiralgrav}). The mechanisms analyzed in those papers, however, are based on the assumption of a parity violating term in the gravitational sector, whereas the present work relies only on parity violation in the matter sector of the theory. Note also that the magnitude of parity violation induced by a gravitational Chern Simons term is limited to small values, according to~\\cite{cherneft}, once one requires validity of the field theoretical description of the system. The paper is organized as follows. Section 2 contains a review of the production of helical gauge modes by an axion-like inflaton. The generation of parity-violating gravitational waves is discussed in section 3. Section 4 deals with the constraints on the parameter space and describes two scenarios, consistent with the current constraints, where parity violating correlation functions would be detectable in future CMB surveys. We conclude in section 5. ", "conclusions": "% The main result of this paper is given by eqs.~(\\ref{main}), that show that a pseudoscalar inflaton, through its natural coupling to gauge fields, can induce a parity-violating component in the spectrum of gravitational waves. The degree of chirality~(\\ref{deltachi}) depends exponentially on the quantity $\\sqrt{\\epsilon}\\,M_P/f$, implying the spectrum of primordial gravitational waves is chiral in a large portion of the parameter space. This scenario represents, to our knowledge, the first example where parity violation can be imprinted on the CMB without invoking new physics in the gravitational sector. Since the mechanism responsible for the generation of chiral tensor modes does also generate an intrinsically non-gaussian component of the scalar perturbations, the simplest version of this scenario cannot produce observable parity violation without exceeding the bounds imposed by non-observation of nongaussianities~\\cite{Barnaby:2010vf}. In section 4, we have however presented two systems where such bounds do not apply, and that provide a proof of the existence of consistent models, compatible with current observations, that might lead to observable violation of parity in the CMB. A few additional comments are in order. In this scenario there is no simple proportionality between the value of the tensor-to-scalar ratio $r$ and the Hubble parameter during inflation. In particular, one can see from figure 1 that, for $\\xi\\gtrsim 3$, $r$ could be observable even for very small values of the Hubble parameter. Also, it is worth noting that parity-violating correlation functions in the CMB could also emerge as an effect of birefringence due to a pseudoscalar quintessence field~\\cite{Carroll:1998zi}. Ref.~\\cite{Gluscevic:2010vv}, however, has shown that it would be possible to discriminate parity-violating correlation functions of primordial origin from those induced by dynamics after the last scattering. Finally, it would be interesting to study whether the system~(\\ref{lag}) can generate a parity-odd component in the CMB bispectrum~\\cite{Kamionkowski:2010rb}. \\smallskip {\\bf Acknowledgments.} It is a pleasure to thank John Donoghue, Nemanja Kaloper, Mikhail Voloshin and especially Marco Peloso for useful discussions. This work is partially supported by the U.S. National Science Foundation grant PHY-0555304." }, "1101/1101.5240_arXiv.txt": { "abstract": "With the aim of studying active region fan loops using observations from the {\\it Hinode} EUV Imaging Spectrometer (EIS) and {\\it Solar Dynamics Observatory} (SDO) Atmospheric Imaging Assembly (AIA), we investigate a number of inconsistencies in modeling the absolute intensities of \\ion{Fe}{8} and \\ion{Si}{7} lines, and address why spectroheliograms formed from these lines look very similar despite the fact that ionization equilibrium calculations suggest that they have significantly different formation temperatures: $\\log\\,(T_e/K)$ = 5.6 and 5.8, respectively. These issues are important to resolve because confidence has been undermined in their use for differential emission measure (DEM) analysis, and \\ion{Fe}{8} is the main contributor to the AIA 131\\,\\AA\\, channel at low temperatures. Furthermore, the strong \\ion{Fe}{8} 185.213\\,\\AA\\, and \\ion{Si}{7} 275.368\\,\\AA\\, lines are the best EIS lines to use for velocity studies in the transition region, and for assigning the correct temperature to velocity measurements in the fans. We find that the \\ion{Fe}{8} 185.213\\,\\AA\\, line is particularly sensitive to the slope of the DEM, leading to disproportionate changes in its effective formation temperature. If the DEM has a steep gradient in the $\\log\\,(T_e/K)$ = 5.6 to 5.8 temperature range, or is strongly peaked, \\ion{Fe}{8} 185.213\\,\\AA\\, and \\ion{Si}{7} 275.368\\,\\AA\\, will be formed at the same temperature. We show that this effect explains the similarity of these images in the fans. Furthermore, we show that the most recent ionization balance compilations resolve the discrepancies in absolute intensities. With these difficulties overcome, we combine EIS and AIA data to determine the temperature structure of a number of fan loops and find that they have peak temperatures of 0.8--1.2MK. The EIS data indicate that the temperature distribution has a finite (but narrow) width $<$ $\\log\\, (\\sigma_{T_e}/K)$ = 5.5 which, in one detailed case, is found to broaden substantially towards the loop base. AIA and EIS yield similar results on the temperature, emission measure magnitude, and thermal distribution in the fans, though sometimes the AIA data suggest a relatively larger thermal width. The result is that both the \\ion{Fe}{8} 185.213\\,\\AA\\, and \\ion{Si}{7} 275.368\\,\\AA\\, lines are formed at $\\log\\, (T_e/K) \\sim$ 5.9 in the fans, and the AIA 131\\,\\AA\\, response also shifts to this temperature. ", "introduction": "To understand how the solar corona is heated to high temperatures, it is important to explain the heating of closed field structures such as active region loops. There have been extensive studies of these structures, and recent progress and outstanding issues have been reviewed by \\citet{klimchuk_2006} and \\citet{reale_2010}. A key diagnostic of the heating of coronal loops is the differential emission measure (DEM) distribution, because many coronal heating models make specific predictions as to its form and shape. For example, nanoflare reconnection models \\citep{parker_1983,parker_1988} predict the presence of a weak high temperature component in the DEM \\citep{cargill_1995}. There have been several recent studies that have tried to detect this emission \\citep{schmelz_etal2009a,reale_etal2009a,testa_etal2010}. The DEM gradient, or proportion of hot and cool material, also sets constraints on impulsive, steady or quasi-steady loop heating models \\citep{warren_etal2010c,tripathi_etal2010}, and the nanoflare model also predicts a spread in the temperature distribution within a loop due to the incoherent heating and cooling of unresolved threads. There has been considerable debate as to whether loops have a multi-thermal temperature distribution because of conflicting measurements by different instruments \\citep{lenz_etal1999,aschwanden_etal1999,schmelz_etal2001}. The discussion hinges on issues such as background subtraction \\citep{delzanna&mason_2003}, methodology \\citep{aschwanden_2002}, or the spatial resolution \\citep{aschwanden_etal2008}, or temperature resolution of the instruments used \\citep{martens_etal2002}. Recent observations by the {\\it Hinode} \\citep{kosugi_etal2007} EUV Imaging Spectrometer \\citep[][EIS]{culhane_etal2007b} suggest that loops formed near 1MK have a narrow temperature distribution, but are not isothermal \\citep{warren_etal2008a}. Doppler velocity measurements are another important diagnostic of the heating process, and again there have been numerous recent studies of flows in active region loops using EIS data \\citep{doschek_etal2007,hara_etal2008,delzanna_2008,brooks&warren_2009}. Some of the signatures of coronal heating models are expected to be quite subtle, however. The nanoflare reconnection model predicts weak downflows at `warm' temperatures \\citep{patsourakos&klimchuk_2006}, and short-lived faint upflows at high temperatures \\citep{patsourakos&klimchuk_2009}. The signatures of nanoflare heating in spectral line profiles therefore depend sensitively on the temperature of the flows, and this also implies that accurate measurements will allow inference of the properties of the energy release. It is important therefore to assign the flow temperatures as accurately as possible. \\begin{figure*} \\centering \\includegraphics[width=0.95\\linewidth]{f1.ps} \\caption{{\\it Hinode} EIS rasters of AR 10978 showing the similarity of \\ion{Fe}{8} and \\ion{Si}{7} images. The \\ion{Mg}{6} and \\ion{Fe}{9} images look to be formed at lower and higher temperatures, respectively. For example, note the increasing vertical extent of the fan structures and cool features to east and west of the active region as you look at the images from top to bottom of the third column (\\ion{Mg}{6}, \\ion{Si}{7}, \\ion{Fe}{9}). The images are scaled linearly. \\label{fig1}} \\end{figure*} A class of loops that have not yet been studied in sufficient detail with the latest instrumentation are the fan structures that appear as partially observed long loops at the edges of active regions. They are seen mostly in the 0.4--1.3MK temperature range \\citep{schrijver_etal1999,delzanna&mason_2003,ugarteurra_etal2009} and have densities greater than $\\log\\, (N_e/cm^{-3})$ = 9 \\citep{delzanna_2003,young_etal2007b}. They also appear to show red-shifted downflows \\citep{winebarger_etal2002,marsch_etal2004}. We examine a sample of these structures in this paper. For this and all the above observational studies, spectral emission from ions of Fe is of major importance. The high elemental abundance of Fe leads to many strong emission lines in spectrometer and imager pass-bands, yet interpretation of the observations is often difficult because of uncertainties in the atomic data used in prediction of emission \\citep{lanzafame_etal2002,young_etal2009}. In principle, the emission from spectral lines of Fe can provide stringent constraints on temperatures and densities in the corona, however, it is of paramount importance that the diagnostic capabilities of Fe lines be assessed critically. Initial results from EIS have indicated a number of problems in interpreting the Fe emission. \\citet{young_etal2007b} noted that spectroheliograms of \\ion{Fe}{8} and \\ion{Si}{7} look very similar in specific active region features such as the fan loops, despite the fact that the temperatures of the peak fractional abundance in ionization equilibrium are significantly different; $\\log\\, (T_e/K)$ = 5.6 and 5.8, respectively, according to \\citet{mazzotta_etal1998}. Since Fe is a considerably more complex atom than Si and is therefore likely to be less well understood, this observation has led to the suggestion that the ionization balance calculations for \\ion{Fe}{8} need to be revised upward to higher temperatures \\citep{young_etal2007b}. Furthermore, this could have an impact on surrounding ions such as \\ion{Fe}{7} and \\ion{Fe}{9}. Recently, \\citet{young&landi_2009} found evidence that this is indeed the case for \\ion{Fe}{7}. In addition to this temperature problem, a number of previous differential emission measure (DEM) studies have found difficulties reproducing the absolute intensities of \\ion{Fe}{8} and \\ion{Si}{7} lines simultaneously. See, for example, the quiet Sun off-limb DEM analysis of \\citet{warren&brooks_2009}, the on disk quiet Sun study of \\citet{brooks_etal2009}, or the analysis of a cool active region feature by \\citet{landi&young_2009}. Such issues are important to resolve as the \\ion{Fe}{8} 185.213\\,\\AA\\, and \\ion{Si}{7} 275.368\\,\\AA\\, lines would provide the best EIS lower temperature constraints on the DEM if we had good confidence in the atomic data. Also, since they are strong, they are the best lines to use for velocity measurements in the transition region, and are present in the majority of EIS observations. As discussed, it is therefore of critical importance that we have confidence in the temperatures we assign to the measured velocities. Changes to the formation temperatures or atomic data for these ions could also have an impact on interpreting the results from imagers. For example, changes to \\ion{Fe}{8} could affect the response functions for the 131\\,\\AA\\, channel of the Atmospheric Imaging Assembly (AIA) on the {\\it Solar Dynamics Observatory (SDO)}. Motivated by our interest in studying the temperature and velocity structure of the fan loops, this situation has led us to take a closer look at the formation of the \\ion{Fe}{8} and \\ion{Si}{7} spectral lines. Recently, there have been a number of revisions to the ionization balance calculations \\citep{bryans_etal2009,dere_etal2009} and we investigate whether they could help resolve these inconsistencies. We also examine the details of spectral line formation in this temperature range in the quiet Sun and in the fan loops. In particular, we examine whether convolving the contribution functions with the temperature distribution of the feature could explain the similarity of \\ion{Fe}{8} and \\ion{Si}{7} images. It is known that the temperature of peak contribution to the line intensity can be shifted from the theoretical peak temperature of the emissivity if the shape of the DEM is taken into consideration \\citep{brosius_etal1996,feldman_etal1999,delzanna_etal2003}. In doing so, we finally show that the ionization equilibrium calculations for Fe may not be the source of this problem. By determining more realistic effective formation temperatures, we present a possible explanation for the observations. We also show that the most recent ionization balance calculations can resolve the discrepancies in the magnitudes of the intensities found in previous studies. With confidence in the atomic data for these ions restored, we perform an emission measure (EM) analysis of a number of fan loops using EIS and AIA and compare the results from the two instruments. ", "conclusions": "Motivated by the desire to study the temperature structure of active region fan loops we have attempted to resolve inconsistencies found in previous work using EIS data. In particular, we have shown that the similarity in EIS \\ion{Fe}{8} 185.213\\,\\AA\\, and \\ion{Si}{7} 275.368\\,\\AA\\, images, that is not expected from the respective temperatures of peak abundance in ionization equilibrium, can be understood when a more accurate calculation of the effective formation temperature in the solar corona is performed. This is done by convolving the contribution functions with the DEM of the target of interest. If the DEM has a steep gradient in the $\\log\\, (T_e/K)$ = 5.6--5.8 range, or is sharply peaked, the two lines will be formed close in temperature. In this work, we compared the effect of this technique on the formation temperatures of these lines in the quiet Sun. The initial separation of $\\log\\, (T_e/K)$ = 5.6--5.8 is reduced to 5.75--5.85, and it is clear that \\ion{Fe}{8} 185.213\\,\\AA\\, has a substantial contribution from emitting material at $\\log\\, (T_e/K)$ = 5.8. To examine whether this explanation could work for the fan loops, we derived the EM distribution along one example in AR 10978. The temperature distribution peaks near 1--1.2MK and narrows along the loop. The peak contribution to the line intensity is $\\log\\, (T_e/K)$ = 5.9 for both the \\ion{Fe}{8} 185.213\\,\\AA\\, and \\ion{Si}{7} 275.368\\,\\AA\\, lines. To investigate whether this effect is generally applicable to other AR fans, we examined a number of other loops. We found that in all cases the two lines are formed at the same temperature ($\\log\\, T_e/K \\sim$ 5.9). This suggests therefore, that the expected difference in images of the fans formed from these lines is a result of an overestimation of the separation in formation temperatures by the approximate method of assuming the lines are formed at the temperature of the peak fractional abundance in ionization equilibrium. Note that other lines may be affected in similar ways. For example, in Figure \\ref{fig1} the \\ion{Mg}{7} and \\ion{Fe}{9} images look different despite the fact that they have similar ionization equilibrium temperatures. We have verified that they are in fact formed at different temperatures in the fan loop of Section \\ref{ftarcfs}. To demonstrate the importance of understanding the formation of the EUV spectrum for broad pass-band imagers we studied the effect of convolving the AIA 131\\,\\AA\\, response function with our EIS fan loop EM distributions. We showed that, as a result, the peak of the dominant lower temperature part of the effective response function shifts up to $\\log\\, (T_e/K) \\sim$ 5.9. It is important to emphasize that we have not shown that this explanation holds in general for all active areas or structures on the Sun. If the DEM slope is shallower (or flat) in the $\\log\\, (T_e/K)$ = 5.6--5.8 range then the \\ion{Fe}{8} 185.213\\,\\AA\\, and \\ion{Si}{7} 275.368\\,\\AA\\, lines should still be formed at a wider separation in temperatures and examples of significantly different images should be found. The apparent lack of such observations for any solar feature, however, provides a stringent constraint on the gradient of the DEM slope in this temperature range all over the Sun. This is consistent with other independent studies that show strong similarities in the shape of the DEM distribution in different areas of the quiet Sun \\citep{lanzafame_etal2005,brooks_etal2009,feldman_etal2009a}. The DEM-gradient resolves the inconsistencies in formation temperatures between the \\ion{Fe}{8} and \\ion{Si}{7} lines, but a related issue is that discrepancies have also been found in the magnitude of the intensities of these lines in previous DEM studies using EIS data. We showed here that these additional issues are resolved when the most recent ionization balance compilation data of \\citet{dere_etal2009} are used for the atomic calculations. From this analysis, therefore, no substantive relative error in the Fe ionization balance is indicated for the specific ions of \\ion{Fe}{8} and \\ion{Fe}{9}. This work has demonstrated that the strong \\ion{Fe}{8} 185.213\\,\\AA\\, and \\ion{Si}{7} 275.368\\,\\AA\\, lines can be used with confidence for DEM studies and velocity work in the transition region. Therefore, we examined the temperature structure of a small sample of fan loops. We found that they have peak temperatures in the range 0.8--1.2MK. One loop was found to be isothermal, but more often the temperature distribution has a narrow width. This result is similar to that found by \\citet{warren_etal2008a} for `warm' active region loops. In one detailed case, the EM distribution is found to broaden considerably towards the base. This could have implications for the location of the heating. We also found that the peak temperatures and emission measures derived from AIA data are in agreement with those derived from EIS. There is also agreement on whether the loops are isothermal or not. The AIA analysis indicates a slightly larger thermal width than EIS when the loops are not isothermal. This is possibly because the EIS data contain observations from consecutive ionization stages of \\ion{Fe}{0} whereas the AIA data only sample every second ionization stage." }, "1101/1101.2526_arXiv.txt": { "abstract": "The computation of the self-force constitutes one of the main challenges for the construction of precise theoretical waveform templates in order to detect and analyze extreme-mass-ratio inspirals with the future space-based gravitational-wave observatory LISA. Since the number of templates required is quite high, it is important to develop fast algorithms both for the computation of the self-force and the production of waveforms. In this article we show how to tune a recent time-domain technique for the computation of the self-force, what we call the Particle without Particle scheme, in order to make it very precise and at the same time very efficient. We also extend this technique in order to allow for highly eccentric orbits. ", "introduction": "\\label{intro} Gravitational Wave Astronomy has the potential to unveil the secrets of physical phenomena not yet understood involving strong gravitational fields. In order to realize this potential there are observatories/detectors either operating or being design/constructed that cover a significant part of the relevant gravitational-wave frequency spectrum. From the very low band ($10^{-9}-10^{-6}$ Hz), where we have the pulsar timing arrays to the high frequency band ($1-10^{4}$ Hz), where most ground-based detectors operate. In low frequency band ($10^{-4}-1$ Hz), not accessible from the ground, we find the future Laser Interferometer Space Antenna (LISA)~\\cite{LISA}, an ESA-NASA mission consisting of three spacecrafts forming a triangular constelations in heliocentric orbit (see~\\cite{Danzmann:2003ad,Prince:2003aa} for details). The targets of LISA are: (i) massive black hole (BH) mergers; (ii) the capture and subsequent inspiral of stellar compact objects (SCO) into a massive BH sitting at a galactic center; (iii) galactic compact binaries; and (iv) stochastic backgrouns of diverse cosmological origin. For many of these systems it is crucial to have a priori gravitational waveform models with enough precision to extract the corresponding signals from the future LISA data stream, and also to estimate the physical parameters of the system with precision. Here we focus in the second type of LISA sources, namely the so-called extreme-mass-ratio inspirals (EMRIs). The EMRIs of interest for LISA consists of massive BHs with masses in the range $M= 10^4-10^7 M_{\\odot}$, and SCOs with masses in the range $m = 1-50 M_{\\odot}$. The SCO moves around the massive BH following highly relativistic motion well inside the strongest field region of the BH. The orbit is not exactly a geodesic of the BH because of the SCO own gravity, which influences its own motion, making the orbit to shrink until it plunges into the BH. During the last year before plunge, it has been estimated~\\cite{Finn:2000sy} that the SCO spends of the order of $10^{5}$ cycles (depending on the type of orbit) inside the LISA band. As a consequence the EMRI GWs carry a detailed map of the BH geometry. This will allow, in particular, to test the spacetime geometry of BHs and even alternative theories of gravity (see, e.g.~\\cite{Schutz:2009zz,Sopuerta:2010zy,Babak:2010ej}). We can also understand better the stellar dynamics near galactic nuclei, populations of stellar BHs, etc. (for a review see~\\cite{AmaroSeoane:2007aw}). All this requires very precise EMRI gravitational waveforms (the precision of the phase should be of the order of one radian per year). Due to the extreme mass ratio of these systems, $\\mu=m/M \\sim 10^{-7} -10^{-3}$, we can describe them in the framework of BH perturbation theory, where the SCO is modeled as a point-like mass and the backreaction effects are described as the effect of a local force, the so-called {\\em self-force}. This self-force is essentially determined by the derivatives of the metric perturbations (retarded field) at the particle location, which need to be regularized due to the singularities introduced by the particle description. The calculation of the retarded field has to be done numerically, either in the frequency or time domain. In this work we concentrate on a recent approach to self-force calculations in the time domain introduced in~\\cite{Canizares:2008dp,Canizares:2009ay} for circular orbits and extended to eccentric orbits in~\\cite{Canizares:2010yx}. This is a multidomain approach in which the point particle is always located at a node between two domains, and hence it has been named the Particle without Particle (PwP) scheme. The main advantage is that the equations to be solved are homogeneous and the particle location enters via junction/matching conditions. In this way we are dealing with smooth solutions at each domain, which is crucial for the convergence of the numerical method used to implement the method, the PseudoSpectral Collocation (PSC) method in our case. In this article we describe how to tune the numerical implementation of the PwP scheme to achieve high precision results with modest computational resources. We focus on two fronts: (i) How to change the framework introduced in~\\cite{Canizares:2009ay,Canizares:2010yx} to compute the self-force for orbits with high eccentricity, and (ii) how to pick the size of the computational domain in order to make our computations much more efficient, so that we can achieve very precise results with a modest amount of computational resources. A similar analysis for a different computational technique has been presented in~\\cite{Thornburg:2010tq}. The plan of the paper is the following: In Sec.~\\ref{scalarparticlearoundsch} we describe briefly the foundations of self-force calculations in a simplified model based on a charged particle orbiting a non-rotating BH. In Sec.~\\ref{pwpscheme} we summarize the PwP scheme and extend the multidomain structure in order to allow for computations in the case of high eccentric orbits. We also show the convergence properties of the PSC method numerical implementation. In Sec.~\\ref{fitmod} we describe how to choose the resolution depending on the mode in order to perform optimal calculations in the sense of computational resources. We finish with conclusions and a discussion in Sec.~\\ref{discussion}, where we also present some numerical results. ", "conclusions": "\\label{discussion} The PwP scheme introduced and developed in~\\cite{Canizares:2008dp,Canizares:2009ay,Canizares:2010yx} provides a framework for precise and efficient computations of the self-force in the time domain. In this paper we have investigated how to tune the method in order to increase its efficiency and accuracy. We have also extended the PwP scheme in order to make it suitable for computations of the self-force on highly eccentric orbits. We have studied how to distribute the domain sizes and the number of collocation points so that we allocate the optimum resolution to each harmonic mode. This is very important as the resolution requirement increases with $m$, despite the fact that high $\\ell$-modes contribute less to the self-force than low $\\ell$ ones. This means that in order to improve the accuracy of the self-force we must have a good control of the resolution, otherwise modes with high $m$ (and hence with high $\\ell$) will limit the precision despite not contributing much to the self-force. Using this information, we have conducted a series of computations of the self-force for the case of a scalar charged particle in circular (geodesic) orbits around a non-rotating BH. The computational parameter space that we have explored is described as follows: The total number of domains that we have used is in the range $d=20-43$. The coordinate size of the large domains (the ones far from the particle and near the boundaries) is in the range $\\Delta\\rsu =50 - 100M$. The number of collocation points has been fixed to $N = 50$. The largest $\\ell$, $\\ell_{\\MAX}$, considered in the computations is in the range $\\ell_{\\MAX}=20-40$. Adapting the resolution for each $m$-mode also allows us to adapt the time step as the CFL condition is proportional to the minimum coordinate physical distance (as measured in terms of the tortoise radial coordinate) between grid points (and inversely proportional to the square of the number of collocation points). This means that the number of time steps required, $N_{t}$, for a simulation is going to be proportional to $m$ as: $N_{t}^m = m\\,N_{t}^{m=1}$. This is assuming that we use the same number of domains for all modes, but it is clear that modes with low $m$ would need less domains than modes with high $m$, and therefore, this is another source of reduction of computational time. We have been able to obtain very precise values of the self-force for a wide range of values within the parameter space. For instance, in the case of the radial component of the regular field, $\\Phi^{\\RR}_r$, the only one that for the circular case needs regularization, we have obtained values like $\\Phi^{\\RR}_r = 1.677282\\times10^{-4}\\, q/M^2\\,,$ which have a relative error of the order of $5\\cdot10^{-5}\\,\\%$ with respect to the values obtained in~\\cite{DiazRivera:2004ik} using frequency-domain methods. The relative error is always in the range $5\\cdot10^{-5}\\,\\% - 5\\cdot10^{-3}\\,\\%$. This constitutes a significant improvement with respect to our own previous estimations presented in~\\cite{Canizares:2009ay}, where the relative errors quoted for $\\Phi^{\\RR}_r$ were of the order of $0.1\\%$. The typical time for a full self-force calculation in a computer with two Quad-Core Intel Xeon processors at $2.27$ GHz is in the range $10-15$ minutes, which is a very significant reduction with respect to our computations in~\\cite{Canizares:2009ay}, specially taking into account that we have also improved the precision of the self-force computations. The calculations we have presented can be further improved in terms of computational time, and perhaps in accuracy by exploring techniques to bring the boundaries closer to the particle without degrading the accuracy of the field values near it. This can be done either by improving the outgoing boundary conditions (see, e.g.~\\cite{Lau:2004as}) or by using some sort of compactification of the physical domain (see, e.g.~\\cite{2010arXiv1008.3809Z}). Another possibility for making the computations faster (although this does not decrease the CPU time) is to parallelize the code and use computers with many cores. This is in principle a simple task as the different modes are not coupled. In any case, the next step in this line of work is to perform a similar phenomenological study for the eccentric case, where not only the azimuthal frequency is important, but also the radial one, which is absent in the circular case." }, "1101/1101.4716_arXiv.txt": { "abstract": "Electromagnetic fields of an accelerated charge are derived from the first principles using Coulomb's law and the relativistic transformations. The electric and magnetic fields are derived first for an instantaneous rest frame of the accelerated charge, without making explicit use of Gauss's law, an approach different from that available in the literature. Thereafter we calculate the electromagnetic fields for an accelerated charge having a non-relativistic motion. The expressions for these fields, supposedly accurate only to a first order in velocity $\\beta$, surprisingly yield all terms exactly for the acceleration fields, only missing a factor $1-\\beta^2$ in the velocity fields. The derivation explicitly shows the genesis of various terms in the field expressions, when expressed with respect to the time retarded position of the charge. A straightforward transformation from the instantaneous rest frame, using relativistic Doppler factors, yields expressions of the electromagnetic fields for the charge moving with an arbitrary velocity. The field expressions are derived without using Li\\'{e}nard-Wiechert potentials, thereby avoiding evaluation of any spatial or temporal derivatives of these potentials at the retarded time. ", "introduction": "The electromagnetic (EM) fields of a moving charge are formally calculated from Li\\'{e}nard-Wiechert potentials.\\cite{1,2,3,12} The mathematics, quite tricky and involved since derivatives of the potentials need to be evaluated at the retarded time, can be quite a deterrent to the initiate. The expression for the fields in the instantaneous rest frame of an accelerated charge can be worked out from physical arguments.\\cite{9,4,5,8} It is generally believed that a relativistic transformation of fields from the instantaneous rest frame to an inertial frame in which charge has an arbitrary velocity could be quite laborious.\\cite{7} EM fields for the special case of acceleration parallel to the velocity vector have been derived\\cite{10,11} without using Li\\'{e}nard-Wiechert potentials. Huang and Lu \\cite{6} attempted the more general case but ended up with wrong expressions for the EM fields though they claimed to be giving ``the exact expression'' for radiation of an accelerated charge. To bypass the mathematical cumbrousness, Padmanabhan\\cite{7} used an alternate approach where the expressions for EM field were derived in an indirect manner, by finding a general covariant 4-vector function of position, velocity and acceleration of the charge which coincided with the EM field values in the instantaneous rest frame. The approach though elegant is not immediately obvious. It may thus be still worthwhile to have EM field expressions transformed directly from the rest-frame values using a 3-vector language, which should be transparent to the reader. We show here that the standard text-book expressions for the EM fields of an accelerated charge can be derived in a fairly easy and straightforward manner, without using Li\\'{e}nard-Wiechert potentials, thereby avoiding evaluation of any spatial or temporal derivatives of these potentials at the retarded time. We start with the radial Coulomb field of a stationary charge and then making use of the relativistic transformations, in particular that of EM fields, we derive the fields for an accelerated charge in its instantaneous rest frame where transverse field components proportional to acceleration show up. The presence of such an electric field component proportional to acceleration was shown by Thomson,\\cite{9} using a physical picture in the pre-relativity days, employing the concept of electric field lines representing Faraday (flux) tubes. A modern derivation using Gauss's law is now available in many text-books.\\cite{4,5,8} We shall derive the transverse components for both electric and magnetic fields, in the same spirit but from a different perspective, without explicitly using Gauss's law. Thereafter we get field expressions for a slowly moving charge using non-relativistic transformations and surprisingly the expression for the acceleration fields turn out to be exactly the same as for a relativistically moving charge, which we derive rigorously in a later section. We have used Gaussian system of units throughout. ", "conclusions": "We have derived the EM fields of a charge with an arbitrary motion. The expressions for both the electric and magnetic fields were first derived for a charge in its instantaneous rest frame using a physical picture. Thereafter we calculate the EM fields of an accelerated charge having a non-relativistic motion. The expressions for these fields, accurate to first order in velocity $\\beta$, when expressed with respect to the time retarded position of the charge, surprisingly yield all terms exactly for the acceleration fields. A Lorentz transformation from the instantaneous rest frame, using relativistic Doppler factors, then led us to the standard expressions for the electromagnetic fields of an arbitrarily moving charge without using Li\\'{e}nard-Wiechert potentials and without resorting to any differentiation that need to be evaluated at the retarded time. The derived expressions of course agree with those derived from Li\\'{e}nard-Wiechert potentials." }, "1101/1101.0842_arXiv.txt": { "abstract": "We report on the quiescent X-ray properties of the recently discovered transiently accreting 11 Hz X-ray pulsar in the globular cluster Terzan 5. Using two archival \\chan\\ observations, we demonstrate that the quiescent spectrum of this neutron star low-mass X-ray binary is soft and can be fit to a neutron star atmosphere model with a temperature of $kT^{\\infty} \\sim 73$~eV. A powerlaw spectral component is not required by the data and contributes at most $\\sim20\\%$ to the total unabsorbed 0.5--10 keV flux of $\\sim 9 \\times 10^{-14}~\\flux$. Such a soft quiescent spectrum is unusual for neutron stars with relatively high inferred magnetic fields and casts a different light on the interpretation of the hard spectral component, which is often attributed to magnetic field effects. For a distance of $5.5$~kpc, the estimated quiescent thermal bolometric luminosity is $\\sim 6\\times10^{32}~\\lum$. If the thermal emission is interpreted as cooling of the neutron star, the observed luminosity requires that the system is quiescent for at least $\\sim100$ years. Alternatively, enhanced neutrino emissions can cool the neutron star to the observed quiescent luminosity. ", "introduction": "\\label{sec:intro} Low-mass X-ray binaries (LMXBs) are binary star systems in which a neutron star or a black hole accretes matter from a (sub-) solar companion. Two phenomena are thought to uniquely identify the compact primary as a neutron star: coherent X-ray pulsations and type-I X-ray bursts. The latter are intense flashes of X-ray emission caused by thermonuclear runaway of the accreted matter on the surface of the neutron star. X-ray pulsations can be observed when the neutron star magnetic field is strong enough to funnel the accretion flow to the magnetic poles. Many neutron star LMXBs are transient and spend most of their lifetime in a quiescent state, during which they are dim with typical 0.5--10 keV X-ray luminosities of $L_q \\sim 10^{31-33}~\\lum$ \\citep[e.g.,][]{heinke2009}. However, occasionally they exhibit outbursts during which their X-ray luminosity increases orders of magnitude to $L_X \\sim 10^{36-38}~\\lum$ \\citep[2--10 keV; e.g.,][]{chen97}. The enhanced activity is ascribed to a sudden strong increase in the mass-accretion rate onto the neutron star, whereas little or no matter is accreted during quiescent episodes. The spectra of quiescent neutron star LMXBs can typically be fitted with a soft thermal model, a hard powerlaw shape, or a combination of both. The hard spectral component has been attributed to non-thermal emission processes related to the magnetic field of the neutron star, e.g., accretion onto the magnetosphere or a pulsar wind mechanism \\citep[e.g.,][]{campana1998}. The soft thermal component is generally interpreted as thermal emission from the neutron star surface. During accretion outbursts, a chain of nuclear reactions release heat deep in the neutron star crust \\citep[e.g.,][]{haensel2008}, which is re-radiated during quiescent episodes. This results in an incandescent luminosity that is set by the long-term averaged mass-accretion rate of the system and the efficiency of neutrino emission processes occurring in the neutron star core \\citep[e.g.,][]{brown1998}. Residual accretion onto the neutron star surface offers an alternative explanation for the quiescent thermal emission \\citep[][]{zampieri1995}. Galactic globular clusters are rich targets for studies of X-ray binaries, which can form in dynamical interactions in the dense cluster environments \\citep[e.g.,][]{pooley2003}. Numerous low-luminosity X-ray sources have been found in globular clusters, amongst which are several candidate quiescent LMXBs \\citep[e.g.,][]{verbunt1995_rosat, grindlay2001,heinke2006_terzan5}. Indeed, a few transient LMXBs have been identified in globular clusters during outburst episodes \\citep[see e.g., the review by][]{verbunt2006}.% \\begin{figure} \\begin{center} \\includegraphics[width=8.0cm]{asm_5day_avg.eps} \\end{center} \\caption[]{{\\rxte/ASM 5-day averaged lightcurve of Terzan 5, showing activity from the globular cluster in 2000, 2003 and 2010.}} \\label{fig:asm} \\end{figure} \\subsection{Terzan 5}\\label{subsec:terzan5} The globular cluster Terzan 5 has long been known to harbour at least one transient neutron star LMXB. In 1980, a number of X-ray bursts were observed with the {\\it Hakucho} satellite \\citep[][]{makishima1981}. Activity from Terzan 5 was also detected in 1984 with \\exosat\\ \\citep[][]{warwick1988}, and in 1990/1991 with \\rosat\\ \\citep[][]{verbunt1995_rosat,johnston1995}. Furthermore, \\rxte/ASM observations have revealed three distinct X-ray outbursts since 1996: in 2000 \\citep[][]{markwardt2000,heinke2003}, 2002 \\citep[][]{wijnands2002_terzan5} and 2010 (see Figure~\\ref{fig:asm}). The limited angular resolution of these instruments precludes pinpointing the X-ray source that causes the outbursts. However, \\chan\\ observations obtained during the 2000 outburst of Terzan 5 allowed for an accurate localization of the transient source that was active at that time \\citep[][]{heinke2003}. Multiple X-ray bursts were observed and the 2000 transient was therefore associated with the X-ray burster that was detected by earlier X-ray missions (\\hakuchoname/\\exo). The quiescent counterpart of this transient neutron star LMXB has an unusually hard spectrum, with no clear evidence for the presence of a thermal emission component \\citep[][]{wijnands2005}. A \\chan\\ study revealed $50$ distinct X-ray point sources within the half-mass radius of the cluster, several of which were proposed to be quiescent LMXBs \\citep[][]{heinke2006_terzan5}. Renewed activity from Terzan 5 was detected in 2010 October during \\inte\\ bulge scan monitoring observations \\citep[][]{bordas2010}. Subsequent pointed \\rxte/PCA observations detected coherent 11 Hz pulsations and type-I X-ray bursts, establishing the nature of the transient X-ray source as an accreting neutron star \\citep[][]{strohmayer2010}. Timing studies of the X-ray pulsations provided a determination of the binary orbital period ($P_{\\mathrm{orb}} = 21.27$~h), and the identification of the mass donor as a $\\sim 0.4 - 1.5 ~\\Msun$ main sequence or slightly evolved star \\citep[][]{papitto2010}. An accurate \\swift\\ localization revealed that the source active in 2010 was likely a different transient than the one detected in 2000 \\citep[][]{kennea2010}. \\chan\\ observations confirmed this, thereby establishing the existence of a second transient neutron star LMXB in Terzan 5, named \\source\\ \\citep[][]{pooley2010}. The angular separation between the two transients is only $\\sim5.5''$ (see Figure~\\ref{fig:ds9}), and due to the limited angular resolution of older X-ray missions it is unclear which of the two was responsible for other active episodes of Terzan 5. The 2010 transient corresponds to the X-ray source CX25 from the \\chan\\ study of \\citet{heinke2006_terzan5}, which was marked as a candidate quiescent LMXB by these authors. ", "conclusions": "\\label{sec:discuss} We report on the spectral analysis of the newly discovered 11 Hz X-ray pulsar \\source\\ (\\intename) in the globular cluster Terzan 5 during quiescence. Using two archival \\chan\\ observations, carried out in 2003 and 2009, we show that the quiescent spectrum is dominated by thermal emission that fits to a neutron star atmosphere model with a temperature of $kT^{\\infty} \\sim 73$~eV. The inferred thermal bolometric luminosity is $L_q \\sim 6 \\times 10^{32}~(D/5.5~\\mathrm{kpc})^2~\\flux$. If a powerlaw is included in the fits, this spectral component contributes $\\lesssim20\\%$ to the total unabsorbed 0.5--10 keV luminosity of $\\sim 3 \\times 10^{32}~(D/5.5~\\mathrm{kpc})^2~\\flux$. There is no evidence for spectral variations in the quiescent emission between the 2003 and 2009 observations. \\subsection{The powerlaw spectral component}\\label{subsec:powerlaw} It is interesting to compare the quiescent spectral properties of the 2010 Terzan 5 transient with results obtained for the accreting millisecond X-ray pulsars (AMXPs). The quiescent spectra of five AMXPs could be studied in detail: \\ngcbron, Aql X-1, \\sax, \\xtepulsar\\ and \\igrpulsar. The former two are both intermittent X-ray pulsars (i.e., the pulsations are detected only sporadically during outburst), that have predominantly soft quiescent X-ray spectra and a 0.5--10 keV luminosity of $\\sim 10^{33}~\\lum$ \\citep[][]{rutledge2001,cackett2005}. \\sax, \\xtepulsar\\ and \\igrpulsar, have quiescent 0.5--10 keV luminosities on the order of $\\sim 10^{32}~\\lum$. Of the three, only \\igrpulsar\\ shows evidence for thermal emission, contributing $\\sim 40\\%$ to the total 0.5--10 keV flux \\citep[][]{heinke2009}, while the quiescent spectra of the other two are completely dominated by a hard spectral component \\citep[][]{campana2002,wijnands05_amxps}. Our limited understanding of the origin of hard quiescent X-ray emission makes it difficult to interpret the difference between the 2010 Terzan 5 transient and other X-ray pulsars. However, if the powerlaw spectral component is related to the magnetic field of the neutron star \\citep[][]{campana1998}, we would have expected to detect a hard X-ray spectrum for the 11 Hz pulsar in Terzan 5, since this source must have a substantial magnetic field \\citep[][Cavecchi et al. in prep.]{papitto2010}. A possible explanation for the soft X-ray spectrum of the 2010 Terzan 5 transient is that little matter is available in quiescence to interact with the neutron star magnetic field, so that the powerlaw emission component that may arise from this process is strongly reduced. The effects of the strong magnetic field on the neutron star spectrum might also play a role \\citep[e.g.,][]{zavlin2002}. \\begin{figure} \\begin{center} \\includegraphics[width=8.0cm]{newtransient_multispec_revised.eps} \\end{center} \\caption[]{{\\chan/ACIS quiescent spectra of the 11 Hz X-ray pulsar in Terzan 5 from 2003 (black) and 2009 (grey) data, along with a neutron star atmosphere model fit (solid lines).}} \\label{fig:spec} \\end{figure} \\subsection{The quiescent thermal emission}\\label{subsec:thermal} The observed thermal quiescent luminosity of the 11 Hz pulsar in Terzan 5 can be used to estimate the duty cycle of the system, by equating this value to the luminosity that is expected to be radiated due to heating of the neutron star. Assuming that the observed accretion luminosity provides a measure for the amount of matter that is accreted onto the neutron star, we can deduce the mass-accretion rate during outburst from the average outburst flux. \\citet{papitto2010} infer a 0.1--100 keV unabsorbed flux of $\\sim 1 \\times 10^{-8}~\\flux$ from fitting \\rxte/PCA and HEXTE spectral data obtained during the 2010 outburst of Terzan 5. This yields an estimated bolometric accretion luminosity of $L_{acc} \\sim 4\\times10^{37}~(D/5.5~\\mathrm{kpc})^2~\\lum$. For an accretion luminosity given by $L_{acc} = (G M_{NS}/R_{NS}) \\langle \\dot{M}_{ob} \\rangle$, the average accretion rate during outburst is $\\langle \\dot{M}_{ob} \\rangle \\sim 2 \\times 10^{17}~\\mathrm{g~s}^{-1} \\sim 3\\times10^{-9}~\\mdot$, for canonical neutron star parameters of $M_{NS} = 1.4~\\Msun$ and $R_{NS} = 10$~km. The nuclear reactions induced by the accretion of matter are expected to result in a quiescent bolometric luminosity that is given by $L_q = \\langle \\dot{M} \\rangle Q_{nuc} / m_u$ \\citep[e.g.,][]{brown1998,colpi2001}. Here, $Q_{nuc} \\sim 2$ MeV is nuclear energy deposited in the crust per accreted baryon \\citep[e.g.,][]{haensel2008}, $m_u=1.66\\times10^{-24}$~g is the atomic mass unit and $\\langle \\dot{M} \\rangle = \\langle \\dot{M}_{ob} \\rangle \\times t_{ob} / t_{rec}$ is the time-averaged mass-accretion rate of the system. The ratio of the outburst duration ($t_{ob}$) and the recurrence time ($t_{rec}$) represents the duty cycle. Using the above equation, the observed quiescent thermal bolometric luminosity of $L_q \\sim 6 \\times 10^{32}~(D/5.5~\\mathrm{kpc})^2~\\lum$ suggests a time-averaged mass-accretion rate of $\\langle \\dot{M} \\rangle \\sim 5\\times10^{-12}~\\mdot$. Combined with the estimated mass-accretion rate during the 2010 outburst, this would imply that the system must have a duty cycle on the order of $\\sim 0.1 \\%$, provided that standard cooling mechanisms are operating in the core. The 2010 outburst commenced around October 10 and is ongoing during \\rxte/PCA pointed observations performed on November 19.\\footnote{Terzan 5 is unobservable with \\rxte\\ between 2010 November 19 and 2011 January 17 due to Sun-angle constraints, nor can it be observed with other X-ray satellites during that epoch.} The outburst thus has a duration of $>6$~weeks, so the estimated duty cycle would imply that the source must spend $\\gtrsim100$~yr in quiescence. This increases further if the outburst is observed to continue for a longer time. Although it is plausible that the 2010 Terzan 5 transient has a very low duty cycle, an alternative explanation is that enhanced neutrino emission mechanisms are operating in the core, which cool the neutron star down to the observed quiescent luminosity. Within our current understanding, this would suggest that the neutron star in this X-ray binary is relatively massive \\citep[e.g.,][]{colpi2001}, or that the core composition includes different forms of matter \\citep[e.g., pions or kaons;][]{yakovlev2004}. It is not obvious that the neutron star in this transient LMXB would be particularly massive, since both its high inferred magnetic field \\citep[$B \\sim10^{9}-10^{10}$~G or possibly even higher;][Cavecchi et al. in prep.]{papitto2010} and relatively slow spin period ($P_s = 11$~Hz) are suggestive of a relatively young system. From modelling the quiescent spectral data we find a neutron star effective temperature, as observed by a distant observer, of $kT^{\\infty}\\sim73$~eV. In the neutron star frame this translates into a value of $kT\\sim95$~eV ($T \\sim 1 \\times 10^{6}$~K) for $M_{\\mathrm{NS}}=1.4~\\Msun$ and $R_{\\mathrm{NS}}=10$~km (corresponding to a gravitational redshift of $1+z=1.3$). Using the model calculations of \\citet{brown08} we can obtain a rough estimate of the neutron star core temperature of $\\sim3\\times10^{7}$~K. For standard (i.e., slow) core neutrino emission processes, such a core temperature yields a neutrino luminosity on the order of $L_{\\nu} \\sim 10^{28}~\\lum$ \\citep[cf. figure 2 of][]{schaab1999}. This is negligible compared to the observed photon emissions of $L_q \\sim 6 \\times 10^{32}~(D/5.5~\\mathrm{kpc})^2~\\lum$. However, for a sample enhanced cooling model (direct Urca), the neutrino luminosity is on the order of $L_{\\nu} \\sim10^{35}~\\lum$ \\citep[cf. figure 2 of][]{schaab1999}, and thus the dominant cooling mechanism. For such neutrino losses a duty cycle of roughly $\\sim25\\%$ is required to explain the observed quiescent thermal luminosity. The time-averaged mass-accretion rate for this scenario is $\\langle \\dot{M} \\rangle \\sim 8\\times10^{-10}~\\mdot$. The outburst duration of the new Terzan 5 transient is unconstrained, but for a typical value of $\\sim 2-12$~months, a duty cycle of $\\sim25\\%$ implies a recurrence time of $\\sim1-4$~yr. This would suggest that some of the previous outbursts from Terzan 5 might have originated from the 11 Hz X-ray pulsar. \\chan\\ observations leave no doubt that the 2000 outburst of Terzan 5 was caused by the other transient LMXB located in this cluster \\citep[cf.][]{heinke2003,pooley2010}. However, lack of high-spatial resolution observations during previous active periods of Terzan 5 (i.e., 1980, 1984, 1990, 1991 and 2002; see Section~\\ref{sec:intro}) do not rule out this possibility. Detailed studies of the X-ray burst behaviour can potentially shed more light on this. We note that the non-detection of quiescent thermal emission for the 2000 Terzan 5 transient led \\citet{wijnands2005} to conclude that this source must either spend hundreds of years in quiescence, or be subject to enhanced neutrino cooling. If both Terzan 5 LMXBs undergo slow core cooling and spend long episodes in quiescence, the other outbursts observed from this globular cluster must have been caused by another transient source. There are several candidate quiescent LMXBs identified in Terzan 5 \\citep[][]{heinke2006_terzan5}. Alternatively, at least one of the neutron star LMXBs undergoes enhanced neutrino core cooling. When inferring the quiescent bolometric luminosity of the 2010 Terzan 5 transient we assumed a distance of $D=5.5$~kpc, as inferred by \\citet{ortolani2007}. We note that \\citet{cohn2002} report a distance towards Terzan 5 of $D=8.7$~kpc, which would increase the thermal bolometric luminosity inferred in this work by a factor of $\\sim2$. However, for a larger distance the inferred accretion luminosity is also increased, so in practise this does not affect the estimated duty cycle. Furthermore, as discussed in \\citet{ortolani2007}, $D=5.5$~kpc can be considered a more reliable distance estimate. We note that if the observed quiescent thermal emission component is due to residual accretion onto the neutron star surface \\citep[][]{zampieri1995}, the interior temperature of the neutron star must be lower than we infer here, thus requiring a smaller duty cycle. The new transient source discovered in Terzan 5 became very bright during its 2010 outburst, reaching up to $L_X \\sim 10^{38}~\\lum$ \\citep[][]{altamirano2010_2}. Combined with the relatively low inferred neutron star temperature during quiescence, this makes the source a potential target to search for cooling of the neutron star crust, once the accretion ceases. \\citet{brown1998} argue that the neutron star crust can become significantly heated for systems in which the outburst luminosity is much higher than the quiescent level. Once the accretion ceases, the thermal relaxation of the neutron star crust might become observable as a gradual decrease in neutron star effective temperature. \\\\ \\noindent {\\bf Acknowledgements.}\\\\ This work was supported by the Netherlands Research School for Astronomy (NOVA) and made use of the \\chan\\ public data archive. RW acknowledges support from a European Research Council (ERC) starting grant. The authors are grateful to the anonymous referee for providing thoughtful comments that helped improve this manuscript. \\vspace{-0.5cm}" }, "1101/1101.5656_arXiv.txt": { "abstract": "There are mainly two different approaches to measure the cosmic star formation history: direct star formation rate density (SFRD) and stellar mass density $\\rhostar$ as functions of redshift. Compilations of current observations seem to show a disparity in the two quantities, in the sense that the integral of SFRD is higher than the observed $\\rhostar$ (after considering gas recycling). Using cosmological smoothed particle hydrodynamics simulations based on the concordance $\\Lambda$ cold dark matter model, we show that the two quantities become more consistent with each other when we consider the observed galaxy mass limit. The comparison between simulations and (dust corrected) observed cosmic SFRD shows a good agreement, while the observed $\\rhostar$ is significantly lower than the simulation results. This can be reconciled if the current high-$z$ galaxy surveys are missing faint low-mass galaxies due to their flux limit. Our simulated GSMFs have steep low-mass end slopes of $\\alpha \\lesssim -2$ at $z>3$, and when these numerous low-mass galaxies are included, the total $\\rhostar$ matches with the integral of SFRD. ", "introduction": "\\label{sec:intro} The cosmic star formation (SF) history is a fundamental quantity that illustrates the buildup of galaxies. It can be estimated by direct measurement of SFRD as a function of redshift through rest-frame H$\\alpha$, UV, and far-IR emission with dust extinction corrections \\citep[e.g.,][]{Hopkins.Beacom:06}. The stellar mass density ($\\rhostar$) at various redshifts is an indirect measurement of the cosmic SF history, because it is indeed the integral of SFRD, after considering the recycling of gas into interstellar medium. Given this simple relationship, in principle the measurements of the two quantities should be consistent with each other. However, in this paper we show that this is not the case for current observational estimates, and other authors have reached similar conclusions \\citep{Nagamine.etal:04, Ouchi.etal:04, vanDokkum:08, Wilkins.etal:08, Dave:08}. Over the past decade, large samples of color selected high-$z$ galaxies such as Lyman break galaxies (LBGs) have enabled us to measure the cosmic SFRD at $3 \\lesssim z \\lesssim 6$ from rest-frame UV luminosity density \\citep[e.g.,][]{Steidal.etal:99}. A number of near-infrared observations have also constrained the galaxy stellar mass function (GSMF) at $z\\lesssim 4$ \\citep[e.g.,][]{Marchesini.etal:09}. Furthermore, the advent of the {\\it Wide Field Camera 3} ({\\it WFC3}) on board the {\\it Hubble Space Telescope} ({\\it HST}) has dramatically improved our ability to measure the rest-frame UV light from galaxies at $z \\gtrsim 6$. Combined with the measurement of rest-frame optical light by the {\\it Infrared Array Camera} ({\\it IRAC}) on {\\it Spitzer Space Telescope}, it is now possible to measure the characteristic stellar mass of galaxies at $z \\gtrsim 6$ using stacked spectra \\citep{Yan.etal:09,Bouwens.etal:10,Labbe.etal:10,Gonzalez.etal:10}. \\begin{table*} \\begin{center} \\begin{tabular}{ccccccc} \\hline Name & Box-size & ${N_{\\rm p}}$ & $m_{\\rm DM}$ & $m_{\\rm gas}$ & $\\epsilon$ & $z_{\\rm end}$ \\\\ \\hline \\hline N216L10 & 10.0 & $2\\times 216^3$ & $5.96 \\times 10^6$ & $1.21 \\times 10^6$ & 1.85 & 2.75 \\cr \\hline N400L34 & 33.75 & $2\\times 400^3$ & $3.49 \\times 10^7$ & $7.31 \\times 10^6$ & 3.38 & 1.0 \\cr \\hline N600L100 & 100.0 & $2\\times 600^3$ & $2.70 \\times 10^8$ & $5.66 \\times 10^7$ & 4.30 & 0.0 \\cr \\hline \\label{table:sim} \\end{tabular} \\caption{ The three different resolution and volume simulations employed in this {\\it Letter}. The box-size is given in units of $h^{-1}$Mpc, ${N_{\\rm p}}$ is the particle number of dark matter and initial gas (hence $\\times\\, 2$), $m_{\\rm DM}$ and $m_{\\rm gas}$ are the masses of dark matter and initial gas particles in units of $\\himsun$, respectively, $\\epsilon$ is the comoving gravitational softening length in units of $h^{-1}$kpc, and $z_{\\rm end}$ is the ending redshift of the simulation. Note that the mass of star particle is a half of the initial gas particle. The value of $\\epsilon$ is a measure of spatial resolution. The name of each simulation is based on the particle count and its box size. } \\end{center} \\end{table*} However, the sources at $z>6$ are very faint, and so far we have only detected the massive end of GSMF. In addition, there are still significant uncertainties in the estimates of SFRD and $\\rhostar$. Given this situation, it would be useful to obtain predictions on SFRD and $\\rhostar$ from theoretical models. In particular, cosmological hydrodynamic simulations have been widely used to investigate cosmic star formation \\citep[e.g.,][]{Cen.Ostriker:92,Katz.etal:96,Springel.Hernquist:03_SFR,Nagamine.etal:06,Dave:08,Schaye.etal:10}. In this {\\it Letter}, we focus on the cosmic SF history at $z>2$ using cosmological simulations. The remaining of this paper is organized as follows. In Section~\\ref{sec:method}, we describe the cosmological hydrodynamic simulations. In Section~\\ref{sec:MF}, we construct the composite GSMF by combining the samples of galaxies from simulations with different resolution and volumes. We study the cosmic SFR and $\\rhostar$ evolution in Section~\\ref{sec:SF}. We summarise and discuss our findings in Section~\\ref{sec:summary}. ", "conclusions": "\\label{sec:summary} Based on the results of self-consistent cosmological hydrodynamic simulations, we argued that the current high-$z$ observations are missing low-mass galaxies with $\\Mstar \\lesssim 10^9\\Msun$ at $z>4$ when accounting for the total $\\rhostar$, and this results in the inconsistency between the observational estimates of SFRD and $\\rhostar$. But are there any other possibilities to explain the above inconsistency, or is it possible that our simulations are incorrect? Below we discuss two points for such possibilities. First possibility is that our simulation might be overproducing the low-mass galaxies at $z>4$. Recent observational and theoretical studies show that the low-mass, high-$z$ galaxies could have lower SF efficiencies than the normal local galaxies \\citep{Wolfe.Chen:06,Gnedin.Kravtsov:10} owing to lower molecular hydrogen fractions in low-metallicity environments. If this is true, our current simulations might be overpredicting the $\\rhostar$ at $z>4$. However if we revise our SF model to account for this effect, the SFRD will also decrease together with $\\rhostar$, and we might underpredict SFRD instead. In the future we will revise our SF model to consider the H$_2$ fraction in high-$z$ galaxies, and evaluate how strong this effect would be. The second possibility is that the IMF at high-$z$ could be different from the local one. All data in Figure~\\ref{fig:sf} assume the \\citet{Salpeter:55} IMF. As we emphasized in the last paragraph of Section~\\ref{sec:method}, the amount of gas converted into stars is a direct output of our simulations, therefore there are no uncertainties as to stellar IMF except the instantaneous gas-recycling fraction, i.e., the ratio between the gas expelled from stars (via stellar winds and supernovae) and the total initial stellar mass. In our simulations, we assume the Salpeter IMF whose instantaneous gas-recycling fraction ($\\beta$) is $\\sim$0.1 \\citep{Springel.Hernquist:03}. If we use the \\citet{Chabrier:03} IMF that has a lower number of low-mass stars, $\\beta$ increases to $\\sim$0.2, which would reduce the simulated $\\rhostar$ by about 10\\%. Note that $\\beta$ here only considers the gas-recycling from massive stars. If we take into account of the contribution from long-lived, low-mass stars, the value of $\\beta$ will increase by a factor of a few. However, this paper focuses on the high-$z$ galaxies, and we can safely ignore the gas-recycling from low-mass stars. Therefore the change from Salpeter to Chabrier IMF cannot fully account for the inconsistency between SFRD and $\\rhostar$. The top-heavy IMF in high-$z$ galaxies has been speculated by many authors \\citep[e.g.,][]{Larson:05,Fardal.etal:07,Dave:08,vanDokkum:08,Wilkins.etal:08,Bailin.etal:10}, and it has a much higher value of $\\beta$. Consequently, the resulting $\\rhostar$ will decrease significantly for a given SFRD, and it may alleviate the discrepancy in $\\rhostar$. However, the current theoretical models and observational support for top-heavy IMF are not very robust; for example, \\citet{vanDokkum.Conroy:10} suggested abundant low-mass stars in high-$z$ star-forming galaxies, which argues against a top-heavy IMF. In summary, we find a good agreement between our self-consistent cosmological simulations and the observational estimates of SFRD and $\\rhostar$ if we limit the comparison to galaxies with $\\Mstar > 10^9\\himsun$. In particular, the consistency between our simulations and the observational estimates by \\citet{Kistler.etal:09}, \\citet{Ouchi.etal:09} and \\citet{Gonzalez.etal:10} at $z>4$ is very encouraging. Our simulations predict the existence of numerous low-mass galaxies with $\\Mstar < 10^9\\himsun$, and these low-mass galaxies are not included in the current observational estimates of $\\rhostar$. This paper demonstrates that the current observational estimates and understanding of the formation of high-$z$ galaxies are still uncertain. To resolve this problem, we need future observations with increasing sensitivity and reduced uncertainties (e.g., JWST) to provide more robust constraints for the abundance of low-mass galaxies at $z>4$, as well as the improvement in the theoretical modeling of galaxy formation, particularly on star formation and its feedback." }, "1101/1101.2289_arXiv.txt": { "abstract": "Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are magnetar candidates, i.e., neutron stars powered by strong magnetic field. If they are indeed magnetars, they will emit high-energy gamma-rays which are detectable by Fermi-LAT according to the outer gap model. However, no significant detection is reported in recent Fermi-LAT observations of all known AXPs and SGRs. Considering the discrepancy between theory and observations, we calculate the theoretical spectra for all AXPs and SGRs with sufficient observational parameters. Our results show that most AXPs and SGRs are high-energy gamma-ray emitters if they are really magnetars. The four AXPs 1E 1547.0-5408, XTE J1810-197, 1E 1048.1-5937, and 4U 0142+61 should have been detected by Fermi-LAT. Then there is conflict between out gap model in the case of magnetars and Fermi observations. Possible explanations in the magnetar model are discussed. On the other hand, if AXPs and SGRs are fallback disk systems, i.e., accretion-powered for the persistent emissions, most of them are not high-energy gamma-ray emitters. Future deep Fermi-LAT observations of AXPs and SGRs will help us make clear whether they are magnetars or fallback disk systems. ", "introduction": "Anomalous X-ray pulsars (AXPs) and soft gamma-ray repeaters (SGRs) are two peculiar kinds of pulsar-like objects. Their persistent X-ray luminosities are in excess of their rotational energy loss rates, while at the same time they show no binary signature (review Mereghetti 2008). They also show recurrent SGR-type bursts (review Hurley 2009). Therefore, the energy budget of AXPs and SGRs is a fundamental problem in their studies. They are supposed to be magnetic field powered, i.e., magnetars (Thompson \\& Duncan 1995, 1996). Another possibility is that they are accretion powered systems, i.e., accretion from supernova fallback disks (Alpar 2001; Chatterjee et al. 2000; Xu et al. 2006). Then, it is of fundamental importance to determine whether they are magnetars or fallback disk systems. Solving this problem is also helpful to other high-energy astrophysical phenomena and related pulsar-like objects (Xu 2007; Tong et al. 2010a) Cheng \\& Zhang (2001) proposed that although AXPs are slowly rotating neutron stars, if their surface dipole magnetic field is strong enough (i.e., if they are really magnetars) then they can accelerate particles and emit high-energy gamma-rays which are detectable by Fermi-LAT according to the outer gap model (Zhang \\& Cheng 1997). However, Sasmaz Mus \\& Gogus (2010) reported a non-detection in a Fermi-LAT observation of AXP 4U 0142+61. This observation is in conflict with the outer gap model. Tong et al. (2010b) proposed that Fermi-LAT observations can help us distinguish between the magnetar model and the fallback disk model. Recently, the Fermi-LAT collaboration have published their observations for all known AXPs and SGRs (five SGRs and eight AXPs), where still no significant detection is reported (Abdo et al. 2010b). Considering this discrepancy between theory and observations, it is then very necessary to do a comprehensive study of this issue. In Cheng \\& Zhang (2001), only five AXPs are considered and the paremeters they used are very uncertain, e.g., the surface temperatures are estimated from the X-ray luminosities, etc. Now, we have very good observational data for more sources (see the McGill AXP/SGR online catalog). On the other hand, there are also developments of the outer gap model (e.g., Takata et al. 2010). In this paper, with up-to-date observational parameters of AXPs and SGRs, we consider the high-energy gamma-ray radiation properties of AXPs and SGRs in the outer gap model (Zhang \\& Cheng 1997; Takata et al. 2010) and compare them with Fermi-LAT observations. Section 2 is application of self-consistent outer gaps to AXPs and SGRs. We consider both the magnetar model and fallback disk model. Discussions and conclusions are presented in Section 3 and Section 4, respectively. ", "conclusions": "In this paper, we calculate the application of self-consistent outer gaps (Zhang \\& Cheng 1997; Takata et al. 2010) to the case of magnetars and compare the results with Fermi-LAT observations of all known AXPs and SGRs (Abdo et al. 2010b). Our calculations show that most AXPs and SGRs will emit high-energy gamma-rays and the gap closure mechanism is dominated by $\\gamma-\\gamma$ pair production process, if they are really magnetars. For the most gamma-ray luminous AXPs 1E 1547.0-5408, XTE J1810-197, 1E 1048.1-5937, and 4U 0142+61, their SEDs are above the Fermi-LAT sensitivity curve and should have been detected by Fermi-LAT. The observational upper limits of 4U 0142+61 are below the theoretical SEDs for large inclination angles. Therefore, there is conflict between outer gap model (Zhang \\& Cheng 1997) in the case of magnetars and Fermi-LAT observations. It is possible that AXPs and SGRs are wind braking, i.e., magnetars without a strong surface dipole field (Harding et al. 1999). It can not be excluded that AXPs and SGRs are fallback disk systems (Alpar 2001; Chatterjee et al. 2000; Xu et al. 2006). Considering the uncertainties in the outer gap modeling (e.g., the solid angle), future deeper Fermi-LAT observations are required. It will help us make clear whether AXPs and SGRs are magnetars or fallback disk systems." }, "1101/1101.3814.txt": { "abstract": "We present a detailed study of a bipolar, possible Type~I planetary nebula (PN), PHR1315-6555 (PN G305.3-03.1), that was discovered as part of the Macquarie/AAO/Strasbourg H$\\alpha$ planetary nebula project (MASH) and that we considered at the time was an excellent candidate for membership of the distant, compact, intermediate-age open cluster, ESO\\,96-SC04. The strong evidence for this association is presented here making this the only known example of a PN physically associated with a Galactic open cluster. Cluster membership is extremely important as it allows for very precise estimates of the fundamental properties of the PN as the cluster is at a known distance. The PN was discovered by one of us (QAP) during systematic MASH searches for new Galactic PNe of the AAO/UKST H$\\alpha$ survey and had been missed in earlier broadband surveys, including specific CCD studies of the host cluster. We present original discovery images and CTIO 4m MOSAIC-II camera follow-up narrow-band images that reveal its bipolar morphology. We also present: (i) low-resolution optical spectra that spectroscopically confirm the PN; (ii) accurate radial velocities of PN and cluster stars from high resolution spectroscopy which show they are consistent; and (iii) a reliable, independent distance estimate to the PN using a robust PN distance indicator which agrees with the published cluster distance to within the errors. We also provide preliminary estimates of basic PN properties and abundance estimates from deeper spectra that show it to be of possible Type~I chemistry. This is also consistent its estimated turn-off mass. Taken together these findings present a powerful case for clear physical association between the PN and host cluster. Results for this association will be of considerable interest to specialists across differing astrophysical disciplines, including PNe, white dwarfs, and open clusters. ", "introduction": "An association between a planetary nebula (PN) and an open star cluster is a very valuable astrophysical tool. This is because the accurate cluster distance, determined from a colour-magnitude diagram (CMD), constrains the physical parameters of the PN and central star (CSPN) to exceptional precision. The age and mass of the progenitor star can be tightly constrained from theoretical cluster isochrones, while CSPN photometry allows a precise determination of its intrinsic luminosity and mass. The progenitor star mass, which can be related to the chemistry of the resulting PN (from spectroscopy), provides a rare additional datum for the fundamental white dwarf (WD) initial-to-final mass relation (IFMR) currently best determined from cluster white dwarfs (e.g. Williams et al. 2004; Ferrario et al. 2005; Kalirai et al. 2008; Dobbie et al. 2009) which intimately links WD properties to their main-sequence progenitors. A robust IFMR is a key component of using WD luminosity functions to constrain the age of the Galactic disk (using the field WD population) and open clusters (using the cluster population) and is also key to mapping the build up of carbon and nitrogen in galaxies. Note that measuring precise stellar masses from PN is also possible, if somewhat controversial -- see the discussion in Gesicki \\& Zijlstra (2007). Unfortunately, the number of PNe that are genuine members of Galactic star clusters of all kinds is extremely small (there are four currently known in Galactic globular clusters; in M15, M22, Pal 6 and NGC 6441; Jacoby et al. 1997). Our discovery (Parker et al. 2006) of a faint bipolar PN (PHR1315-6555) within 23~arcseconds of the projected centre of the distant, compact, intermediate-age Galactic open cluster ESO 96-SC04, is therefore of great interest.\\footnote{ At least two PN candidates have also been found in globular clusters in external galaxies, e.g. Larsen (2008), for a peculiar PN candidate in a globular cluster in the Fornax dwarf spheroidal galaxy and Minniti \\& Rejkuba (2002) for a PN candidate in globular cluster G169 in Centaurus A (NGC5128).} %This PN open cluster association, is, as we will show, currently the only proven example in our Galaxy. Many other PN-cluster candidates (e.g. Bonatto, Bica \\& Santos 2008) have been shown to be either likely or definite line-of-sight superpositions. For example Majaess, Turner \\& Lane (2007) showed that 50 per cent of known possible PN-cluster associations were spatial coincidences, while Frew (2008) showed that a case against association can be made for nearly all previous implied PN-cluster pairings. Indeed, one of the better candidates for a PN-cluster association, that between the PN NGC~2438 and the open cluster M~46 (e.g. Pauls \\& Kohoutek 1996; Majaess, Turner \\& Lane 2007; Bonatto, Bica \\& Santos 2008) has recently been eliminated via a detailed radial velocity study of the PN and large numbers of cluster stars (Kiss et al. 2008). In this paper we present preliminary discovery and detailed subsequent follow-up data on our remarkable find and its host cluster. These are used to build a compelling and robust case for the veracity of the PN-cluster association. ", "conclusions": "We have found a faint, bipolar, high-excitation planetary nebula (PHR1315-6555) of possible Type~I chemistry in the intermediate-age open cluster ESO\\,96-SC04, based on our initial discovery from the AAO/UKST H$\\alpha$ Survey and subsequent confirmatory spectroscopy. The PN was missed on earlier broadband CCD imaging studies of the cluster as it is compact and of relatively low surface brightness. We have several key arguments and other corroborating strands of evidence that together present an extremely compelling case for a physical association between the PN and the cluster. These arguments comprise very close angular separation of PN to the compact cluster core, excellent agreement of PN and cluster radial velocities, compatible, independent distance determinations to PN and cluster, consistent reddening estimates, %consistent sub-solar metallicity estimates for PN and several cluster stars which are also consistent for a Hyades age cluster and Galactic scale height arguments. Finally, the estimated PN physical properties are consistent with the cluster turn-off mass and distance. PHR1315-6555 is currently the only bona-fide PN known to be unequivocally associated with an open cluster in the Galaxy. The importance of this physical association to the WD initial-to-final mass relation is stressed. Follow-up observations are planned to unambiguously identify and measure the CSPN and address the question of its possible binarity." }, "1101/1101.2901_arXiv.txt": { "abstract": "The last few years have seen tremendous progress in our understanding of cataclysmic variable stars. As a result, we are finally developing a much clearer picture of their evolution as binary systems, the physics of the accretion processes powering them, and their relation to other compact accreting objects. In this review, I will highlight some of the most exciting recent breakthroughs. Several of these have opened up completely new avenues of research that will probably lead to additional major advances over the next decade. ", "introduction": "The study of cataclysmic variables (CVs) -- close binary systems containing an accreting white dwarf (WD) primary -- has been undergoing a renaissance over the last few years. As also recently noted by Paul Groot \\cite{groot2010}, the field had experienced a boom in the 80s and early 90s, but then seemed to suffer a bit of a slump. This seems to have been caused partly by the need to shift focus from what used to be a mostly ``object-centered'' view of the field to one that is more ``population-centered''. As I will try to show in this review, this slump is most definitely behind us. In fact, the last few years have seen a series of breakthroughs that are dramatically improving our understanding of CV evolution, accretion physics and the connection between CVs and related systems, such as accreting neutron stars (NSs) and black holes (BHs). Let me start, however, by providing some context for these advances. ", "conclusions": "" }, "1101/1101.1959.txt": { "abstract": "With the coming data deluge from synoptic surveys, there is a growing need for frameworks that can quickly and automatically produce calibrated classification probabilities for newly-observed variables based on a small number of time-series measurements. In this paper, we introduce a methodology for variable-star classification, drawing from modern machine-learning techniques. We describe how to homogenize the information gleaned from light curves by selection and computation of real-numbered metrics (\u00d2features\u00d3), detail methods to robustly estimate periodic light-curve features, introduce tree-ensemble methods for accurate variable star classification, and show how to rigorously evaluate the classification results using cross validation. On a 25-class data set of 1542 well-studied variable stars, we achieve a 22.8\\% overall classification error using the random forest classifier; this represents a 24\\% improvement over the best previous classifier on these data. This methodology is effective for identifying samples of specific science classes: for pulsational variables used in Milky Way tomography we obtain a discovery efficiency of 98.2\\% and for eclipsing systems we find an efficiency of 99.1\\%, both at 95\\% purity. We show that the random forest (RF) classifier is superior to other machine-learned methods in terms of accuracy, speed, and relative immunity to features with no useful class information; the RF classifier can also be used to estimate the importance of each feature in classification. Additionally, we present the first astronomical use of hierarchical classification methods to incorporate a known class taxonomy in the classifier, which further reduces the catastrophic error rate to 7.8\\%. Excluding low-amplitude sources, our overall error rate improves to 14\\%, with a catastrophic error rate of 3.5\\%. ", "introduction": "\\label{sec:intro} Variable star science (e.g., \\citealt{em08}) remains at the core of many of the central pursuits in astrophysics: {\\it pulsational} sources probe stellar structure and stellar evolution theory, {\\it eruptive and episodic} systems inform our understanding of accretion, stellar birth, and mass loss, and {\\it eclipsing} systems constrain mass transfer, binary evolution, exoplanet demographics, and the mass-radius-temperature relation of stars. Some eclipsing systems and many of the most common pulsational systems (e.g., RR Lyrae, Cepheids, and Mira variables) are the fundamental means to determine precise distances to clusters, to relic streams of disrupted satellites around the Milky Way, and to the local group of galaxies. They anchor the measurement of the size scale of the Universe. See \\citet{2009astro2010S.307W} for a recent review. The promise of modern synoptic surveys \\citep{2007AJ....134..973I}, such as the Large Synoptic Survey Telescope (LSST), is the promise of discovery of many new instances of variable stars \\citep{2007AJ....134.2236S}, some to be later studied individually with greater photometric and spectroscopic scrutiny\\footnote{High-precision photometry missions (Kepler, MOST, CoRoT, etc.) are already challenging the theoretical understanding of the origin of variability and the connection of some specific sources to established classes of variables.} and some to be used as ensemble probes to larger volumes. New classes (with variability reflecting physics not previously seen) and rare instances of existing classes of variables are almost certainly on the horizon (e.g., \\citealt{2007AJ....134.2398C}). Classification of variable stars---the identification of a certain variable with a previously identified group (``class'') of sources presumably of the same physical origin---presents several challenges. First, time-series data alone (i.e., without spectroscopy) provides an incomplete picture of a given source: this picture is even less clear the more poorly sampled the light curve is both in time and in precision. Second, on conceptual grounds, the observation of variability does not directly reveal the underlying physical mechanisms responsible for the variability. What the totality of the characteristics {\\it are} that define the nature of the variability may in principle be known at the statistical level. But {\\it why} that variability is manifest relies on an imperfect mapping of an inherently incomplete physical model to the data. (For example, the periodic dimming of a light curve may be captured with a small number of observable parameters but the inference that that source is an eclipsing one requires a theoretical framework.) This intermingling of observation and theory has given rise to a taxonomy of variable stars (for instance, defined in the GCVS\\footnote{General Catalog of Variable Stars, {\\tt http://www.sai.msu.su/groups/cluster/gcvs/gcvs/}}) that is based on an admixture of phenomenology and physics. Last, on logistical grounds, the data volume of time-series surveys may be too large for human-intensive analysis, follow-up, and classification (which benefits from domain-specific knowledge and insight). While the data deluge problem suggests an obvious role for computers in classification\\footnote{Not discussed herein are the challenges associated with {\\it discovery} of variability. See \\citet{2009MNRAS.400.1897S} for a review.}, the other challenges also naturally lend themselves to algorithmic and computational solutions. Individual light curves can be automatically analyzed with a variety of statistical tools and the outcome of those analyses can be handled with machine-learning algorithms that work with existing taxonomies (however fuzzy the boundary between classes) to produce statistical statements about the source classification. Ultimately, with a finite amount of time-series data we wish to have well-calibrated probabilistic statements about the physical origin and phenomenological class of that source. While straightforward in principle, providing a machine-learned classifier that is accurate, fast, and well-calibrated is an extraordinarily difficult task on many fronts (see discussion in \\citealt{2008AIPC.1082..257E}). There may be only a few instances of light curves in a given class (``labelled data'') making training and validation difficult. Even with many labelled instances, in the face of noisy, sometimes spurious, and sparsely sampled data, there is a limit to the statistical inferences that can be gleaned from a single light curve. Some metrics (called, in machine-learning parlance,``features'') on the light curve may be very sensitive to the signal-to-noise of the data and others, particularly frequency-domain features, may be sensitive to the precise cadences of the survey (\\S \\ref{ss:surveydependence}). For computationally intensive feature generation (e.g., period searches) fast algorithms may be preferred over slower but more robust algorithms. Machine learning in variable star classification has been applied to several large time-series datasets \\citep{2004AJ....128.2965W,2007arXiv0712.2898W,2007debo,2008AN....329..288M,2009A&A...494..739S,2010blom}. A common thread for most previous work is application of a certain machine-learning framework to a single survey. And, most often, the classification is used to distinguish/identify a small set of classes of variables (e.g., Miras and other red giant variability). \\citet{2007debo} was the first work to tackle the many-class ($>20$) problem with multiple survey streams. \\citet{2007debo} also explored several classification frameworks and quantitatively compared the results. The purpose of this work is to build a many-class classification framework by exploring in detail each aspect of the classification of variable stars: proper feature creation and selection in the presence of noise and spurious data (\\S \\ref{sec:features}), fast and accurate classification (\\S \\ref{sec:methods}), and improving classification by making use of the taxonomy. We present a formalism for evaluating the results of the classification in the context of expected statistical risk for classifying new data. We use data analyzed by Debosscher et al.\\ to allow us to make direct comparison with those results (\\S \\ref{sec:results}). Overall, we find a 24\\% improvement in the misclassification rate with the same data. The present work only makes use of metrics derivable from time-domain observations in a single bandpass; color information and context (i.e., the location of the variable in the Galaxy and with respect to other catalog sources) are not used. In future work, we will explore how machine-learned classifiers can be applied across surveys (with different characteristics) and how context and time-domain features can be used in tandem to improve overall classification. \\\\ \\pagebreak ", "conclusions": "We have presented a thorough study of automated variable star classification from sparse and noisy single-band light curves. In the 25-class problem considered by \\citet{2007debo}, which includes all of the most important variable star science classes, we obtain a 24\\% improvement over their best classifier in terms of misclassification error rate. We attribute this improvement to all of the following advances: \\begin{itemize} \\item {\\bf Better periodic feature estimation.} Our Lomb-Scargle period-fitting code is both fast and accurate. With the same random forest classifier, the average error rate using our periodic feature estimates is 23.8\\%, compared to an error rate of 26.7\\% using only Debosscher's period feature estimates, representing an improvement of 11\\%. \\item {\\bf Use of predictive non-periodic features.} Simple summary statistics and more sophisticated model parameters give a significant improvement. Using both our periodic and non-periodic features, the random forest error rate is 22.8\\%, a 4\\% improvement over using only our periodic features. \\item {\\bf More accurate classification methods.} All of the methods considered in this paper, save the single-tree models, achieve a statistically-significant improvement over Debosscher's best classifier. Our random forest classifier, applied to the exact features used by that paper, achieves an 11\\% improvement over their best error rate, 30\\%. \\end{itemize} Another contribution is our discussion of model validation through estimation of the expected prediction error on new data by cross-validation. We presented cross-validation as a statistically-rigorous way to both tune a classifier and select between competing methods. Indeed, all of the numbers quoted in this paper are 10-fold cross-validation estimates. We have shown the adeptness of tree-based classifiers in the problem of variable star classification. We demonstrated the superiority of the random forest classifier in terms of error rates, speed, and immunity to features with little useful classification information. We outlined how to calculate the optimal probability threshold to obtain pure and complete samples of specified sub-classes, and showed that the multi-class random forest is often superior to the one-versus-all random forest in this problem. We advocate the continued use of this method for other classification problems in astronomy. Furthermore, we described how the random forest classifier can be used to estimate the importance of each feature by computing the expected classification gains versus replacing that feature with random noise. In the variable star classification problem, it was found that several non-periodic features have high importance. A classifier built only on the non-periodic features still performs quite well, attaining 27.6\\% error rate. Finally, this paper is the first to use the known variable-star taxonomy both to train a classifier and evaluate its result. We introduced two different classification methods to incorporate a hierarchical taxonomy: HSC, which builds a different classifier in each non-terminal node of the taxonomy, and HMC, which fits a single classifier, penalizing errors at smaller depths in the taxonomy more heavily. We demonstrated that both of these methods perform well, in terms of classification rate and catastrophic error rate. The class taxonomy was also used to construct the notion of catastrophic error rate, which considers as catastrophic any error made at the top level of the hierarchy. Several open questions remain with regard to the automated classification of astronomical time series. Many of these questions will be addressed by us in future publications, where we will expand on the methodology presented here and attempt to classify data from other surveys, such as the All Sky Automated Survey (ASAS), SDSS Stripe 82, and the Wide Angle Search for Planets (WASP). Some of the questions that we will address are: \\begin{itemize} \\item If we train a classifier on a set of objects from one (or multiple) survey(s), will that classifier be appropriate to predict the classes of objects from the new survey? This question is of great importance because presumably a set of known (labeled) variable stars will be compiled from previous surveys to train a classifier for use on a new survey. \\item What features are robust across a wide range of different surveys, each with different cadences? If some sets of features are robust to survey design and cadence, those should be used in lieu of survey-dependent features in a classifier. In this paper, we have excluded any feature that was blatantly survey-dependent (such as any that used mean flux), but this does not guarantee that some features will not have survey dependence. \\item How does mis-labelled training data affect the classifier accuracy? Can mis-labelled data be effectively detected and cleaned from the classification set? \\item How can a classifier be trained to efficiently identify outliers/new types of variables? Future surveys will unquestionably discover new science classes that do not fit under any of the training classes. Recent methodology has been developed in the statistics literature for outlier discovery in random forest classifiers, and we plan to adapt this for variable star classification. \\item How are the error rates of a classifier affected by computational limitations (where, perhaps some CPU-intensive or external server-dependent features are not used)? In automated classification of astronomical sources, there is often a time sensitivity for follow-up observations. Presumably there are more useful features for classification than the ones that we employed in this paper, but they may be expensive to compute or retrieve for each observation. This trade-off between error rate and computation time must be explored. \\end{itemize} Finally, as a longer-term goal, we are striving to develop methodology that can be used on a LSST-caliber survey. This means that our methods must be fast enough to compute features and class probabilities for thousands of objects per night, work well at an LSST cadence, be applicable to multi-band light curves, and perform classification for all types of astronomical objects, including transients, variable stars, and QSOs. Our task, looking forward, is to address each of these problems and develop methodology for fast and accurate classification for LSST. % acknowledgements" }, "1101/1101.0629_arXiv.txt": { "abstract": "{Optical modules for X-ray telescopes comprise several double-reflection mirrors operating in grazing incidence. The concentration power of an optical module, which determines primarily the telescope's sensitivity, is in general expressed by its on-axis effective area as a function of the X-ray energy. Nevertheless, the effective area of X-ray mirrors in general decreases as the source moves off-axis, with a consequent loss of sensitivity. To make matters worse, the dense nesting of mirror shells in an optical module results in a mutual obstruction of their aperture when an astronomical source is off-axis, with a further effective area reduction.} {To ensure the performance of X-ray optics for new X-ray telescopes (like NuSTAR, NHXM, ASTRO-H, IXO), their design entails a detailed computation of the effective area over all the telescope's field of view. While the effective area of an X-ray mirror is easy to predict on-axis, the same task becomes more difficult for a source off-axis. It is therefore important to develop an appropriate formalism to reliably compute the off-axis effective area of a Wolter-I mirror, including the effect of obstructions.} {Most of collecting area simulation for X-ray optical modules has been so far performed along with numerical codes, involving ray-tracing routines, very effective but in general complex, difficult to handle, time consuming and affected by statistical errors. In contrast, in a previous paper we approached this problem from an analytical viewpoint, to the end of simplifying and speeding up the prediction of the off-axis effective area of unobstructed X-ray mirrors with any reflective coating, including multilayers.} {In this work we extend the analytical results obtained: we show that the analytical formula for the off-axis effective area can be inverted, and we expose in detail a novel analytical treatment of mutual shell obstruction in densely nested mirror assemblies, which reduces the off-axis effective area computation to a simple integration. The results are in excellent agreement with the findings of a detailed ray-tracing routine.} {} ", "introduction": "Optical modules for X-ray telescopes consist of a number of grazing incidence mirror shells with a common axis and focus. In a widespread design, the Wolter's, the mirrors comprise two consecutive segments, a paraboloid and a hyperboloid, in order to concentrate X-rays by means of a double reflection. The two reflections occur at the same incidence angle for X-rays coming from an on-axis source at astronomical distance (Van~Speybroeck~\\& Chase~\\cite{VanSpeybroeck}). The optical design of the module is primarily dependent on the required effective area, which determines the telescope's sensitivity. The most representative indicator of the concentration power of an optical module is assumed in general to be the effective area on-axis. However, the effective area in general decreases as the X-ray source moves off-axis. This in turn diminishes the telescope's sensitivity, hence may represent a severe limitation to its {\\it field of view}. For this reason, the prediction of the on- and off-axis effective area modules is a very important task in the development of new X-ray telescopes such as NuSTAR (Hailey~et al.~\\cite{Hailey}), NHXM (Basso~et al.~\\cite{Basso}), ASTRO-H (Kunieda~et al.~\\cite{Kunieda}), and IXO (Bookbinder~\\cite{Bookbinder}). While the effective area on-axis is in general easy to calculate for a typical Wolter-I mirror module configuration, the computation becomes more difficult off-axis, because of the variable incidence angles over the two reflecting surfaces and the variable fraction of singly-reflected X-rays that are not focused and contribute to the stray light (see e.g., Cusumano~et al.~\\cite{Cusumano}). To make things worse, the assembled mirrors can shade each other if they are not spaced enough, which contributes to an even steeper decrease in the collecting area off-axis. A widespread method for performing the calculation, accounting for all these factors, has made use of accurate ray-tracing routines (see, e.g., Mangus~\\& Underwood~\\cite{Mangus}; Zhao~et~al.~\\cite{Zhao}), which reconstruct the paths of a selection of X-rays impinging a mirror module. These numerical codes are in general accurate, because they simulate the real incidence of rays on the optical system. However, they are complex and time consuming, especially whenever the simulation includes wide-band multilayer coatings to extend the reflectivity beyond 10 keV (Joensen~et~al.~\\cite{Joensen}; Tawara~et~al.~\\cite{Tawara}). The reason is that the multilayer reflectivity computation, which has to be performed for every ray traced, is a complex procedure especially for wideband multilayers, which comprise many ($\\sim$ 200) couples of layers. For this reason, although the ray-tracing approach should not be disregarded, it is interesting to derive analytical formulae for the effective area off-axis. In attempting to achieve this, Van~Speybroeck~\\& Chase (\\cite{VanSpeybroeck}) discovered, by analyzing the results of ray-tracing simulations, that the geometric collecting area of a Wolter-I mirror decreases with the off-axis angle $\\theta$ of the source, according to the formula \\begin{equation} A_{\\infty}(\\theta) = A_{\\infty}(0)\\,\\left(1-\\frac{2\\theta}{3\\alpha_0}\\right), \\label{eq:SC_formula} \\end{equation} where $A_{\\infty}(0)$ is the on-axis geometric area and $\\alpha_0$ is the incidence angle for an astronomical source on-axis. More recently, a semi-analytical method for computing the off-axis effective area has been applied to solve the problem of optimizing multilayer recipes to the telescope's field of view (Mao~et~al.~\\cite{Mao99}; Mao~et~al.~\\cite{Mao00}; Madsen~et~al.~\\cite{Madsen}). In this approach, the effective area is computed by integrating the mirror reflectivity over the incidence angles, after weighting it over an appropriate function, $W_{\\mathrm{inc}}$, derived from a ray-tracing. In a previous paper (Spiga~et~al.~\\cite{Spiga2009}), we already derived a {\\it completely} analytical method for computing the off-axis effective area of a Wolter-I mirror with any reflective coating. In a subsequent article (Spiga~\\&~Cotroneo~\\cite{Spiga2010}), we developed this formalism, deriving the analytical expression of the aforementioned $W_{\\mathrm{inc}}$ function, which represents the distribution of the off-axis effective area over the incidence angles of rays, and we also used it to face the problem of multilayer optimization (Cotroneo~et~al.~\\cite{Cotroneo2010}). The results were accurately verified as well, by comparison with the outcomes of a ray-tracing program. However, in these previous works, we did not consider the mutual shading (also known as {\\it vignetting}) of mirrors, which may occur in mirror modules when shells are densely nested together. While in general the mirror module is designed to avoid any vignetting on-axis, the problem may arise for sources off-axis and cause a further loss of effective area. Consequently, the results could be applied only to single mirror shells, or to mirror assemblies for which the mutual obstruction is known to be negligible over all the field of view. In this work, we overcome these limitations and extend the developed formalism to the general case of obstructed Wolter-I mirrors in X-ray optical modules. We still assume that the mirror profile can be approximated by a double cone as far as the sole effective area is concerned, a condition in general fulfilled by optics with large $f$-numbers. In Sect.~\\ref{single}, we briefly review the results obtained for unobstructed single mirror shells, and in addition we show how the analytical formalism can be inverted to derive the product of the two reflectivities from the desired effective area variation with the off-axis angle. In Sect.~\\ref{Vign}, we describe the geometrical parameters driving the nested mirror obstructions, derive the expression of the vignetting coefficients, and obtain {\\it an analytical formula for the off-axis effective area of an obstructed mirror} (Eq.~(\\ref{eq:area_total_fin})). We then derive in Sect.~\\ref{Application} some analytical expressions for the obstructed geometric area, and we apply the results to the geometrical optimization of the module. In Sect.~\\ref{Validation}, we prove the validity of the analytical formulae by means of a ray-tracing routine. The results are briefly summarized in Sect.~\\ref{Conclusions}. ", "conclusions": "We have developed the analytical formalism for the off-axis effective area of Wolter-I mirror shells, in double cone approximation, which we began to describe in a previous paper (Spiga~et~al.~\\cite{Spiga2009}). We have shown that the analytical expression of the effective area off-axis can be inverted to derive the product of the reflectivity of the two segments (Sect.~\\ref{invcomp}). This might be useful to future developments for computing a suitable multilayer recipe to return the desired effective area trend off-axis. We have found analytical expressions for the vignetting coefficients (Sect.~\\ref{obscoeff}), for the three possible sources of obstruction in nested mirror modules, as a function of the azimuthal coordinate of the mirror surface. Using the vignetting coefficients, we have derived an integral formula (Eq.~(\\ref{eq:area_total_fin})) for the obstructed effective area of a Wolter-I X-ray mirror in double cone approximation, with any reflective coating, including multilayers. The computation only requires the standard routines for the reflectivity of the coating, and an integration over the azimuthal coordinate of the mirror shell. We have obtained analytical expressions of the obstructed geometric area (Sect.~\\ref{geometric}) for the case of a source at infinite distance, and applied the result to the problem of designing an optical module that does not suffer from the mutual obstructions of mirrors (Sect.~\\ref{design}). Finally, the results have been validated by means of a comparison with the findings of a detailed ray-tracing (Sect.~\\ref{Validation}). As a final application, we note that each vignetting coefficient can be adapted to estimate the unwanted vignetting caused by collimators aimed at reducing the {\\it stray light} in mirror modules (see, e.g., Cusumano~et~al.~\\cite{Cusumano}). For example, $V_1$ would quantify the vignetting of the baffle at the entrance pupil, if $R^*_{\\mathrm M}$ is interpreted as the outer radius of the collimator ring and $L_1^*$ as its distance from the intersection plane. With analogous substitutions, $V_3$ would represent the vignetting of a baffle located at the exit pupil, and $V_2$ would express the vignetting of the baffle at the intersection plane, even though this kind of baffle can be designed to avoid any obstruction of focused rays (Sect.~\\ref{no_obst})." }, "1101/1101.5532_arXiv.txt": { "abstract": "Using the publicly available VESPA database of SDSS Data Release 7 spectra, we calculate the stellar Mass Weighted Age (hereafter MWA) as a function of local galaxy density and dark matter halo mass. We compare our results with semi-analytic models from the public Millennium Simulation. We find that the stellar MWA has a large scatter which is inherent in the data and consistent with that seen in semi-analytic models. The stellar MWA is consistent with being independent (to first order) with local galaxy density, which is also seen in semi-analytic models. As a function of increasing dark matter halo mass (using the SDSS New York Value Added Group catalogues), we find that the average stellar MWA for member galaxies increases, which is again found in semi-analytic models. Furthermore we use public dark matter Mass Accretion History (MAH) code calibrated on simulations, to calculate the dark matter Mass Weighted Age as a function of dark matter halo mass. In agreement with earlier analyses, we find that the stellar MWA and the dark matter MWA are anti correlated for large mass halos, i.e, dark matter accretion does not seem to be the primary factor in determining when stellar mass was compiled. This effect can be described by down-sizing. ", "introduction": "The spectra of galaxies encodes information about the histories of the component stellar populations, dust, and star formation. Various tools have been developed to extract this information \\citep[e.g.,][]{Heavens:1999am,Tojeiro:2007wt} and previous works have compared the resulting extracted information with both extrinsic and intrinsic galaxy properties. These approaches rely on the assumption that the evolution of the stellar populations are well understood and that the current modeling of stellar population is accurate. The MOPED \\citep{Heavens:1999am} routine implements the general process of reforming a complex dataset (e.g., a galaxy spectra) into a set of parameters (e.g., star formation rate, metallicity) and parameter combinations, assuming uncorrelated noise, such that the data compression is loss less \\citep[see also][]{2010arXiv1010.5907G}. \\cite{Mateus:2008wf} used MOPED-derived stellar masses and luminosities of SDSS Data Release 3 \\citep[][]{SDSSDR3} galaxies to build marked correlations functions, and compared with marked correlations of semi analytic models of galaxies in the Millennium Simulation. More recently, \\cite{Ferreras:2010wn} applied a Principal Component Analysis technique, to decompose the spectra of low redshift ($z<0.1$) SDSS early-type galaxies into two quantities: the average stellar age for the galaxy, assuming metal-rich stellar populations, and the fraction of recent star formation. They found little dependence (and a large scatter) of recent star formation and average stellar age on host halo mass, but find a stronger correlation on local processes, e.g., galactic velocity dispersion. An easily accessible, robust code, is the VErsatile SPectral Analysis\\footnote{http://www-wfau.roe.ac.uk/vespa/} \\citep[hereafter VESPA, see][for more details]{Tojeiro:2007wt,Tojeiro:2009kk} package, which recovers star formation and metallicity histories of the galactic spectra using synthetic stellar population models. The software recovers histories in adaptive age bins according to the signal-to-noise of the galaxy spectrum on a case by case basis and addresses the age-metallicity relation. Two popular synthetic stellar population models are included in the VESPA output, those of \\citet[][hereafter BC03]{2003MNRAS.344.1000B}, and the \\cite{Maraston:2004em} \\& \\citet[][]{Maraston:2008nn} hereafter M05, which differ in their respective resolutions, and the use of empirical libraries to model the thermally pulsating asymptotic giant branch. Furthermore VESPA corrects for galactic extinction using the dust maps of \\cite{dustmaps}, and fits for the dust in each galaxy using either a one or two parameter dust model. VESPA has been used by, e.g., \\cite{Tojeiro:2009kk} to compare the BC03 and M05 models, in particular, the recovered total stellar mass today (finds good agreement) and the mass averaged metallicity (finding less agreement). \\cite{Tojeiro:2010up} used VESPA to model the color evolution of high signal-to-noise SDSS Data Release 7 Luminous Red Galaxies as a function of color, luminosity and redshift. In this paper we compare the results of VESPA, in particular the typical time scale of stellar mass assembly, the stellar Mass Weighted Age (hereafter MWA), with similar results from N-body simulations and semi analytic models. In particular we address the following questions: \\begin{enumerate} \\item Does the stellar MWA depend on the local density? \\item Are similar trends seen in semi analytic models? \\end{enumerate} To this end, we examine the correlation of local galaxy density with the VESPA calculated stellar MWA, and compare with the stellar MWA of the semi analytic models in the Millennium Simulation \\citep{Springel:2005nw,BoylanKolchin:2009nc}. \\begin{enumerate} \\setcounter{enumi}{+2} \\item Is the stellar MWA correlated with the mass of the dark matter halo that the galaxy inhabits? \\end{enumerate} In this case, we match the galaxies used in VESPA with those from the New York Value Added Group catalogue \\citep[][]{Blanton:2004aa,Yang:2007yr} and explore the average stellar MWA of member galaxies as a function of dark matter halo mass. \\begin{enumerate} \\setcounter{enumi}{+3} \\item Is the stellar MWA correlated with the dark matter MWA? \\end{enumerate} We address this by comparing the results of the above analysis with the results of Mass Accretion History code \\citep{Zhao:2008wd}, rebinned in terms of VESPA time bins for direct comparison to the observations. The layout of the paper is as follows: in \\S\\ref{data} we describe the observational and simulated data and the process of reconstructing the stellar MWA of a galaxy. We continue in \\S\\ref{method} by detailing our measurement of local galaxy density and dark matter halo mass, for both the observed and semi analytic samples, and present the results in \\S\\ref{results}. We conclude and discuss in \\S\\ref{conclusions}. Throughout the paper we employ a flat $\\Lambda$CDM cosmology with ($h,\\, \\Omega_{\\Lambda},\\,\\Omega_{m},\\,\\sigma_8$) given by ($0.7 \\,\\km\\, \\s^{-1}\\,\\Mpc,0.7,0.3,0.8$) ", "conclusions": "\\label{conclusions} We used the results of VESPA \\citep{Tojeiro:2007wt,Tojeiro:2009kk} which analyzed $10^6$ SDSS DR7 \\citep{SDSSDR7} galaxy spectra, to calculate the stellar Mass Weighted Age (MWA). To remove peculiarities in the data, we chose to only use high signal-to-noise spectra which allowed the measurement of stellar histories in greater than $4$ VESPA age bins. We then calculated the local galaxy density of the all the SDSS galaxies in cylinders of radius $2.25\\,\\h^{-1}\\,\\Mpc$ and length $4.5\\h^{-1}\\,\\Mpc$, and applied the standard $V_{max}$ correction to account for the magnitude limits of the SDSS. We also cross-matched the VESPA galaxies with the New York Value Added Group Catalog \\citep{Yang:2007yr}, which includes estimates of the parent dark matter halo mass. Using the public code of \\cite{Zhao:2008wd}, we determined the dark matter MWA as a function of halo mass. We additionally obtained the stellar MWA, local galaxy density, and dark matter halo masses of galaxies from within the Millennium Simulation \\citep{Springel:2005nw}, populated according to the semi analytic models of \\cite{DeLucia:2005yk}. We found a large scatter in the recovered stellar MWA, which we determined to be partially due ($\\sim 20\\%$) to the differences between the BC03 \\citep[][]{2003MNRAS.344.1000B} and M05 \\citep{Maraston:2008nn} stellar populations models, by comparing the medians of our distributions, and also by examining the results of \\cite{Tojeiro:2010up}, who compared the differences between models when applied to stacked LRG spectra with high signal-to-noise. A large dispersion is also observed in the stellar MWA of the semi-analytic models, suggesting that the dispersion is also inherent to the data. We found that the recovered stellar MWA using the different stellar population models agree to within $1.5$Gyrs and that the stellar MWA of most ($60\\%$ to $90\\%$ depending on stellar population model and galaxy sample) galaxies was older than $8$ Gyrs, independent of redshift, local galaxy density, dark matter halo mass or galaxy type. This is expected from the star formation history of the Universe, which peaks at around $z=1-2$, corresponding to $8-10$ Gyrs lookback time \\citep[e.g., see][]{2004Natur.428..625H}. We now return to the questions posed in the introduction; \\begin{enumerate} \\item Does the stellar MWA depend on the local density? \\item Are similar trends seen in semi-analytic models? \\end{enumerate} To first order, we found that stellar MWA does not appear to be related to local galaxy density in either the observed or simulated data. We found similar dispersions in the observed and simulated data and apparent flatness across local galaxy density \\citep[although see][]{2010MNRAS.402.1942C}. \\begin{enumerate} \\setcounter{enumi}{+2} \\item Is the stellar MWA correlated with the mass of the Dark Matter Halo that the galaxy inhabits? \\end{enumerate} We did find a correlation of older stellar MWA with increasing dark matter halo mass in the observed galaxy sample, independent of stellar population model, which was also seen in the semi-analytic models, albeit for a smaller range in dark matter halo mass. \\begin{enumerate} \\setcounter{enumi}{+3} \\item Is the stellar MWA correlated with the dark matter MWA? \\end{enumerate} We found that the dark matter MWA became anti correlated with the stellar MWA as the mass of dark matter halos increases. This is an observation of a ``downsizing'' effect \\citep[][]{1996AJ....112..839C} which describes the idea that the dark matter halo mass at which star formation is highest, shifts to lower masses at later times \\citep[see e.g., Fig 8 \\& 9 of ][]{2009ApJ...696..620C,2005MNRAS.356..495J}. Another meaning of ``downsizing'' is that more massive galaxies compiled more of their stars at higher redshift (and in a shorter time scale) than less massive galaxies as seen in Fig. \\ref{AMW_haloMassa1aa} and e.g., \\cite{2005ApJ...621..673T,2007ApJ...669..947J}. These results have been seen previously by measuring the Specific Star Formation Rates (the amount of star formation per solar mass) of sets of similar mass galaxies over a range of redshifts. Here, the downsizing effect is observed by the decrease in the stellar MWA as a function of decreasing stellar and dark matter halo mass, using VESPA stellar history reconstructions of the galaxy spectra. These conclusions are in agreement with other measures of ``downsizing'' using MOPED to reconstruct the Star Formation Rate (SFR) histories of SDSS DR3 galaxies \\citep[][found the SFR of massive galaxies peaked earlier than less massive galaxies]{2007MNRAS.378.1550P}. We find VESPA is well suited for reconstructing stellar histories using the included synthetic stellar population models, and agrees well with semi-analytical models drawn from larger simulations." }, "1101/1101.0121_arXiv.txt": { "abstract": "We present a search for radio transients in the field of the bright radio source 3C 286 using archival observations from the Very Large Array. These observations span 23 years and include 1852 epochs at 1.4 GHz in the C and D configurations. We find no transients in the field. The sensitivity of the observations is limited by dynamic range effects in the images. At large flux densities ($> 0.2$ Jy), single epoch observations provide a strong limit on the transient surface density. At flux densities near the dynamic range threshold, we use the requirement that transient sources must appear in consecutive epochs to be confirmed as real. This sets the sensitivity at low flux densities to transient durations of $\\tau \\sim 1$ day, while $\\tau > 1$ minute for high flux densities. At 70 mJy, we find a 1-$\\sigma$ limit on the surface density $\\Sigma < 3 \\times 10^{-3}$ deg$^{-2}$. At 3 Jy, we find a 1-$\\sigma$ limit $\\Sigma < 9 \\times 10^{-4}$ deg$^{-2}$. A future systematic search of the VLA archives can provide one to two orders of magnitude more sensitivity to radio transients. ", "introduction": "Radio transient (RT) sources probe the high energy population of the Universe. Known hosts to transient radio emission include neutron stars, black holes, supernovae, gamma-ray bursts, and highly magnetized stars and planets. Most of what is known about RTs has been learned from follow-up of events discovered at high-energy or optical wavelengths or through serendipitous discovery \\citep[e.g.,][]{2009A&A...499L..17B}. Systematic searches for RTs have been conducted but vast parameter space remains unexplored. Blind searches for RTs are an important scientific goal for major new radio telescope facilities such as the Allen Telescope Array \\citep{2010ApJ...719...45C,2010ApJ...725.1792B}, LOFAR \\citep{2009ASPC..407..318H}, ASKAP \\citep{2010PASA...27..272M}, and the Long Wavelength Array \\citep{2010AJ....140.1995L}. But there is significant opportunity to identify RTs through analysis of archival radio data. \\citet{2007ApJ...666..346B} used nearly 1000 observations of a blank field observed by the Very Large Array (VLA) over 20 years to identify a set of RTs that have no apparent counterpart at radio or optical wavelengths. \\citet{2010ApJ...711..517O} have suggested that these may be due to neutron stars. \\citet{2002ApJ...576..923L} and \\citet{2006ApJ...639..331G} conducted a search for RTs at 1.4 GHz through a comparison of the VLA NVSS and FIRST survey catalogs that identified a radio supernova in the nearby galaxy NGC 4216. \\citet{2010arXiv1011.0003B} recently completed a search at 843 MHz of the Molongolo Observatory Synthesis Telescope (MOST) data archives that uncovered 15 RTs and a larger number of variable sources. Some of the sources discovered in the Molongolo search appear similar to the RTs found by \\citet{2007ApJ...666..346B} in that they have no faint radio or optical counterpart. \\citet{2010AJ....140..157B} found a population of faint, variable radio sources in the galactic plane, the majority of which are without multi-wavelength counterparts. We describe here a search for radio transients in the field of the quasar 3C 286 with analysis of 1.4 GHz archival data from the VLA. One of the standard flux calibrators for the VLA, 3C 286 has been observed thousands of times over the life time of the array. This search builds on the work of \\citet{2007ApJ...666..346B}; however, observations in the vicinity of the bright (15 Jy) source 3C 286 place limits on the sensitivity that can be achieved. The dynamic range of these observations is limited by calibration errors and other systematic effects rather than by statistical noise. These calibration and systematic errors can produce an apparent source in the image that is well above the theoretical detection threshold. Accordingly, we require stronger evidence (such as appearance in consecutive epochs and comparisons with other cataloged sources in the field) to demonstrate the existence of an RT. We present the data and its analysis in \\S 2, a simulation demonstrating the ability to identify sources with these methods in \\S 3, our source detections and transient identification efforts in \\S 4, limits in transient surface density in \\S 5, and a summary in \\S 6. ", "conclusions": "We have presented an analysis of 1852 epochs of VLA observations spanning 23 years of the bright calibrator 3C 286 at 1.4 GHz. This data set provides an important search for radio transients brighter than 70 mJy. We do not find any transients, in contradiction with optimistic estimates of transient surface density from M09 but consistent with limits from other surveys. Differences between surveys may be a function of observing frequency, regions of sky covered, and systematic problems in recovering transients. The results are limited significantly by dynamic range of the imaging. If systematic errors had not contributed to the imaging, our surface density limit would apply at flux densities that are an order of magnitude lower. Nevertheless, the results demonstrates that searches around bright calibrators can provide unique information. Future searches of the VLA archives can improve on these results through the use of the larger number of observations of other standard calibrators such as 3C 48 and of observations at other frequencies. The examination of fainter calibrators is an important way to get closer to the statistical noise limits under the assumption of a fixed dynamic range limit. Finally, more complete models used in self-calibration may permit higher dynamic range imaging. If 1\\% of the data from the VLA archives consist of calibrators suitable for transient searching, we will have 2000 hours of usable data from the past 25 years, corresponding to an order of magnitude increase in sensitivity to radio transients." }, "1101/1101.2254_arXiv.txt": { "abstract": "Structural parameters are normally extracted from observed galaxies by fitting analytic light profiles to the observations. Obtaining accurate fits to high-resolution images is a computationally expensive task, requiring many model evaluations and convolutions with the imaging point spread function. While these algorithms contain high degrees of parallelism, current implementations do not exploit this property. With ever-growing volumes of observational data, an inability to make use of advances in computing power can act as a constraint on scientific outcomes. This is the motivation behind our work, which aims to implement the model-fitting procedure on a graphics processing unit (GPU). We begin by analysing the algorithms involved in model evaluation with respect to their suitability for modern many-core computing architectures like GPUs, finding them to be well-placed to take advantage of the high memory bandwidth offered by this hardware. Following our analysis, we briefly describe a preliminary implementation of the model fitting procedure using freely-available GPU libraries. Early results suggest a speed-up of around $10\\times$ over a CPU implementation. We discuss the opportunities such a speed-up could provide, including the ability to use more computationally expensive but better-performing fitting routines to increase the quality and robustness of fits. ", "introduction": "Recent trends in commodity computing hardware have seen a dramatic shift first from single-core processors to multi-core and then to accelerated platforms like graphics processing units (GPUs). GPUs were originally designed to speed up 3D graphics calculations for video games, but their immense memory bandwidth and arithmetic capabilities have seen them re-purposed for the needs of scientific computing. While unquestionably powerful, their radically different, massively-parallel architectures have shaken up the software community. Astronomy is one of many fields trying to adapt to these changes in computing hardware. While the area is still in its infancy, GPUs have already been shown to provide significant speed-ups across a range of astronomy problems. These include direct N-body simulation (e.g., \\citealt{HamadaEtal2009}), adaptive mesh refinement hydrodynamics (e.g., \\citealt{SchiveEtal2010}), galaxy spectral energy density calculations \\citep{JonssonPrimack2010}, gravitational microlensing \\citep{BateEtal2010}, correlation for radio telescopes (e.g., \\citealt{WaythEtal2009}) and coherent pulsar dedispersion \\citep{vanStratenBailes2010}. The approach taken in each of these cases has, however, been \\textit{ad hoc} in nature -- the transition to the GPU has been guided largely by hardware-specific documentation, code samples and simple trial and error. While such an approach has proven very successful for these early adopters, it is not clear that it will remain effective when it comes to more complex algorithms. Furthermore, in some cases the cost of re-implementing a code may be too large to gamble on a return (i.e., a speed-up) of unknown magnitude. In this paper we discuss the potential for accelerating the process of galaxy fitting using GPUs. Rather than tackling the challenge blind, we instead use a generalised method based on algorithm analysis as outlined in \\citet{BarsdellEtal2010}. The galaxy fitting process is described in Section \\ref{sec:GalaxyFitting}, which is followed by a full analysis of the problem in Section \\ref{sec:AlgorithmAnalysis}. A preliminary implementation and results are described in Section \\ref{sec:Results} before our summary discussion in Section \\ref{sec:Discussion}. ", "conclusions": "\\label{sec:Discussion} Many-core architectures like GPUs are now an important part of the computing landscape. While many software challenges remain, a generalised approach to analysing astronomy problems has proven very useful in tackling new GPU codes. Galaxy fitting looks to be a promising application of GPU technology. Significant speed-ups present the opportunity to perform faster fits, which may be crucial for the next generation of galaxy surveys. Alternatively, the additional processing speed could be fed back into the fitting routine to provide fits of much better quality in the same length of time, helping to overcome common problems such as local minima and unphysical results. While useful as a profiling tool, our prototype GPU code requires significant further development before it can be considered a viable alternative to other galaxy fitting codes in use by the astronomy community. Future work will address such development. Given the generality of our analysis, it is likely that other fitting problems in astronomy would also benefit from GPU acceleration. If one allows flexibility in the dimensionality of the problem, procedures such as spectral line or cube fitting become possible. Such problems will also be the subject of future work." }, "1101/1101.1899_arXiv.txt": { "abstract": "{In this paper, the CoRoT Exoplanet Science Team announces its 14th discovery. Herein, we discuss the observations and analyses that allowed us to derive the parameters of this system: a hot Jupiter with a mass of $7.6 \\pm 0.6$ Jupiter masses orbiting a solar-type star (F9V) with a period of only 1.5 d, less than 5 stellar radii from its parent star. It is unusual for such a massive planet to have such a small orbit: only one other known exoplanet with a higher mass orbits with a shorter period. ", "introduction": "Transiting exoplanets offer greater opportunities for the study and understanding of exoplanetary systems than those discovered by radial velocity measurements. Analysis of transit light curves yields planetary radii and enables tests for rings \\citep{barnes2004}, moons \\citep{sart1999}, and other planets through transit timing variations \\citep{mac2010}, while high-precision observations of primary and secondary transits can reveal some details of planetary atmospheres (which is not currently possible for non-transiting planets) and albedos \\citep[e.g.]{dem2009}, which is easier for transiting exoplanets but still possible for others. The potential of transiting exoplanets has inspired considerable effort towards their discovery, both from the ground and from space. While ground-based searches have discovered the majority of known transiting exoplanets to this point, space-based missions offer the greatest potential for discovery. Observing from space allows nearly continuous sampling and much better photometric precision, which is adversely affected by the atmosphere. This makes it possible to detect long-period transiting exoplanets, whose transits can easily be longer than a typical night, and smaller exoplanets, whose transits are too shallow to be detected from ground. The CoRoT (CO{\\it vection} RO{\\it tation and planetary }T{\\it ransits} space mission was the first space mission dedicated primarily to searching for transits \\citep{bag2009}. The mission has successfully demonstrated the advantages to space; given its orbit and the lack of atmosphere, it can observe the same field continuously for up to five months with remarkably high relative precision. This enabled the discovery of both the first transiting 'Super-Earth' \\citep[CoRoT-7b:]{leg2009,que2009} and the first temperate transiting gas giant \\citep[CoRoT-9b]{dee2010}. In this paper, we announce the discovery of the 14th transiting planet discovered by CoRoT; an unusually massive exoplanet orbiting an F9V star with metallicity consistent with Solar. In Sec. 2, we detail the CoRoT photometry. In Sec. 3, we describe the ground-based follow-up observations that we used to confirm the planetary nature of CoRoT-14b. In Sec. 4, we discuss our analysis of the light curves to extract the transit parameters and present the inferred planetary parameters. In Sec. 5, we analyze the parent star. Finally, we conclude our paper in Sec. 6, where we discuss how the properties CoRoT-14b compare to the ensemble of known transiting planets. ", "conclusions": "The most interesting quality of CoRoT-14b is its mass relative to its period -- only WASP-18b is both more massive and dense while being closer to its parent star. Figure~\\ref{per_ecc_mass} demonstrates this, plotting period vs. eccentricity for the know exoplanets with periods less than 10 days. When examining this plot, another characteristic of massive planets becomes apparent: they have a strong tendency towards elliptical orbits -- only 3 of the 12 ($\\sim$25\\%)transiting exoplanets that have masses greater than 2 $M_J$ and periods less than 10 days have $e=0$, not including those planets with unknown eccentricity, while 3 more of these orbit stars too faint to allow the orbital eccentricity to be measured readily. By contrast, transiting planets with masses less than 2 $M_J$ and periods less than 10 days have only a $\\sim$21\\% chance (13/63) of having a non-zero eccentricity. While it is impossible to draw any definitive conclusions with such a small sample size, these numbers suggest that that more massive planets may in truth have longer periods and higher eccentricities than less massive planets, although it is possible that some of these non-zero eccentricities are artifacts arising from the small number of RV measurements \\citep{she2008}. An examination of the theory for tidal circularization and orbital decay, arising from tides induced by the parent star on the exoplanet, shows that this is not unexpected (see Figure~\\ref{circtime}). Both of these phenomena have timescales that go as $Q M_pM_\\star^{2/3}P^{13/3}R_p^{-5}$ \\citep[see e.g.]{dob2004,ferr2008}. Assuming that $Q$, the quality factor, is approximately equal for all gas giants, we would expect that high mass planets with small radii will maintain their eccentricity (and semi-major axis) longer -- a tendency further accentuated by the fact that more massive planets have higher surface gravity, allowing them to resist inflation caused in part by by incident radiation from the parent star and therefore having smaller radii. However, this does not explain the circular orbit of the high mass planet/brown dwarf CoRoT-3b \\citep{bou2008} -- its circularization timescale is significantly longer than the age of the universe. It is possible that CoRoT-3b might be eccentric -- the RV observations used to measure this parameter are scattered over a year, making it difficult to rule out small, non-zero eccentricities. If both the adopted zero eccentricity and Q factor are correct, the properties of CoRoT-3b would be indicative of {\\it in situ} formation rather than migration, the generally accepted process by which short-period planets end up where they are. By contrast, CoRoT-14b is less massive and closer to its host star, leading to a much shorter circularization timescale. The observations of CoRoT14b are currently consistent with a circularly orbit -- it would therefore come as no surprise if this turns out to be the case in the end." }, "1101/1101.4191_arXiv.txt": { "abstract": "The Hercules Thick Disk Cloud \\citep{lar08} was initially discovered as an excess in the number of faint blue stars between quadrants 1 and 4 of the Galaxy. The origin of the Cloud could be an interaction with the disk bar, a triaxial thick disk or a merger remnant or stream. To better map the spatial extent of the Cloud along the line of sight, we have obtained multi-color UBVR photometry for 1.2 million stars in 63 fields approximately 1 square degree each. Our analysis of the fields beyond the apparent boundaries of the excess have already ruled out a triaxial thick disk as a likely explanation \\citep{lar10}. In this paper we present our results for the star counts over all of our fields, determine the spatial extent of the over density across and along the line of sight, and estimate the size and mass of the Cloud. Using photometric parallaxes, the stars responsible for the excess are between 1 and 6 kiloparsecs from the Sun, 0.5 -- 4 kpc above the Galactic plane, and extends approximately 3-4 kiloparsecs across our line of sight. It is thus a major substructure in the Galaxy. The distribution of the excess along our sight lines corresponds with the density contours of the bar in the Disk, and its most distant stars are directly over the bar. We also see through the Cloud to its far side. Over the entire 500 square degrees of sky containing the Cloud, we estimate more than 5.6 million stars and 1.9 million solar masses of material. If the over density is associated with the bar, it would exceed 1.4 billion stars and more than than 50 million solar masses. Finally, we argue that the Hercules-Aquila Cloud \\citep{bel07} is actually the Hercules Thick Disk Cloud. ", "introduction": "Studies of both stars and gas in the Galaxy are revealing significant structure and asymmetries in its motions and spatial distributions. Some examples of recent structure include the bar of stars and gas in the Galactic bulge (\\citealt{1991ApJ...379..631B}, \\citealt{1994ApJ...429L..73S}), the evidence from infrared surveys for a larger stellar bar in the inner disk (\\citealt{1992ApJ...384...81W}, \\citealt{1997MNRAS.292L..15L}, \\citealt{2005ApJ...630L.149B}), the outer ring \\citep{2003ApJ...588..824Y}, the discovery of the Sagittarius dwarf \\citep{1994Natur.370..194I,1995MNRAS.275..591I} and a significant asymmetry of unknown origin in the distribution of faint blue stars in Quadrant 1 (Q1) of the inner Galaxy \\citep{lar96}. Each of these observations provides a significant clue to the history of the Milky Way. When combined with the growing evidence for Galactic mergers in addition to the Sagittarius dwarf, i.e. the Monoceros stream (\\citealt{2002ApJ...569..245N}, \\citealt{2003MNRAS.340L..21I}), the Canis Major merger remnant (\\citealt{2004MNRAS.355L..33M}), the Virgo stream (\\citealt{2001ApJ...554L..33V}, \\citealt{2005ApJ...633..205M}) and the recent Hercules--Aquila cloud \\citep{bel07}, we now realize that the structure and evolution of our Galaxy have been significantly altered by mergers with other systems. Indeed the population of the Galactic Halo and possibly the Thick Disk as well, may be dominated by mergers with smaller systems. Larsen and Humphrey's asymmetry involves faint bluer stars in Quadrant 1 (Q1) of the inner Galaxy ($l = 20\\degree - 45\\degree$ at intermediate latitudes) characterized by an overdensity of $\\approx$ 30\\% when compared with complementary longitudes in the Quadrant 4 (Q4 , $l = 315\\degree - 340\\degree$). The initial discovery was made using star counts from the Minnesota Automated Plate Scanner Catalog of the POSS I (MAPS, \\cite{cab03} \\footnote{http://aps.umn.edu}). A more spatially complete survey \\citep{par03} with 40 contiguous fields above and below the plane in Q1 and in Q4 above the plane confirmed the star count excess and found that the asymmetry in Q1, while somewhat irregular in shape, was also fairly uniform and covered several hundred square degrees in Q1. It is therefore a major substructure in the Galaxy due to more than small scale clumpiness. The stars responsible for the excess were probable Thick Disk stars typically 1 -- 2 kpc from the Sun. \\cite{par04} also found an associated kinematic signature. The Thick Disk stars in Q1 have a much slower effective rotation rate $\\omega$, compared to the corresponding Q4 stars, with a significant lag of 80 to 90 km s$^{-1}$ in the direction of Galactic rotation, greater than the expected lag of 30 -- 50 km s$^{-1}$ for the Thick Disk population. The asymmetry is now designated the Hercules Thick Disk Cloud \\citep{lar08} (hereafter the Hercules Cloud). The release of the SDSS Data Release 5 (DR5) photometry in the same direction as the observed asymmetry in Q1 led to the discovery of another feature at much fainter magnitudes, the so-called Hercules-Aquila cloud \\citep{bel07}, however we suggest (\\S {4.3} that the over density is actually closer and at the same distances as the nearer Hercules Cloud. A second analysis \\citep{jur08} of the same dataset and led to the confirmation of our nearer Hercules Cloud at its approximate distances, though it was initially attributed to a possible stellar ring above the plane. Our comparison of the stellar density distributions in Q1 and Q4 above the plane \\citep{lar08} demonstrated that the excess is in Q1 only and is therefore not consistent with a ring. With the increasing evidence for Galactic mergers \\citep{1994Natur.370..194I, 1995MNRAS.275..591I, 2002ApJ...569..245N, 2003ApJ...596L.191N, 2004MNRAS.355L..33M, 2005ApJ...633..205M, 2003ApJ...588..824Y, 2006ApJ...639L..13W}, we now realize that the population of the Galactic Halo, and possibly the Thick Disk as well, may be dominated by mergers with smaller systems. The Hercules Cloud has no spatial overlap with the path of the Sagittarius dwarf through the Halo \\citep{2001ApJ...547L.133I}, and the predicted path of the Canis Major dwarf \\citep{2004MNRAS.355L..33M}, so its association with either of these well-studied features is unlikely. Our line of sight to the asymmetry is also interestingly in the same general direction as the stellar bar in the Disk \\citep{1992ApJ...384...81W, 1997MNRAS.292L..15L, 2000MNRAS.317L..45H, 2005ApJ...630L.149B}, but the bar is approximately 5 kpc from the Sun in this direction. Thus the stars showing the excess were mostly between the Sun and the bar, not directly above it. However the maximum extent of the star count excess along our line of sight was not known. Interpretation of the Hercules Cloud is not clear-cut. While it might well be the fossil remnant of a merger, the star count excess is also consistent with a triaxial Thick Disk or inner Halo as well as a dynamical interaction with the stellar bar especially given the corresponding asymmetry in the kinematics \\citep{par04}. A rotating bar in the Disk could induce a gravitational ``wake'' that would trap and pile up stars behind it \\citep{1992ApJ...400...80H, 1998ApJ...493L...5D}. Thus in response to the bar, there would not only be an excess of stars in Q1 over Q4, but those stars may show a measurable lag in their rotational velocities as observed in Q1. While similar, a triaxial Thick Disk could also yield different effective rotation rates because of noncircular streaming motions along its major axis. If the Thick Disk is triaxial, we would expect to observe the star count excess out to greater longitudes, but it appears to terminate near {\\it l} $\\sim$ 55$\\degree$ \\citep{par03}. To search for the asymmetry at greater longitudes from the Galactic center, our Paper I \\citep{lar10} extended the star counts to fainter magnitudes, corresponding to greater distances. Our results do not support the triaxial interpretation of the asymmetry. We find a statistically significant excess of faint blue stars for the two innermost Q1 fields at {\\it l} of 45$\\degree$ and 50$\\degree$, but the fields at the greater longitudes (55$\\degree$, 60$\\degree$, 65$\\degree$ and 75$\\degree$) show no significant excess including the faintest magnitude intervals. One of the greatest uncertainties concerning the nature of the Hercules Cloud is its spatial extent along the line of sight. Our earlier work \\citep{lar96,par03,par04} used photographic data having completeness limits of 18 -- 18.5 mag. To further explore its possible origins, we have mapped the extent of the spatial asymmetry to greater distances as a function of Galactic longitude and latitude. In Paper I we described our CCD observing program to much fainter limiting magnitudes. In the next section we present a brief summary of the observations and the data. In \\S {3} we describe our analysis of the star counts and population separation. The resulting map of the star count excess from photometric parallaxes and the size and mass of the Hercules Cloud are presented in \\S {4} and in the last section the implications for the origin of the asymmetry and the Hercules Cloud are discussed. ", "conclusions": "The star count asymmetry in Q1, the Hercules Cloud, is associated with a Thick-Disk-like density function having an upper boundary at $b\\approx40\\degree$. The strength of the feature decreases at lower latitudes, but then reappears below the plane, indicating that it is not an isolated stream confined to one side of the Disk. The photometric parallaxes for our ``Blue'' population shows the star count excess extending along various sight lines from $\\approx$ 1 to 6 kpc from the Sun. The regions showing the strongest excess along these lines of sight have a very interesting association with the increasing density in the direction of the bar in the Galactic plane and may have an associated kinematic signature \\citep{paperIII}. Thus the stars participating in the over-density in Q1 may be either members of a Thick Disk population associated with the bar or result from a dynamical response to the bar's passage. From the calculations presented in Section 4.2 we infer a total mass for this feature which ranges from a conservative estimate of 2\\e{6} $M_\\odot$ if it is a local overdensity to a potential upper limit of 5\\e{8} $M_\\odot$ if it is associated with the bar. Finally, we argue that the Hercules-Aquila Cloud is much closer to the Sun than previously reported and associated with the Hercules Thick Disk Cloud. We have also identified an excess in a populations of faint red stars in Q1, but much closer to the Sun. It is possible that these stars may extend into the Solar neighborhood and could be related to the local Hercules stream \\citep{1998A&A...335L..61R, 1999ApJ...524L..35D, 2000AJ....119..800D, 2007ApJ...655L..89B, 2009IAUS..254..139W}. This is certainly an interesting possibility but one not easily resolved. The local moving groups are defined by nearby bright stars distributed in a volume of space surrounding the Sun, while in our approach to mapping the asymmetry with relatively faint stars, direction is important. Further work on this asymmetry in Q1, or the Hercules Cloud, must concentrate on several questions. First of all, areas on the sky more than 1 square degree must be observed along each line of sight to increase the statistical significance of the detection. Far more fields in Q4 and also below the plane and in other directions not covered by SDSS are needed for Galactic structure studies in general, and also to improve our mapping of the excess below the plane. Large scale surveys at the lower latitudes may be required to statistically isolate the Hercules Cloud from the Galactic Disk. Does the Cloud or asymmetry extend towards the Galactic center? \\citet{jur08} traced the over-density associated with the Cloud to $l=355\\degree$, but it disappears by $l=340\\degree$ \\citep{lar08}. If it is associated with the bar or due to a gravitational interaction with the bar, as we suspect, then the excess would be expected to extend into Q4, but at much greater distances and fainter magnitudes. And finally, is the asymmetry related to the local Hercules Stream which passes through the Solar neighborhood? In Paper III, we discuss the kinematics of the associated stellar population and the possible origins of the Hercules Thick Disk Cloud." }, "1101/1101.4158_arXiv.txt": { "abstract": "A close substellar companion has been claimed to orbit the bright sdB star HD $149382$ with a period of $2.391$ d. In order to check this important discovery we gathered $26$ high resolution spectra over $55$ days with the {\\sc hermes} spectrograph on the $1.2$m Mercator telescope on La Palma, and analyzed the resulting radial velocities. Our data show no sign of any significant radial-velocity periodicities, and from the high precision of our measurements we rule out any RV variations with amplitudes higher than $0.79$ km/s on periods shorter than $50$ days. ", "introduction": "HD 149382 was recognised to be a subdwarf B star by \\cite{MacConnell1972}, and the spectroscopic analysis of \\cite{saffer1994} place it on the hot end of the extreme horizontal branch (EHB). With a magnitude of V\\,=\\,8.9 it is the brightest EHB star in the sky, and one of just a handful of such subdwarfs that can be easily observed with high-resolution spectroscopy on 1-m class telescopes. The first hints of a binary nature for HD\\,149382 came with the study of \\cite{ulla1998}, where they used infrared $JHK$ colours as an indicator of binarity, and concluded that HD\\,149382 has a very signficant IR excess, compatible with a companion of class K1 to G2. Such binaries are predicted to form via stable Roche lobe overflow, and result in orbital periods of hundreds of days \\citep{Han2003}. A red companion was clearly detected about 1\" away from the subdwarf in high resolution H-band imaging obtained with the {\\sc naomi} adaptive optics system at the William Herschel Telescope on La Palma, by \\cite{roy2005} (see Fig. \\ref{Ingrid}). At $\\sim$75\\,pc, this separation corresponds to $\\sim$75\\,AU making it too wide for the two stars to have interacted during the evolution of the primary. There is also no significant change in the separation between the images obtained in 2002 and 2004, so any orbital period must be on the order of decades, if these objects are gravitationally bound at all. \\begin{figure} \\includegraphics[height=.3\\textheight]{Ingrid.eps} \\caption{A red companion was clearly detected about 1\" away from the subdwarf in high resolution H-band imaging obtained with the {\\sc naomi} adaptive optics system at the William Herschel Telescope on La Palma, by \\cite{roy2005}} \\label{Ingrid} \\end{figure} Later, a close companion was claimed to be found around HD\\,149382 by \\cite{geier2009}, based on 15 high-resolution spectra ($R$\\,=\\,30000--48000) on three different spectrographs (ESO-2.2m/{\\sc feros}, CAHA-2.2m/{\\sc foces}, and McDonald-2.7m/Coud\\'e) taken within four years and one additional VLT/{\\sc uves} spectrum ($R$\\,=\\,80000). This was also the first claim that a close substellar companion, able to influence the stellar evolution of the host star, was found, thus challenging evolutionary experts to come up with a theory in which single stars with substellar companions were able to produce a subdwarf. This led to the profound question of what the actual influence of the companion might be, as such a low mass companion would be able to aid the formation of a hot subdwarf. ", "conclusions": "Our additional data clearly shows that there is as yet no evidence for a close substellar companion surrounding the bright sdB star HD\\,149382. We gathered 22 spectra with the {\\sc hermes} spectrograph on the Mercator Telescope. Our data show no sign of any radial-velocity periodicities, and from the high precision of our measurements we rule out any RV variations with amplitudes higher than 0.79 km/s on periods shorter than 50 days." }, "1101/1101.3913_arXiv.txt": { "abstract": "{Partially ionized plasma is usually described by single-fluid approach, where the ion-neutral collision effects are expressed by Cowling conductivity in the induction equation. However, the single-fluid approach is not valid for the time-scales less than ion-neutral collision time. For these time-scales the two-fluid description is better approximation.}{To derive the dynamics of magnetohydrodynamic waves in two-fluid partially ionized plasmas and to compare the results with those obtained under single-fluid description.}{Two-fluid magnetohydrodynamic equations are used, where ion-electron plasma and neutral particles are considered as separate fluids. Dispersion relations of linear magnetohydrodynamic waves are derived for simplest case of homogeneous medium. Frequencies and damping rates of waves are obtained for different parameters of background plasma.}{We found that two- and single-fluid descriptions give similar results for low frequency waves. However, the dynamics of MHD waves in two-fluid approach is significantly changed when the wave frequency becomes comparable or higher than ion-neutral collision frequency. Alfv\\'en and fast magneto-acoustic waves attain their maximum damping rate at particular frequencies (for example, the peak frequency equals 2.5 ion-neutral collision frequency for 50 $\\%$ of neutral Hydrogen) in wave spectrum. The damping rates are reduced for higher frequency waves. The new mode of slow magneto-acoustic wave appears for higher frequency branch, which is connected to neutral hydrogen fluid. }{The single-fluid approach perfectly deals with slow processes in partially ionized plasmas, but fails for time-scales smaller than ion-neutral collision time. Therefore, two-fluid approximation should be used for the description of relatively fast processes. Some results of single-fluid description, for example the damping of high-frequency Alfv\\'en waves in the solar chromosphere due to ion-neutral collisions, should be revised in future. } ", "introduction": "Astrophysical plasmas often are partially ionized. Neutral atoms may change the plasma dynamics due to collisions with charged particles. The ion-neutral collisions may lead to different new phenomena in plasma, for example the damping of magnetohydrodynamic (MHD) waves (Khodachenko el al. \\cite{Khodachenko2004}, Forteza et al. \\cite{Forteza2007}). Solar photosphere, chromosphere and prominences contain significant amount of neutral atoms, therefore the complete description of plasma processes requires the consideration of partial ionization effects. Braginskii ({\\cite{Braginskii1965}}) gave the basic principles of transport processes in plasma including the effects of partial ionization. Since this review, numerous papers addressed the problem of partial ionization in the different regions of solar atmosphere. Khodachenko and Zaitsev (\\cite{Khodachenko2002}) studied the formation of magnetic flux tube in a converging flow of solar photosphere, while Vranjes et al. (\\cite{Vranjes2008}) studied the Alfv\\'en waves in weakly ionized photospheric plasma. Leake and Arber (\\cite{Leake2005}) and Arber et al. ({\\cite{Arber2007}}) studied the effect of partially ionized plasma on emerging magnetic flux tubes and concluded that the chromospheric neutrals may transform the magnetic tube into force-free configuration. Haerendel (\\cite{Haerendel1992}), De Pontieu and Haerendel (\\cite{De Pontieu1998}), James and Erd\\'elyi (\\cite{James2002}), James et al. (\\cite{James2004}) considered the damping of Alfv\\'en waves due to ion-neutral collision as a mechanism of spicule formation. Khodachenko el al. (\\cite{Khodachenko2004}) and Leake et al. (\\cite{Leake2006}) studied the importance of ion-neutral collisions in damping of MHD waves in the chromosphere and prominences. Forteza et al. (\\cite{Forteza2007,Forteza2008}), Soler et al. (\\cite{Soler2009a, Soler2009b,Soler2010}) and Carbonell et al. (\\cite{Carbonell2010}) studied the damping of MHD waves in partially ionized prominence plasma with and without plasma flow. All these papers considered the single-fluid MHD approach, when inertial terms in the momentum equation of relative velocity between ions and neutrals are neglected. The partially ionized plasma effects are described by generalized Ohm's law with Cowling conductivity, which leads to the modified induction equation (Khodachenko el al. \\cite{Khodachenko2004}). Ambipolar diffusion is more pronounced during the transverse motion of plasma with regards to magnetic field, therefore the Alfv\\'en and fast magneto-acoustic wave are more efficiently damped. The slow magneto-acoustic waves are weakly damped in the low plasma beta case. Moreover, Forteza et al. (\\cite{Forteza2007}) found that the damping rate of slow magneto-acoustic waves derived through normal mode analysis is different from that estimated by Braginskii ({\\cite{Braginskii1965}}). The problem of discrepancy between normal mode analysis (Forteza et al. \\cite{Forteza2007}) and energy consideration (Braginskii {\\cite{Braginskii1965}}) is still an open question and the present study attempts to shed light on it. The single-fluid approach has been shown to be valid for the time-scales which are larger than ion-neutral collision time. However, the approximation fails for the shorter time scales, therefore the two-fluid approximation, which means the treatment of ion-electron and neutral gases as separate fluids, should be considered. The two-fluid approximation is valid for the time-scales larger than ion-electron collision time, which is significantly short due to Coulomb collision between ions and electrons. In this paper, we study MHD waves in two-fluid partially ionized plasma. The particular attention is paid to the wave damping due to ion-neutral collisions and comparison between the wave dynamics in single and two-fluid approximations. We derive the two-fluid MHD equations from initial three-fluid equations and solve the linearized equations in the simplest case of a homogeneous plasma. ", "conclusions": "\\begin{enumerate} \\item Frequencies and damping rates of low frequency MHD waves in the two-fluid description are similar to those obtained in the single-fluid approach. But high-frequency waves (with higher frequency than the ion-neutral collision frequency) have completely different behavior. \\item Alfv\\'en and fast magneto-acoustic waves have maximal damping rates at some frequency interval peaking at particular frequency. The peak frequency is $2.5 \\, \\nu_{in}$, where $\\nu_{in}$ is the ion-neutral collision frequency, for 50$\\%$ of neutral hydrogen. For 10$\\%$ of neutral hydrogen, the peak frequency is shifted to$10 \\, \\nu_{in}$. The damping rate is reduced for higher frequencies, therefore the damping of high-frequency Alfv\\'en waves in the solar chromosphere with realistic height profile of ionization degree needs to be revised in future. \\item There are two types of slow magneto-acoustic waves in the high-frequency part of wave spectrum: one connected with the ion-electron fluid and another with the fluid of neutrals. \\item There is no cut-off frequency of Alfv\\'en waves due to ambipolar diffusion. The cut-off frequency found in the single-fluid approach is caused by neglecting the inertial terms in the momentum equation of relative velocity. \\item The damping rate of slow magneto-acoustic waves is similar to Braginksii (\\cite{Braginskii1965}) in low plasma $\\beta$ approximation. The deviation from the Braginskii formula found by normal mode analysis in single-fluid approach (Forteza et al. \\cite{Forteza2007}) is probably caused by neglecting the inertial terms. \\end{enumerate}" }, "1101/1101.3314_arXiv.txt": { "abstract": "Cosmic reionization is expected to be complex, extended and very inhomogeneous. Existing constraints at $z\\sim6$ on the volume-averaged neutral hydrogen fraction, $\\avenf$, are highly model-dependent and controversial. Constraints at $z<6$, suggesting that the Universe is highly ionized, are also model-dependent, but more fundamentally are invalid in the context of inhomogeneous reionization. As such, it has recently been pointed out that there is no conclusive evidence that reionization has completed by $\\zfivesix$, a fact that has important ramifications on the interpretation of high-redshift observations and theoretical models. We present the first direct upper limits on $\\avenf$ at $\\zfivesix$ using the simple and robust statistic of the covering fraction of dark pixels in the Ly$\\alpha$/$\\beta$ forests of high redshift quasars. With a sample of 13 Keck ESI spectra we constrain $\\avenf\\lsim0.2$ at $5r_{\\rm N}) \\propto r_{\\rm N}^{-q}$, where $r_{\\rm N}$ is the radius of the nuclei and $q$ is the slope. We include a large number of optical observations published by ourselves and others since the comprehensive review in the {\\it Comets II} book \\citep{Lamy-chapter}, and make use of an improved fitting method. We assess the uncertainty on the CSD due to all of the unknowns and uncertainties involved (photometric uncertainty, assumed phase function, albedo and shape of the nucleus) by means of Monte Carlo simulations. In order to do this we also briefly review the current measurements of these parameters for JFCs. Our final CSD has a slope $q=1.92\\pm 0.20$ for nuclei with radius $r_{\\rm N} \\ge 1.25$ km. ", "introduction": "The size distribution of any population of solar system small bodies is of critical importance in constraining their formation and subsequent collisional evolution. \\citet{Dohnanyi69} showed that a collisionally relaxed population of self-similar bodies, with the same strength per unit mass, has a characteristic power law size distribution with a slope of 2.5. For a collisional population of gravity controlled (strengthless) bodies \\citet{OBrien+Greenberg03} demonstrated that the expected size distribution has a shallower slope, 2.04. Happily, the size distribution is also one of the more straightforward characteristics of the population to determine, as at least reasonable estimates of the size of bodies can be made with snap-shot observations. Time-series data allow better measurements of their sizes, as this removes uncertainties due to their rotational light curve. For Jupiter Family comets (JFCs) these observations are generally made when the comet is at a large distance from the Sun, and therefore more likely to be inactive, so a brightness measurement for the bare nucleus can be made. Converting the measured optical magnitude to the size of an object depends on its albedo and phase function. There are only 7 JFCs for which both of these are independently well measured (2P/Encke, 9P/Tempel~1, 10P/Tempel~2, 19P/Borrelly, 28P/Neujmin~1, 67P/Churyumov-Gerasimenko and 81P/Wild~2; three of which have a size measurement from resolved imaging by spacecraft [see Tables \\ref{beta_table} and \\ref{albedo_table} for references]). A few more comets have measurements of one or the other. Where these parameters are not known, we are forced to assume values for them; typically 4\\% for the albedo and a linear phase function of $\\beta = 0.035$ mag deg$^{-1}$. Previous measurements of the size distribution of JFCs have followed similar assumptions. The distribution is generally plotted (fig.~\\ref{ref_CSD}) as a cumulative size distribution (CSD), expressed in terms \\begin{equation} N(>r_{\\rm N}) \\propto r_{\\rm N}^{-q} \\end{equation} $N$ is the number of nuclei with radius, $r_{\\rm N}$, larger than $r_{\\rm N}$, and $q$ is the slope. A number of measurements of these distribution co-efficients have been made, generally based on snap-shot observations of a large number of nuclei. \\citet{Lowry03} estimated $q = 1.6\\pm0.1$ based on a sample of 33 comets, and \\citet{Weissman+Lowry03} updated this to $q=1.59\\pm0.03$, based on 41 JFCs with $r_{\\rm N} \\ge 1.4$ km. They chose to fit the size distribution to only those nuclei with $r_{\\rm N} \\ge 1.4$ km as the slope of the CSD is approximately constant above this radius, while below it there is a sharp cut-off. This break may imply a relative paucity of small nuclei compared with the expected number from a continuation of the power law from larger sizes; \\citet{Meech04} and \\citet{Fitzsimmons10} show that such a break is most likely real and not due to observational biases (as suggested by \\citet{Lamy-chapter}) by modelling realistic observational surveys. A more recent size distribution from \\citet{Weissman09} has a value of $q=1.94\\pm0.07$ based on 41 JFCs with $r_{\\rm N} > 1.4$ km. Another estimate comes from \\citet{FernandezJ99}, who used selected data in quality classes 1-3 (uncertainties on $m_V(1,1,0)$ up to $\\pm 1$ mag.) from the catalogue presented by \\citet{Tancredi00}, with cut-offs in both absolute magnitude and perihelion distance. The discrepancy between \\citeauthor{FernandezJ99}'s estimate of $q=2.65\\pm0.25$ and those of \\citeauthor{Lowry03} can be explained by these cut-offs (leaving only 12 comets) and the large uncertainties on magnitudes in \\citeauthor{Tancredi00}'s catalogue. \\citet{Tancredi06} presented an updated catalogue and find $q=2.7\\pm0.3$ for $r_{\\rm N} \\ge 1.5$ km. \\citet{Meech04} estimate $q=1.45\\pm0.05$ over the range $1 \\le r_{\\rm N} \\le 10$ km, and a steeper $q=1.91\\pm0.06$ in the range $2 \\le r_{\\rm N} \\le 5$ km, showing the large dependence on the choice of size range. \\citet{Hicks07} estimate $q=1.50\\pm0.08$ from Near Earth Asteroid Tracking (NEAT) survey data, although this makes use of observations of active comets and a coma subtraction technique, so is not included with the other surveys of distant inactive comets in Table \\ref{previous}. Finally, \\citet{Lamy-chapter} collated the data from most of these catalogues, together with their own unpublished results, those from \\citet{Licandro00b} and also from other papers on individual comets, to calculate $q=1.9\\pm0.3$ for JFC nuclei with $r_{\\rm N} \\ge 1.6$ km. \\begin{table} \\caption{Previous JFC size distribution estimates.} \\begin{center} \\begin{tabular}{l c c c} Reference & $N_{\\rm fit} / N_{\\rm tot}$ & $r_{\\rm N}$ range & $q$ \\\\ \\hline \\citet{FernandezJ99} & 12/64 & $2.1 - 3.3$ & $2.7 \\pm 0.3$ \\\\ \\citet{Lowry03} & 16/19 & $1.4 - 3.6$ & $1.6 \\pm 0.1$ \\\\ \\citet{Weissman+Lowry03} & 41/54 & $1.4 - 15$ & $1.59 \\pm 0.03$ \\\\ \\citet{Meech04} & 38/48 & $1.0 - 10$ & $1.45 \\pm 0.05$ \\\\ \\citet{Meech04} & 21/48 & $2.0 - 5$ & $1.91 \\pm 0.06$ \\\\ \\citet{Lamy-chapter} & 29/65 & $1.6 - 15$ & $1.9 \\pm 0.3$ \\\\ \\citet{Weissman09} & 41/67 & $1.4 - 6.0$ & $1.94 \\pm 0.07$ \\\\ \\citet{Tancredi06} & 32/72 & $1.7 - 4.5$ & $2.7 \\pm 0.3$ \\\\ \\hline \\end{tabular} \\end{center} \\label{previous} \\end{table}% These previous estimates are listed in Table \\ref{previous}, where we list the number of comets included in the fit (and the total number considered in the survey), the range in sizes over which the authors fit the linear part of the CSD, and the resulting slope $q$. The uncertainty on $q$ is that quoted by each author, and is generally the formal uncertainty from a least squares (or similar) fit to the line, despite the fact that technically one cannot use such a fitting technique as the points in a cumulative distribution are not independent. These previous works also make no attempt to assess the uncertainty on $q$ contributed by the assumptions on phase function and albedo or the uncertainty from the photometry. In this paper we present an updated size distribution based on new photometry of distant JFC nuclei, using a censored data analysis technique to produce the CSD, and make a rigourous assessment of the uncertainty on $q$ due to all the unknown factors. \\begin{figure} \\centering \\includegraphics[angle=-90,width=0.45\\textwidth]{ref_CSD.ps} \\caption{Reference CSD with our data set and the usual assumptions. This shows the normal shape of the CSD, with a linear part and a cut-off. The error bars shown are calculated using the Kaplan-Meier statistic. Each fit within a MC run produces a CSD like this with a slightly different distribution of points. } \\label{ref_CSD} \\end{figure} ", "conclusions": "\\begin{table*} \\begin{minipage}{156mm} \\caption{Inputs and results from Monte Carlo runs.} \\begin{center} \\begin{tabular}{l|c l l p{2.7cm} c c|c c c} \\hline \\# & Mag\\footnote{% Vary input magnitude within photometric error bars? } & Phase function & Albedo & Shape\\footnote{% Assumed shape distribution for nuclei. } & Fix\\footnote{% Use fixed values for $\\beta$ and $A_R$ when known for a given comet? If no, all comets have values selected from the specified distribution, even those where these values are independently constrained. } & $N$ & $q_1$\\footnote{% $q_1$ is the best fit at $r_{\\rm N} \\le$ cut-off. } & $q_2$\\footnote{% $q_2$ is the best fit at $r_{\\rm N} >$ cut-off, i.e. the quoted $q$. } & Cut-off\\\\ \\hline \\multicolumn{4}{l}{\\it Fixed assumed values:}\\\\ 1 & N & 0.035 & 0.04 & Spherical & N & 1 & 0.19 & 1.97 & 1.26 \\\\ % 2 & N & 0.035 & 0.04 & Spherical & Y & 1 & 0.19 & 2.00 & 1.26 \\\\ % \\multicolumn{4}{l}{\\it Adjust photometry:}\\\\ 3 & Y & 0.035 & 0.04 & Spherical & N & 1000 & 0.18 $\\pm$0.01 & 1.93 $\\pm$0.10 & 1.24 $\\pm$0.07 \\\\ % \\multicolumn{4}{l}{\\it Adjust phase function:}\\\\ 4 & N & 0.025 & 0.04 & Spherical & N & 1 & 0.18 & 1.94 & 1.17 \\\\ % 5 & N & 0.045 & 0.04 & Spherical & N & 1 & 0.19 & 1.98 & 1.35 \\\\ % 6 & N & 0.065 & 0.04 & Spherical & N & 1 & 0.20 & 1.99 & 1.50 \\\\ % 7 & N & 0.053 $\\pm$ 0.016 & 0.04 & Spherical & N & 1000 & 0.19 $\\pm$0.01 & 1.95 $\\pm$0.07 & 1.38 $\\pm$0.07 \\\\ % 8 & N & 0.053 $\\pm$ 0.005 & 0.04 & Spherical & N & 1000 & 0.20 $\\pm$0.00 & 1.99 $\\pm$0.02 & 1.42 $\\pm$0.03 \\\\ % 9 & N & 0.053 $\\pm$ 0.025 & 0.04 & Spherical & N & 1000 & 0.19 $\\pm$0.02 & 1.91 $\\pm$0.11 & 1.35 $\\pm$0.09 \\\\ % 10 & N & 0.053 $\\pm$ 0.016 & 0.04 & Spherical & Y & 1000 & 0.19 $\\pm$0.01 & 2.05 $\\pm$0.06 & 1.39 $\\pm$0.06 \\\\ % 11 & N & 0.010 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.15 $\\pm$0.01 & 1.72 $\\pm$0.06 & 1.02 $\\pm$0.06 \\\\ % 12 & N & 0.020 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.16 $\\pm$0.02 & 1.81 $\\pm$0.11 & 1.13 $\\pm$0.11 \\\\ % 13 & N & 0.030 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.18 $\\pm$0.01 & 1.94 $\\pm$0.09 & 1.26 $\\pm$0.09 \\\\ % 14 & N & 0.040 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.19 $\\pm$0.01 & 1.97 $\\pm$0.05 & 1.32 $\\pm$0.06 \\\\ % 15 & N & 0.050 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.19 $\\pm$0.01 & 1.97 $\\pm$0.04 & 1.38 $\\pm$0.06 \\\\ % 16 & N & 0.060 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.20 $\\pm$0.01 & 1.97 $\\pm$0.05 & 1.43 $\\pm$0.06 \\\\ % 17 & N & 0.070 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.20 $\\pm$0.01 & 1.95 $\\pm$0.05 & 1.47 $\\pm$0.06 \\\\ % 18 & N & 0.080 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.20 $\\pm$0.01 & 1.92 $\\pm$0.05 & 1.49 $\\pm$0.05 \\\\ % 19 & N & 0.090 $\\pm$ 0.010 & 0.04 & Spherical & N & 1000 & 0.20 $\\pm$0.01 & 1.90 $\\pm$0.04 & 1.54 $\\pm$0.04 \\\\ % \\multicolumn{4}{l}{\\it Adjust albedo:}\\\\ 20 & N & 0.035 & 0.044 $\\pm$ 0.013 & Spherical & N & 1000 & 0.17 $\\pm$0.02 & 1.85 $\\pm$0.15 & 1.16 $\\pm$0.10 \\\\ % 21 & N & 0.035 & 0.020 $\\pm$ 0.013 & Spherical & N & 1000 & 0.16 $\\pm$0.03 & 1.64 $\\pm$0.21 & 1.64 $\\pm$0.21 \\\\ % 22 & N & 0.035 & 0.060 $\\pm$ 0.013 & Spherical & N & 1000 & 0.18 $\\pm$0.02 & 1.89 $\\pm$0.12 & 1.01 $\\pm$0.08 \\\\ % 23 & N & 0.035 & 0.040 $\\pm$ 0.005 & Spherical & N & 1000 & 0.18 $\\pm$0.01 & 1.93 $\\pm$0.09 & 1.26 $\\pm$0.08 \\\\ % 24 & N & 0.035 & 0.040 $\\pm$ 0.020 & Spherical & N & 1000 & 0.16 $\\pm$0.03 & 1.72 $\\pm$0.19 & 1.18 $\\pm$0.13 \\\\ % 25 & N & 0.035 & 0.044 $\\pm$ 0.013 & Spherical & Y & 1000 & 0.17 $\\pm$0.02 & 1.85 $\\pm$0.14 & 1.16 $\\pm$0.10 \\\\ % \\multicolumn{4}{l}{\\it Adjust shape:}\\\\ 26 & N & 0.035 & 0.04 & Flat, 1 $\\le a/b \\le$ 3 & N & 1000 & 0.18 $\\pm$0.02 & 1.92 $\\pm$0.12 & 1.17 $\\pm$0.09 \\\\ % 27 & N & 0.035 & 0.04 & Flat, 1 $\\le a/b \\le$ 5 & N & 1000 & 0.17 $\\pm$0.02 & 1.87 $\\pm$0.15 & 1.09 $\\pm$0.09 \\\\ % 28 & N & 0.035 & 0.04 & Gaussian, 1.5 $\\pm$ 0.6 & N & 1000 & 0.18 $\\pm$0.01 & 1.95 $\\pm$0.10 & 1.22 $\\pm$0.08 \\\\ % 29 & N & 0.035 & 0.04 & Gaussian, 1.5 $\\pm$ 0.3 & N & 1000 & 0.18 $\\pm$0.01 & 1.96 $\\pm$0.08 & 1.23 $\\pm$0.07 \\\\ % 30 & N & 0.035 & 0.04 & Gaussian, 1.5 $\\pm$ 0.9 & N & 1000 & 0.18 $\\pm$0.01 & 1.94 $\\pm$0.11 & 1.20 $\\pm$0.08 \\\\ % 31 & N & 0.035 & 0.04 & Gaussian, 2 $\\pm$ 0.6 & N & 1000 & 0.18 $\\pm$0.01 & 1.91 $\\pm$0.11 & 1.18 $\\pm$0.08 \\\\ % 32 & N & 0.035 & 0.04 & Power law, $x =$ -5.6, 1 $\\le a/b \\le$ 3 & N & 1000 & 0.18 $\\pm$0.01 & 2.01 $\\pm$0.07 & 1.26 $\\pm$0.06 \\\\ % 33 & N & 0.035 & 0.04 & Power law, $x =$ -5.6, 1 $\\le a/b \\le$ 5 & N & 1000 & 0.18 $\\pm$0.01 & 2.02 $\\pm$0.07 & 1.26 $\\pm$0.07 \\\\ % 34 & N & 0.035 & 0.04 & Power law, $x =$ -4.6, 1 $\\le a/b \\le$ 5 & N & 1000 & 0.18 $\\pm$0.01 & 2.00 $\\pm$0.08 & 1.25 $\\pm$0.07 \\\\ % 35 & N & 0.035 & 0.04 & Power law, $x =$ -6.6, 1 $\\le a/b \\le$ 5 & N & 1000 & 0.18 $\\pm$0.01 & 2.02 $\\pm$0.06 & 1.27 $\\pm$0.06 \\\\ % 36 & N & 0.035 & 0.04 & Power law, $x =$ -2.6, 1 $\\le a/b \\le$ 5 & N & 1000 & 0.18 $\\pm$0.01 & 1.97 $\\pm$0.11 & 1.22 $\\pm$0.08 \\\\ % 37 & N & 0.035 & 0.04 & Power law, $x =$ -8.6, 1 $\\le a/b \\le$ 5 & N & 1000 & 0.18 $\\pm$0.00 & 2.02 $\\pm$0.04 & 1.27 $\\pm$0.05 \\\\ % \\multicolumn{4}{l}{\\it Adjust all parameters:}\\\\ 38 & Y & 0.035 $\\pm$ 0.020 & 0.040 $\\pm$ 0.040 & Gaussian, 1.5 $\\pm$ 0.6 & N & 10000 & 0.16 $\\pm$0.04 & 1.51 $\\pm$0.22 & 1.17 $\\pm$0.19 \\\\ % 39 & Y & 0.053 $\\pm$ 0.016 & 0.044 $\\pm$ 0.013 & Power law, $x =$ -5.6, 1 $\\le a/b \\le$ 3 & N & 10000 & 0.18 $\\pm$0.03 & 1.92 $\\pm$0.20 & 1.25 $\\pm$0.13 \\\\ % 40 & Y & 0.053 $\\pm$ 0.016 & 0.044 $\\pm$ 0.013 & Power law, $x =$ -5.6, 1 $\\le a/b \\le$ 3 & Y & 10000 & 0.19 $\\pm$0.02 & 2.01 $\\pm$0.20 & 1.28 $\\pm$0.12 \\\\ % \\hline \\end{tabular} \\end{center} \\end{minipage} \\label{big_MCrun_table} \\end{table*}% \\begin{figure} \\centering \\includegraphics[angle=-90,width=0.45\\textwidth]{density_plot.ps} \\caption{Probability map for MC run 39, which gives our final result by allowing all parameters to vary. The shading shows the average shape of the 10,000 CSDs: darker areas are the bins in $R$--$P(r_{\\rm N} > R)$ space where the majority of CSDs passed through, lighter areas show the outlying areas explored at the ends of the distributions. A dashed line showing the average slope from this run is over-plotted.} \\label{fig:density} \\end{figure} Table \\ref{big_MCrun_table} shows input parameters and results for all the MC runs, including all of the ones varying only a single parameter at a time described in the previous section and also runs allowing all parameters to vary. This shows that, relative to our reference CSD with $q=1.97$, the largest change in slope is found when allowing the albedo to vary. This is mostly due to the very shallow slope that results from any albedo distribution that includes exceptionally dark ($A_R \\approx 0$) nuclei, which are unlikely to be realistic. In any case, this demonstrates the importance of better constraining the albedo distribution of comets, an important result that the SEPPCoN survey will provide \\citep{FernandezY08,Lowry10}. The only other MC runs with a large difference from the reference ($\\Delta q \\ge 0.1$) are the extreme shape distributions, which are also unlikely to represent reality. We can conclude from this that the uncertainty on the input parameter distributions actually has only a small effect, as the variation in slopes always falls within the typical uncertainty found when varying a single parameter at a time. Allowing all parameters to vary within the best-fit distributions gives a final value of $q = 1.92 \\pm 0.20$. This value is remarkably close to our reference value, again implying that the combined effect of the various uncertain parameters is actually relatively small. This helps to validate the assumptions made in previous CSD estimates, but the real significance of this work is that the uncertainty on this value has been rigourously determined by including all individual uncertainties on each radius. Figure \\ref{fig:density} shows how the CSDs varied in this MC simulation. Comparing this with the previous CSD estimates listed in Table~\\ref{previous} shows that the value we find for $q$ is in a similar range to recent results, if slightly steeper than the consensus value, and when we consider the true uncertainty we find the majority of results are encompassed within $2\\sigma$. While previous works assumed constant values for $A_R$ and $\\beta$ and did not consider the uncertainties on these, we show that the CSD slope including all sources of uncertainty agrees with these earlier results. It is worth noting that our testing of various distributions for the assumed parameters never produced a slope as steep as the considerably larger $q$ found by \\citet{FernandezJ99} and \\citet{Tancredi06}, even for extreme distributions we regard as unlikely to be realistic. We suspect that the difference is due to the higher accuracy of the photometric measurements used in our catalogue, all of which come from large aperture telescopes. It is possible that a similar MC analysis of the uncertainty on $q$ due to the photometric uncertainties in the \\citet{Tancredi06} catalogue would give a large error bar, making it consistent with the other results, however this catalogue does not include all the necessary information (in its currently published form) to apply our technique. The CSD slope we find is very close to the theoretical value of $q=2.04$ found by \\citet{OBrien+Greenberg03} for a collisionally relaxed population of strengthless bodies. However, the large uncertainty we find means that we cannot claim that this is evidence that JFC nuclei came from collisional disruption of larger bodies; we can simply state that the CSD slope is consistent with that interpretation. It is important to remember that the activity of comets alters their sizes as they lose significant amounts of material with each perihelion passage, and therefore we would not necessarily expect to see a CSD that matched a collisional one, even if the comets are collisional fragments. Both `normal' cometary activity and major mass loss events (nucleus splitting or outbursts) are highly variable between comets and from one orbit to another, so the change in size due to mass loss from a nucleus is a highly complicated process. More detailed theoretical models describing the expected CSD of nuclei for various formation (and evolution) scenarios are required for comparison with observations, however the large uncertainty on $q$ found by this paper shows that we must first provide better observational constraints on all physical properties of nuclei (albedos, phase functions and shape models) to produce observed CSDs with sufficient accuracy to compare with these models." }, "1101/1101.0701_arXiv.txt": { "abstract": "We consider a set of $M$ images, whose pixel intensities at a common point can be treated as the components of a $M$-dimensional vector. We are interested in the estimation of the modulus of such a vector associated to a compact source. For instance, the detection/estimation of the polarized signal of compact sources immersed in a noisy background is relevant in some fields like Astrophysics. We develop two different techniques, one based on the Maximum Likelihood Estimator (MLE) applied to the modulus distribution, the modulus filter (ModF) and other based on prefiltering the components before fusion, the filtered fusion (FF), to deal with this problem. We present both methods in the general case of $M$ images and apply them to the particular case of three images (linear plus circular polarization). Numerical simulations have been performed to test these filters considering polarized compact sources immersed in stationary noise. The FF performs better than the ModF in terms of errors in the estimated amplitude and position of the source, especially in the low signal-to-noise case. We also compare both methods with the direct application of a matched filter (MF) on the polarization data. This last technique is clearly outperformed by the new methods. ", "introduction": "The detection and estimation of the intensity of compact objects --i.e. signals with a compact support either in time or space domains-- embedded in a background plus instrumental noise is a problem of interest in many different areas of science and engineering. A classic example is the detection of point-like extragalactic objects such as galaxies and galaxy clusters in sub-millimetric Astronomy. Regarding this particular field of interest, different techniques have proven useful in the literature. Some of the proposed techniques are frequentist, such as the standard matched filter \\cite{MF_radio92}, the matched multifilter \\cite{herr02a} or the recently developed matched matrix filters \\cite{herranz08a}. Other frequentist techniques include continuous wavelets like the standard Mexican Hat \\cite{wsphere} and other members of its family \\cite{MHW206} and, more generally, filters based on the Neyman-Pearson approach using the distribution of maxima \\cite{can05b}. All these filters have been applied to real data of the Cosmic Microwave Background (CMB), like those obtained by the WMAP satellite \\cite{NEWPS07} and CMB simulated data \\cite{challenge08} for the experiment on board the \\emph{Planck} satellite \\cite{planck_tauber05}, which has recently started operations. Besides, Bayesian methods have also been recently developed \\cite{psnakesI}. Although we have chosen the particular case of sub-millimetric Astronomy as a means to illustrate the problem of compact source detection, the methods listed above are totally general and can be used to any analogous image processing problem. In most cases one is interested only in the intensity of the compact sources. In other cases, however, there are other properties of the signal that may be of interest. Such is the case, for example, of sources that emit electromagnetic radiation that is at least partially polarized. Polarization of light is conventionally described in terms of the Stokes parameters $Q$, $U$ and $V$. Let us consider a monochromatic, plane electromagnetic wave propagating in the $z$-direction. The components of the wave's electric field vector at a given point in the space can be written as \\begin{eqnarray} E_x & = & a_x(t) \\cos \\left[ \\omega_0 t - \\theta_x(t) \\right] \\nonumber \\\\ E_y & = & a_y(t) \\cos \\left[ \\omega_0 t - \\theta_y(t) \\right]. \\end{eqnarray} \\noindent If some correlation exists between the two components in the previous equation, then the wave is \\emph{polarized}. The Stokes parameters are defined as the time averages \\begin{eqnarray} I & \\equiv & \\langle a_x^2 \\rangle + \\langle a_y^2 \\rangle , \\nonumber \\\\ Q & \\equiv & \\langle a_x^2 \\rangle - \\langle a_y^2 \\rangle , \\nonumber \\\\ U & \\equiv & \\langle 2 a_x a_y \\cos \\left( \\theta_x - \\theta_y \\right) \\rangle , \\nonumber \\\\ V & \\equiv & \\langle 2 a_x a_y \\sin \\left( \\theta_x - \\theta_y \\right) \\rangle . \\end{eqnarray} The parameter $I$ gives the intensity of the radiation which is always positive, while the other three parameters define the polarization state of the wave and can have either sign. $Q$ and $U$ are the linear polarization parameters and $V$ indicates the circular polarization of the wave. Unpolarized radiation is described by $Q=U=V=0$. While the total intensity of the wave is independent of the orientation of the $x$ and $y$ axes, the values of the other Stokes parameters are not invariant with respect to changes of the orientation of the receivers. On the other hand, the \\emph{total polarization}, defined as \\begin{equation} \\label{eq:P} P \\equiv \\sqrt{Q^2+U^2+V^2}, \\end{equation} \\noindent is invariant with respect to the relative orientation of the receivers and the direction of the incoming light, and therefore it is a quantity with a clear physical meaning. For the case of purely linear polarization, $V=0$ and the previous expression reduces to $P \\equiv (Q^2+U^2)^{1/2}$. Note that in order to get $P$ from its components $Q$, $U$, $V$ it is necessary to perform a non-linear operation. Although strictly speaking $Q$, $U$ and $V$ are the components of a tensor, the invariant combination (\\ref{eq:P}) gives a quantity that can be seen as the modulus of a vector. In this paper we will introduce a methodology that can be applied to any problem in which a set of images contain signals whose individual intensities can be considered as the components of a vector, but where the quantity of interest is the modulus of such a vector. For illustrative purposes, throughout the paper we will use as an example the case of light polarization, but the methods we will introduce are not limited to the example. Another possible application could be the determination of the modulus of a complex-valued signal. For example, in \\cite{oberto} the case of the estimation of the modulus of a complex-valued quantity whose components follow a Gaussian distribution was addressed. The techniques presented in their paper are related to our methods, but restricted to the two-dimensional case. In \\cite{simmons85}, four methods: MLE estimator, median estimator, mean estimator and Wardle and Kronberg's estimator \\cite{wardle74} are applied to the estimation of polarization (two-dimensional case). However, these methods are applied in the cited paper to a single data and cannot be, except for the MLE, generalized to the detection of a signal with a given profile in a pixelized image, as considered in our paper. This is due to the fact that these methods lead to a system of possibly incompatible equations. Going back to the case of point-like extragalactic objects in sub-millimetric Astronomy, the polarization of the sources plays and important role in the cosmological tests derived from CMB observations. Standard cosmological models predict that CMB radiation is linearly polarized. However, some cosmological models predict in addition a possible circular polarization of CMB radiation \\cite{cooray03}. In order to better constrain the cosmological model with observations, it is crucial to determine the degree of polarization of not only the CMB radiation but also the other astrophysical sources whose signals are mixed with it. For an excellent review on CMB polarization, see \\cite{kamion97}. We have treated the application to the detection of linearly polarized sources in CMB maps elsewhere \\cite{paco09}, but it is known that extragalactic radio sources can also indeed show circular polarization \\cite{kirk06}. Besides, circular polarization occurs in many other astrophysical areas, from Solar Physics \\cite{tris07} to interstellar medium \\cite{cox07}, just to put a few examples. Therefore, in this paper we aim to address the general case of linear plus circular polarization of compact sources. In the general case we have three images $Q$, $U$ and $V$. Different approaches can be used to deal with detection/estimation of point-like sources embedded in a noisy background. On the one hand, one can try to get the source polarization amplitude $A$ directly on the $P$-map. In this approach, we will consider one filter, obtained through the MLE applied to the modulus distribution (ModF). On the other hand, we can operate with three matched filters, each one on $Q$, $U$ and $V$ followed by a quadratic fusion and square root (FF). We are trying to compare the performance of the two techniques for estimating the position and polarization amplitude of a compact source. Of course, in the case we have only the map of the modulus of a vector and the components are unknown, the FF cannot be applied but we can still use the ModF. Finally, we also apply a matched filter (MF) directly on the $P$-map and compare this simple method with the two new methods introduced in this paper. In Section~\\ref{sec:methodology} we will develop the methodology for the case of $M$ images, because of the possible interesting applications to the general $M$-dimensional case and in particular to the 3-dimensional case (polarization). In Section~\\ref{sec:simresults} we will show the results when applying these techniques to numerical simulations of images of $Q$, $U$ and $V$ that are relevant for the detection of compact polarized sources in Astrophysics. Finally, in Section~\\ref{sec:conclusions} we give the main conclusions. ", "conclusions": "\\label{sec:conclusions} In this paper, we deal with the detection and estimation of the modulus of a vector, a problem of great interest in general and in particular in astrophysics when we consider the polarization of the cosmic microwave background (CMB), extragalactic sources, the interestellar medium or the Sun. The total polarization intensity $P$ is defined as $P\\equiv (Q^2+U^2+V^2)^{1/2}$, where $Q$ , $U$ ( linear polarization) and $V$ (circular polarization) are the Stokes parameters. We consider the case of images in $Q$, $U$ and $V$ consisting of a compact source with a profile $\\tau(\\vec{x})$ immersed in Gaussian uncorrelated noise. We intend to detect the source and estimate its polarization amplitude by using three different methods, two new and a standard matched filter. a) a filter operating on the modulus $P$ (ModF) and based on the maximization of the corresponding log-likelihood and b) a filtered fusion (FF) procedure, i.e. the application of the matched filter (MF) on the images of $Q$, $U$ and $V$ and the combination of the corresponding estimates by making the non-linear fusion $\\hat{A}\\equiv (Q_{MF}^2+U_{MF}^2+V_{MF}^2)^{1/2}$ and c) a matched filter (MF) operating on the $P$-map. We present the three filters in Section~\\ref{sec:methodology} for the general case of the modulus of a $M$-vector, and for the case of three-dimensional vectors, deriving the corresponding expressions , eqs. (\\ref{eq:estimator_A}) and (\\ref{eq:estimator_A3}), for the estimation of the modulus of a vector with the ModF. Since we are interested in the detection of polarized signals in Astrophysics, we have only considered the three-dimensional vector case in our simulations. Note, however that the methods can be applied in the general case of combination of $M$ images to obtain the image of the modulus of a vector. We have compared the performance of the filters when applied to simulated images consisting of the $Q$, $U$ and $V$ components of a Gaussian-shaped signal with different intensities plus Gaussian uncorrelated stationary noise. For each simulation and for the two methods, we have estimated the source amplitude and position. Besides, we have calculated the detection power for a fixed significance $\\alpha=0.05$ and the errors of the estimated amplitude and position. We find that the performance of the FF is the best for low signal-to-noise, the ModF performs like the FF for $A\\geq 3$. The MF produces very high errors in the polarization estimation, making this filter unsuitable for the treated problem. We want to point out the good performance of the ModF, which could be an interesting alternative to the FF when we have an image of the modulus, but we do not know the components of the vector. Finally, it would be interesting to generalize our methods to the case of images in which the noise at different pixels is not uniform or, more generally, is correlated. This could be important for CMB polarization observations. We leave this problem for further work. \\begin{table} \\caption{First column: triplets of values of $A_Q$, $A_U$ and $A_V$ and the corresponding value of $A$ in $\\sigma$ units. Second and third columns: detection power for the ModF and FF (percentage) for a significance level $\\alpha=0.05$. Fourth and fifth columns: flux relative errors for the two filters. Sixth and seventh columns: absolute value of the relative error. Eighth and ninth columns: position errors in numbers of pixels. In all the cases we have performed the average of 100 simulations.} \\label{tb:table1} \\centering \\scriptsize \\begin{tabular}{|c|c|c|c|c|c|c|c|c|} \\hline & & & & & & & & \\\\ \\textbf{($A_Q,A_U,A_V;A$)}& \\textbf{pow$_{ModF}$}& \\textbf{pow$_{FF}$} & \\textbf{err$_{ModF}$}& \\textbf{err$_{FF}$} & \\textbf{$|$err$|$$_{ModF}$} & \\textbf{$|$err$|$$_{FF}$} & \\textbf{pos$_{ModF}$} & \\textbf{pos$_{FF}$} \\\\ & & & & & & & & \\\\ \\hline \\textbf{(0.00,0.00,0.50;0.50)} & 6 & 12 & 3.35 & 1.58 & 3.35 & 1.58 & 4.36 & 3.83 \\\\ \\textbf{(0.00,0.50,0.50;0.71)} & 1 & 22 & 1.90 & 0.81 & 1.90 & 0.81 & 11.31 & 2.45 \\\\ \\textbf{(0.50,0.50,0.50;0.87)} & 7 & 41 & 1.34 & 0.52 & 1.34 & 0.52 & 6.68 & 2.19 \\\\ \\textbf{(0.00,0.00,1.00;1.00)} & 5 & 47 & 1.18 & 0.37 & 1.18 & 0.37 & 3.80 & 1.99 \\\\ \\textbf{(0.00,0.50,1.00;1.12)} & 6 & 71 & 0.91 & 0.24 & 0.91 & 0.24 & 3.44 & 1.23 \\\\ \\textbf{(0.50,0.50,1.00;1.22)} & 6 & 76 & 0.71 & 0.16 & 0.71 & 0.17 & 3.11 & 1.14 \\\\ \\textbf{(0.00,1.00,1.00;1.41)} & 17 & 89 & 0.49 & 0.10 & 0.49 & 0.14 & 3.24 & 0.87 \\\\ \\textbf{(0.00,0.00,1.50;1.50)} & 24 & 96 & 0.45 & 0.09 & 0.45 & 0.15 & 2.34 & 0.70 \\\\ \\textbf{(0.50,1.00,1.00;1.50)} & 25 & 95 & 0.44 & 0.09 & 0.44 & 0.15 & 2.17 & 0.65 \\\\ \\textbf{(0.00,0.50,1.50;1.58)} & 27 & 96 & 0.34 & 0.06 & 0.34 & 0.14 & 1.45 & 0.62 \\\\ \\textbf{(0.50,0.50,1.50;1.66)} & 27 & 97 & 0.32 & 0.04 & 0.32 & 0.13 & 1.37 & 0.61 \\\\ \\textbf{(1.00,1.00,1.00;1.73)} & 43 & 100 & 0.25 & 0.03 & 0.25 & 0.11 & 1.53 & 0.70 \\\\ \\textbf{(0.00,1.00,1.50;1.80)} & 41 & 99 & 0.25 & 0.05 & 0.25 & 0.14 & 1.39 & 0.43 \\\\ \\textbf{(0.50,1.00,1.50;1.87)} & 50 & 100 & 0.18 & 0.04 & 0.18 & 0.12 & 1.33 & 0.53 \\\\ \\textbf{(0.00,0.00,2.00;2.00)} & 59 & 100 & 0.17 & 0.03 & 0.17 & 0.12 & 0.83 & 0.30 \\\\ \\textbf{(0.00,0.50,2.00;2.06)} & 63 & 100 & 0.12 & 0.02 & 0.13 & 0.12 & 0.72 & 0.40 \\\\ \\textbf{(1.00,1.00,1.50;2.06)} & 65 & 100 & 0.11 & 0.03 & 0.11 & 0.10 & 0.92 & 0.41 \\\\ \\textbf{(0.00,1.50,1.50;2.12)} & 76 & 100 & 0.11 & 0.04 & 0.12 & 0.10 & 0.61 & 0.33 \\\\ \\textbf{(0.50,0.50,2.00;2.12)} & 74 & 100 & 0.10 & 0.03 & 0.12 & 0.10 & 0.66 & 0.35 \\\\ \\textbf{(0.50,1.50,1.50;2.18)} & 76 & 100 & 0.09 & 0.02 & 0.11 & 0.11 & 0.73 & 0.33 \\\\ \\textbf{(0.00,1.00,2.00;2.24)} & 81 & 100 & 0.06 & 0.02 & 0.10 & 0.10 & 0.62 & 0.33 \\\\ \\textbf{(0.50,1.00,2.00;2.29)} & 81 & 100 & 0.04 & 0.00 & 0.09 & 0.10 & 0.79 & 0.31 \\\\ \\textbf{(1.00,1.50,1.50;2.35)} & 86 & 100 & 0.04 & 0.02 & 0.11 & 0.10 & 0.90 & 0.22 \\\\ \\textbf{(1.00,1.00,2.00;2.45)} & 85 & 100 & 0.03 & 0.01 & 0.09 & 0.09 & 0.56 & 0.24 \\\\ \\textbf{(0.00,0.00,2.50;2.50)} & 96 & 100 & 0.04 & 0.04 & 0.10 & 0.09 & 0.51 & 0.18 \\\\ \\textbf{(0.00,1.50,2.00;2.50)} & 96 & 100 & 0.03 & 0.02 & 0.10 & 0.09 & 0.50 & 0.25 \\\\ \\textbf{(0.00,0.50,2.50;2.55)} & 95 & 100 & 0.04 & 0.02 & 0.11 & 0.09 & 0.48 & 0.24 \\\\ \\textbf{(0.50,1.50,2.00;2.55)} & 94 & 100 & 0.03 & 0.00 & 0.10 & 0.10 & 0.68 & 0.20 \\\\ \\textbf{(0.50,0.50,2.50;2.60)} & 99 & 100 & 0.03 & 0.02 & 0.10 & 0.09 & 0.58 & 0.17 \\\\ \\textbf{(1.50,1.50,1.50;2.60)} & 95 & 100 & 0.03 & 0.02 & 0.11 & 0.10 & 0.62 & 0.17 \\\\ \\textbf{(1.00,1.50,2.00;2.69)} & 99 & 100 & 0.02 & 0.03 & 0.10 & 0.08 & 0.49 & 0.16 \\\\ \\textbf{(0.00,1.00,2.50;2.69)} & 97 & 100 & 0.00 & 0.01 & 0.09 & 0.08 & 0.50 & 0.14 \\\\ \\textbf{(0.50,1.00,2.50;2.74)} & 98 & 100 & 0.01 & 0.02 & 0.10 & 0.09 & 0.38 & 0.15 \\\\ \\textbf{(0.00,2.00,2.00;2.83)} & 100 & 100 & 0.00 & 0.01 & 0.09 & 0.06 & 0.58 & 0.15 \\\\ \\textbf{(1.00,1.00,2.50;2.87)} & 99 & 100 & -0.01 & 0.00 & 0.10 & 0.08 & 0.44 & 0.17 \\\\ \\textbf{(0.50,2.00,2.00;2.87)} & 98 & 100 & 0.02 & 0.02 & 0.11 & 0.09 & 0.41 & 0.15 \\\\ \\textbf{(0.00,1.50,2.50;2.92)} & 100 & 100 & -0.01 & 0.01 & 0.08 & 0.06 & 0.32 & 0.09 \\\\ \\textbf{(1.50,1.50,2.00;2.92)} & 100 & 100 & 0.00 & 0.00 & 0.08 & 0.08 & 0.35 & 0.10 \\\\ \\textbf{(0.50,1.50,2.50;2.96)} & 100 & 100 & 0.01 & 0.01 & 0.09 & 0.08 & 0.33 & 0.11 \\\\ \\textbf{(1.00,2.00,2.00;3.00)} & 100 & 100 & 0.01 & 0.02 & 0.08 & 0.07 & 0.27 & 0.07 \\\\ \\textbf{(1.00,1.50,2.50;3.08)} & 100 & 100 & 0.01 & 0.01 & 0.08 & 0.06 & 0.27 & 0.05 \\\\ \\textbf{(0.00,2.00,2.50;3.20)} & 100 & 100 & 0.00 & 0.00 & 0.08 & 0.07 & 0.33 & 0.12 \\\\ \\hline \\end{tabular} \\end{table}" }, "1101/1101.2368_arXiv.txt": { "abstract": "We present new results on the physical nature of infrared-luminous sources at $0.5 0.9$ mJy. We find many ($\\sim 60\\%$) of our sources to possess an important bulge and/or central point source component, most of which reveal additional underlying structures after subtraction of a best-fit sersic (or sersic+PSF) profile. Based on visual inspection of the NIC2 images and their residuals, we estimate that $\\sim 80\\%$ of all our sources are mergers. We calculate lower and upper limits on the merger fraction to be 62\\% and 91\\% respectively. At $z < 1.5$, we observe objects in early (pre-coalescence) merging stages to be mostly disk and star formation dominated, while we find mergers to be mainly bulge-dominated and AGN-starburst composites during coalescence and then AGN-dominated in late stages. This is analogous to what is observed in local ULIRGs. At $z \\ge 1.5$, we find a dramatic rise in the number of objects in pre-coalescence phases of merging, despite an increase in the preponderance of AGN signatures in their mid-IR spectra and luminosities above $10^{12.5} L_{\\odot}$. We further find the majority of mergers at those redshifts to retain a disk-dominated profile during coalescence. We conclude that, albeit still driven by mergers, these high-$z$ ULIRGs are substantially different in nature from their local counterparts and speculate that this is likely due to their higher gas content. Finally, we observe obscured ($\\tau_{9.7\\mu m} > 3.36$) quasars to live in faint and compact hosts and show that these are likely high-redshift analogs of local dense-core mergers. We find late-stage mergers to show predominantly unobscured AGN spectra, but do not observe other morphological classes to occupy any one specific region in the $\\tau_{9.7\\mu m}$ vs. PAH equivalent width (or Spoon) diagram. This suggests a high degree of variation in the PAH emission and silicate absorption properties of these mergers, and possibly throughout the merging process itself. ", "introduction": "Since their discovery in the {\\em IRAS} all-sky survey more than 25 years ago, ultra-luminous infrared galaxies (ULIRGs, $L_{IR} > 10^{12} L{\\odot}$) have been thought to represent a key evolutionary link between normal galaxies and quasars. Early evidence suggested that they were young quasars fed and obscured by large amounts of gas and dust funneled towards the center of a merger remnant of two gas-rich spirals \\citep{Sanders88}. Intense star formation rather than black hole accretion was later shown to be the primary energy source of most ULIRGs \\citep{Rigopoulou96, Genzel98}, but black hole accretion has remained the dominant mechanism at higher luminosities \\citep[$L_{IR} \\gtrsim 10^{12.3} L_{\\odot}$;][]{Lutz98,Veilleux99,Tran01,Farrah03,Veilleux09a}. Ground-based as well as high resolution {\\em HST} imaging, meanwhile, revealed that more than 95\\% of ULIRGs originate in a merger, but that the ULIRG phase can also appear much before final coalescence \\citep{Murphy96, Veilleux02}. Simulations have yielded strong support to the picture of quasars originating in the merger of two gas-rich spirals after a phase of intense star formation and rapid black hole growth giving rise to the ULIRG phenomenon. \\citet{Barnes91} and \\citet{Mihos94} showed that the merger of two equal-mass disk galaxies can rapidly dissipate angular momentum, causing the gas to fall to the center of the galaxy and create a starburst of ULIRG proportions. They also demonstrated how the stellar component formed tails and streams much like the ones observed in ULIRGs. Refinements in hydrodynamical simulations and the introduction of AGN feedback by \\citet{Springel05} then showed how the AGN, once triggered, can expel remaining gas, quench star formation and become a true quasar within a remnant elliptical galaxy \\citep{diMatteo05,Hopkins08}. Finally, the addition of radiative transfer confirmed the exceptional infrared luminosity associated with the whole event, and in particular with the final coalescence \\citep{Jonsson06,Li08,Younger09,Narayanan09}. Although ULIRGs and quasars are extremely rare locally \\citep{Soifer87}, observations at sub-millimeter, mid-IR, optical and X-ray wavelengths have demonstrated that their number and luminosity densities increase rapidly with redshift \\citep{Chapman05, LeFloch05, Richards06, Hasinger05}. In particular, the advent of the {\\em Spitzer Space Telescope} has enabled sensitive and fast imaging at 24\\um\\ with MIPS, yielding the detection of a large number of infrared-luminous galaxies at $z<3$ \\citep[e.g.][]{Perez-Gonzalez05}. With such a rise in prominence, it becomes important for our understanding of galaxy/quasar evolution to ask whether, or how many of, these numerous high-$z$ ULIRGs are triggered through the same physical mechanisms as their low-redshift counterparts, and whether they also represent a transition towards quasars. Given the higher gas fractions \\citep{Noterdaeme09}, star formations rates \\citep{Hopkins06} as well as specific star formation rates \\citep{Zamojski07, Brinchmann00} of the overall galaxy population at these redshifts, quiescent star formation is expected to contribute increasingly more to the infrared luminosity of galaxies \\citep{Hopkins10a}. Certainly at $z \\sim 2$, some ULIRGs have been found to exhibit disk-like kinematics \\citep{Forster-Schreiber09, Carilli10, Bothwell10}. The answer as to what extent mergers are still necessary to explain the origin of the infrared-luminous population at higher redshifts is, therefore, unclear. Meanwhile, \\iso studies of local ULIRGs have demonstrated that mid-IR spectroscopy, through the resolution of PAH emission complexes and measurement of their strength relative to the underlying continuum, is the most effective single tool for identifying which of star formation or AGN activity is responsible for the observed mid-IR radiation of an object. The launch of the {\\em Spitzer Space Telescope} with its infrared spectrograph ({\\em IRS}) sensitive to fluxes of $S_{24\\mu m} \\gtrsim 1$ mJy has, thus, brought a flurry of mid-IR spectroscopic surveys of bright, high-redshift $24\\mu m$ galaxies aimed at addressing the origin of their infrared luminosity \\citep{Houck05, Yan05, Weedman06, Yan07, Sajina07, Farrah08, Dasyra09, Desai09}. The results of these efforts have demonstrated that, unlike sub-millimeter galaxies that are primarily powered by star formation \\citep{Pope08,Menendez09}, $24\\mu m$-selected objects appear to be more analogous to local ULIRGs in that they display both types of spectra (as well as various combinations thereof). Also in analogy to local ULIRGs, the AGN contribution to their mid-IR flux increases with total IR-luminosity \\citep{Sajina07, Dey08, Desai09}. Our group carried out two such mid-IR spectroscopic programs of $24\\mu m$-bright galaxies, the first of which yielded spectra for 52 objects at $z$\\,$\\simgt$\\,1\\,--\\,3 \\citep{Yan07,Sajina07,Sajina08}, while the second targeted 150 objects spanning a redshift range of $z$\\,$\\sim$\\,0.3\\,--\\,2.5 and peaking at $z=1$ \\citep{Dasyra09}. Both were conducted in the \\spitz Extragalactic First Look Survey (XFLS). Analysis of this combined data has revealed the presence of an obscured AGN in $\\simgt$\\,75\\% of our objects. The answer as to whether the infrared luminosity of these $24\\mu m$-bright galaxies is due to mergers, however, has, thus far, not been conclusive. Initial morphological studies have yielded mixed results \\citep{Dasyra08, Bussmann09, Melbourne09}, but part of this variation could be due to small sample sizes and selection criteria. In an effort to address this question more fully, we have obtained HST/NICMOS imaging of a sample of 135 bright high-redshift $24\\mu m$-selected galaxies that combines the previously published data for 33 sources from our first program \\citep{Dasyra08} to new data for 102 sources from our second program. In this paper, we expand the analysis of \\citet{Dasyra08} and perform a systematic search of merger signatures by uncovering underlying structures, resulting in the discovery that $\\sim 80\\%$ of our objects are fueled by either an ongoing or recent ($\\lesssim 0.5$ Gyr) merger event. We also artificially redshift local ULIRGs to quantify the detectability of merging signatures at high redshift and to create a comparison sample for our sources. We examine similarities and differences with local ULIRGs. We then combine our morphological results with our IRS spectra \\citep{Sajina07, Dasyra09} and the SED analysis of \\citet{Sajina08} and Sajina et al. (in preparation) to investigate the link between morphology/merging and, both, the relative strength of the AGN and starburst components at infrared wavelengths, as well as the degree of obscuration. We discuss our results in the context of other LIRG and ULIRG samples at both high and low redshift. Finally, we confront our observations with simulations, and discuss their consequences for our understanding of galaxy evolution. The paper is organized such that we first describe our sample and analysis in \\S 2, which we follow by a discussion of our simulated observations of redshifted local galaxies in \\S 3. We present our results in \\S 4 and compare them to that of other ULIRG samples in \\S 5. We end section 5 by discussing implications for the theory of galaxy evolution, and conclude with a summary (section~6). We use a $\\Lambda$CDM cosmology with $H_{0} = 70 \\mbox{ km s}^{-1} \\mbox{ Mpc}^{-1}$, $\\Omega_{m} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ", "conclusions": "} We have shown that a high fraction, about 80\\%, of bright ($S(24\\mu m) > 0.9$ mJy), high-redshift $24\\mu m$-sources are likely to be ongoing mergers in various stages of the process. We further demonstrated that although our sources show a broad range of morphologies and populate the entire merging sequence, the relative number of early mergers increases substantially at $z \\ge 1.5$. We showed that star formation activity as probed by the equivalent width of PAH features, decreases, on average, along the merging sequence to the profit of black-hole accretion. We finally demonstrated that obscured quasars, in our sample, live in \\textquotedblleft faint \\& compact\\textquotedblright galaxies. We mirrored our analysis on that of \\citet{Veilleux02, Veilleux06, Veilleux09b} on the 1-Jy sample, and simulated high-redshift NICMOS observations of local ULIRGs in order to address the question as to whether our sources are high-redshift analogs of local ULIRGs. To a large extent, we presented our comparison of the two samples concurrently with our results in section~4, but discuss that question in a broader context in this section. The 1-Jy sample, however, is arguably not the best point of comparison, since the sample is selected at a substantially different rest-frame wavelength. Furthermore, comparison with local galaxies gives us only part of the picture; comparison with other high-redshift samples is also highly desirable. We propose to carry such comparisons in this section. We will discuss our results in the context of five types of infrared-luminous objects: local ULIRGs (from the 1-Jy sample), other bright $24\\mu m$-selected samples at high-redshift, sub-mm galaxies, and finally $z \\sim 1$, $70\\mu m$-selected objects. Lastly, we compare our results with theoretical predictions and discuss implications for our understanding of galaxy evolution. \\subsection{Are bright high-redshift $24\\mu m$ galaxies analogs of local ULIRGs? \\label{sec:local_analogs}} Our observations suggest that, up to $z=1.5$, they are analogs of, not only local ULIRGs, but also top-end LIRGs ($L_{IR} > 10^{11.5} L_{\\odot}$). This is evidenced by their large infrared luminosities ($L_{3-1000\\mu m} \\sim 10^{11.5}\\textendash 10^{12.5} L_{\\odot}$) and their high merger fraction (up to 83\\% at $z < 1.5$), which are the two defining characteristics of high-end LIRGs and ULIRGs locally \\citep{Murphy96,Veilleux02,Ishida04}. Our analysis, however, has revealed that our $24\\mu m$-selected sample at $z < 1.5$ tends to draw from objects in more advanced stages of merging, likely owing to the shorter wavelength selection. When compared to redshifted galaxies from the 1-Jy sample, our objects also tend to show stronger, brighter merger signatures such as more visible tidal features and a higher incidence of detected residuals. Spiral galaxies make for only 17\\% of our objects at $z < 1.5$, thus rejecting the hypothesis that our sources are scaled up versions of lower luminosity local LIRGs \\citep[which have disk morphologies in $\\gtrsim 40\\%$ of cases;][]{Wang06}, even though some of our spirals do show slightly elevated IR-luminosities (up to $1.6 \\times 10^{12} L_{\\odot}$). At redshifts above 1.5, we observe a dramatic rise in the number of early mergers among our bright $24\\mu m$ sources despite the fact that we are probing, at those redshift, only extremely high luminosities ($L_{IR} > 10^{12.5} L_{\\odot}$), and despite the fact that $\\gtrsim 90\\%$ of these objects have PAH EWs indicative of the presence of an AGN ($\\mbox{EW}_{7.7\\mu m} < 1.2$). Locally, this combination of high luminosity and low EWs can only be found among more advanced mergers. These early mergers that represent, at $z \\ge 1.5$, half of our sample, thus form a new type of sources. Morphologically, albeit very messy, they do not appear fundamentally different from lower redshift luminous mergers at similar stages. Their PAH EWs, however, indicate that they are primarily AGN-dominated or composite systems. It is, therefore, an earlier triggering, and sustained fueling, of a strong obscured AGN that distinguishes them from local ULIRGs. As we discuss in section~\\ref{sec:theory}, we think this might be the consequence of the elevated gas content of these objects in comparison to that of local ULIRGs \\citep{Yan10}. The infrared-luminous aspect of galaxy evolution thus appears to proceed differently at high redshifts, in a way that is not well represented by local ULIRGs. \\subsection{Comparison with other samples of $24\\mu m$-selected galaxies} \\citet{Dasyra08} presented a first analysis of a $z > 1.5$ subset of the sample used in this paper that was selected from objects in our first \\spitz/IRS program. By looking for merging pairs, they were able to place a lower limit on the number of interactions in $z\\sim 2$ bright $24\\mu m$ galaxies at 52\\%. We have expanded those results by performing a systematic search for merging signatures and find that the other half of our $z > 1.5$ objects all show signs of possibly being later-stage mergers (from coalescence onward). 60\\% of those objects are identified with high confidence, being of confidence level ~4 or less (cf. \\S~2.5.2). \\citet{Dasyra08} further found hints that these high-redshift bright $24\\mu m$ sources might be disk-dominated rather than bulge-dominated like local ULIRGs. Our analysis reveals that this is indeed the case and that two thirds of our objects at $z>1.5$, in fact, possess disk-dominated profiles ($B/D < 0.5$ or $n<2.3$). This can be largely attributed to the high fraction of early mergers at those redshifts, but many, more advanced, mergers also show rather low $B/D$ ratios or sersic indices. We speculate that this is due to a combination of two effects: high gas fractions in those objects that act to maintain a disk around the central object for longer (and perhaps indefinite) periods of time \\citep[][see also \\S~\\ref{sec:theory}]{Robertson06}, and band-shifting that causes us to see those objects in their rest-frame $R$, $V$ or $B$ band where their disks can be readily detected. Another type of bright $24\\mu m$-selected galaxies are the so called DOGs (for dust-obscured galaxies). These are $z \\sim 2$ objects with very red observed $F(24\\mu m) / F(R) > 1000$ colors. \\citet{Bussmann09} presented NICMOS (as well as ACS/WFPC2) morphologies of 31 such sources selected from the larger sample of \\citet{Dey08} and chosen to have $24\\mu m$ fluxes in excess of 0.8 mJy. \\citet{Melbourne09} presented NIR morphologies of 15 additional DOGs obtained through $K$-band AO observations, eleven of which have $24\\mu m$ fluxes above 0.8 mJy. Both papers find DOGs to possess a high fraction of disk-like profiles ($\\sim 90\\%$ and 50\\%, respectively), and argue, based on their size, sersic index, axis ratio and morphological parameters that DOGs are consistent with being late-stage mergers, transitioning from the chaotic coalescence towards becoming relaxed ellipticals. \\begin{deluxetable*} {cccc} \\tablecolumns{4} \\tablewidth{0pt} \\tablecaption{DOG morphologies compared to $24\\mu m$-galaxies\\tablenotemark{a}} \\tablehead{ \\colhead{Merging sequence} & \\colhead{Morphological class} & \\colhead{All $z \\ge 1.5$} & \\colhead{DOGs} \\\\ \\colhead{} & \\colhead{} & \\colhead{$24\\mu m$-galaxies} & \\colhead{} } \\startdata Isolated objects & All spirals & 0 & 0 \\\\ \\\\ First approach & Close pairs (phase I) & 2\\%(1) & 0 \\\\ \\\\ \\multirow{3}{*}{Early mergers} & First contact (phase II) & 2\\%(1) & 0 \\\\ & Pre-mergers (phase III) & 45\\% (27) & 46\\% (12) \\\\ & Triplets & 3\\% (2) & 0 \\\\ \\\\ \\multirow{3}{*}{Coalescence} & Advanced & & \\\\ & Mergers (phase IV) & 26\\% (15.4) & 19\\% (5) \\\\ & Faint \\& Compact & 7\\% (4.4) & 12\\% (3) \\\\ \\\\ \\multirow{3}{*}{Late mergers} & Old mergers (phase V) & 13\\% (8) & 19\\% (5) \\\\ & Regular bulges & 2\\% (1) & 4\\% (1) \\\\ & Pure point sources & 1\\% (0.5) & 0 \\\\ \\enddata \\tablenotetext{a}{After redistribution of objects with unknown redshifts in the proportions found for objects with known redshifts.} \\label{tbl:dogs} \\end{deluxetable*} There are 35 DOGs in our sample (Figure~\\ref{fig:selection}), 75\\% of which are at $z \\ge 1.5$. Not counting $z < 1.5$ objects, we observe that about half of our DOGs are actually pre-mergers (phase III objects), while only half are in more advanced stages, contrary to the hypothesis that they are predominantly late-stage objects. We agree, however, that most of them have disk-dominated profiles, and find that fraction to be 2/3. The morphological properties of DOGs, summarized in Table~\\ref{tbl:dogs}, suggest that they are not very different from the overall $24\\mu m$-bright population at $z \\ge 1.5$. They only tend to draw a little bit more towards AGN-dominated morphological classes, that is faint \\& compact objects and late-stage mergers. As a consequence, we also find them to be slightly more compact, AGN-dominated and obscured than average bright $24\\mu m$-selected galaxies at $z \\ge 1.5$. \\subsection{Comparison with sub-mm selected galaxies} Sub-millimeter selected galaxies (SMGs) form another type of galaxies known to lie at $z \\approx 2$ \\citep{Chapman05} and to be extremely luminous in the infrared \\citep{Kovacs06}. These objects are thus closely related to our bright $24\\mu m$-sources. Unlike our $24\\mu m$-sources, however, SMGs possess PAH-dominated mid-IR spectra in $\\gtrsim 80\\%$ of cases \\citep{Pope08, Menendez09}. They would, therefore, largely populate the high-EW, high-$z$ part of Figure~\\ref{fig:ew_vs_z}: the only region not occupied by our galaxies. The two populations, thus, complement one another in covering the broader high-redshift ULIRG population. Nevertheless, many SMGs show composite spectra, so that there is also significant overlap between the two types of sources (Sajina et al. in preparation). High-resolution HST observations of sub-millimeter galaxies (SMGs) at both optical \\citep{Smail98, Conselice03b, Pope05, Dunlop09, Swinbank10} and near-infrared wavelengths \\citep{Swinbank10} have revealed that a large fraction of them ($40\\% \\textendash 90\\%$) show highly disturbed or multiple component morphologies indicative of merging activity. CO kinematics have similarly confirmed the merging nature of a majority of these objects \\citep{Neri03, Greve05,Tacconi06, Tacconi08, Schinnerer08,Iono09, Bothwell10}, although large star-forming disks have certainly also been found to exist among the sub-millimeter population \\citep{Carilli10, Bothwell10}. Most models also favor a merger origin for SMGs \\citep[e.g.][]{Baugh05, Chakrabarti08, Narayanan09b, Narayanan10}, even though regular disk star formation fueled by cold flow accretion also appears to be a viable path \\citep{Dave10}. For our purposes, we mostly refer to the results of \\citet{Swinbank10}, who published the first morphological analysis of sub-millimeter galaxies at near-infrared wavelengths. Although they did not use the same approach towards morphological classification as we do here, we can still infer, from their $H$-band images, that 40\\% of objects in their sample have two or more distinguishable components. This makes the fraction of early mergers among SMGs very close to that of our $z \\ge 1.5$ galaxies. We can also see that their objects display many of the same kind of disturbed morphologies as ours. We might, thus, expect the morphological distribution of SMGs not to be hugely different from that of our high-$z$ objects, especially given that CO studies, as mentioned above, have found a number of those objects to possess rather chaotic kinematics characteristic of the more advanced merging stages. Nonetheless, we know that purely AGN-dominated objects are not found in sub-millimeter samples, and, based on the results presented in section~\\ref{sec:midIR_vs_morph}, we would therefore expect faint \\& compact objects, as well as very late-stage objects that appear as regular ellipticals, to be absent from SMG samples. This does appear to be true though more detailed analysis would be needed to confirm it. Size and structural analysis, which we postpone for subsequent publication, might also reveal more subtle differences between $24\\mu m$ and sub-millimeter galaxies. \\subsection{Comparison with far-IR selected galaxies} \\citet{Kartaltepe10} have conducted the most comprehensive study of the role of mergers in far-IR selected galaxies to date using the full 2 sq. deg. of HST/ACS imaging and multi-wavelength coverage in the COSMOS field to investigate the morphological properties of a complete sample of 1500 $70\\mu m$-selected galaxies. They too adopted the morphological classification of \\citet{Surace98} based on the merging sequence, enabling direct comparison with our results. This comparison, however, is a little bit complicated by the fact that they use optical rather than NIR imaging, causing $\\sim 20\\%$ of their objects at $z > 1$ and $L_{IR} > 10^{12} L_{\\odot}$ to be too faint to classify morphologically. This caveat is sufficiently minor though to be circumvented, and we demonstrate below that interesting conclusions can still be drawn from a comparison of the two samples. Visually, the $70\\mu m$-galaxies of \\citet{Kartaltepe10} resemble a lot more our low-redshift ($z < 1.5$) sample than either our high-redshift galaxies or the SMGs. Since their sample includes many low-redshift objects with moderate IR-luminosities, we use, for our comparison, only those objects in their sample with a $\\log (L_{IR}/L_{\\odot}) > 11.5$. We also limit our sample to $z < 1.5$ objects in order to ensure that both samples have comparable distributions in both redshift and total IR-luminosity. Table~\\ref{tbl:comparison} shows the detailed comparison by morphological class. In their classification, \\citet{Kartaltepe10} make the distinction between minor and major mergers whereas we do not. Then, many of their sources have unknown morphologies because they are too faint to be detected in the observed optical. For comparison purposes, we, therefore, redistributed their minor mergers, for lack of a better prior, between phases III \\& IV in the same proportion as that of their major mergers, and their unknown objects equally among phase III and \\textquotedblleft faint \\& compact\\textquotedblright objects, since these are the two classes our faint objects fall in. We show in table~\\ref{tbl:comparison} both the observed and redistributed numbers. \\begin{deluxetable*} {ccccc} \\tablecolumns{5} \\tablewidth{0pt} \\tablecaption{Comparison of morphologies of 24 and 70$\\mu m$-selected galaxies} \\tablehead{ \\colhead{Merging sequence} & \\colhead{Morphological class} & \\colhead{$24\\mu m$-selected\\tablenotemark{a}} & \\colhead{$70\\mu m$-selected\\tablenotemark{b}} & \\colhead{$70\\mu m$-selected\\tablenotemark{c}} \\\\ \\colhead{} & \\colhead{} & \\colhead{} & \\colhead{(observed)} & \\colhead{(estimated)} } \\startdata Isolated objects & All spirals & 16\\% & 22\\% & 22\\% \\\\ \\\\ First approach & Close pairs (phase I) & 4\\% & 0 & 0 \\\\ \\\\ \\multirow{3}{*}{Early mergers} & First contact (phase II) & 3\\% & 1\\% & 1\\% \\\\ & Pre-mergers (phase III) & 15\\% & 21\\% & 36\\% \\\\ & and Triplets & & & \\\\ \\\\ \\multirow{3}{*}{Coalescence} & Advanced & & & \\\\ & Mergers (phase IV) & 34\\% & 15\\% & 22\\% \\\\ & Faint \\& Compact & 9\\% & \\nodata & 5\\% \\\\ \\\\ \\multirow{3}{*}{Late mergers} & Old mergers (phase V) & 15\\% & 3\\% & 3\\% \\\\ & Regular bulges & 4\\% & 9\\% & 9\\% \\\\ & Pure point sources & 1\\% & 2\\% & 2\\% \\\\ \\\\ \\nodata & Minor mergers & \\nodata & 16\\% & \\nodata \\\\ \\nodata & Unknowns & \\nodata & 11\\% & \\nodata \\\\ \\enddata \\label{tbl:comparison} \\tablenotetext{a}{$z < 1.5$ objects only} \\tablenotetext{b}{Data from \\citet{Kartaltepe10}, including only objects with $\\log (L_{IR}/L_{\\odot}) > 11.5$} \\tablenotetext{c}{After redistribution of minor mergers and unknowns.} \\end{deluxetable*} Table~\\ref{tbl:comparison} confirms the astonishing correspondence between our $z < 1.5$ sample and their bright $70\\mu m$-sources. It is remarkable that the proportions in all phases of the merging process differ by no more that $\\sim 15\\%$ (of the total sample), lending great support to both our results. Nevertheless, the increase in the number of isolated spirals and early-phase mergers among $70\\mu m$ sources and the corresponding decline in the number of coalescence and post-coalescence mergers when compared to $24\\mu m$ sources is clear, and probably a simple consequence of the higher proportion of star-formation dominated galaxies among $70\\mu m$-selected sources. The broad correspondence between $24\\mu m$ and $70\\mu m$-selected galaxies suggests, as does the similarity between $z \\sim 2$, $24\\mu m$-selected and submillimeter-selected galaxies, that most of the observed morphological specificity in these different samples is common and intrinsic to {\\em all} infrared-luminous objects, largely irrespective of whether they are powered by star formation or an AGN (although these can be responsible for small differences between the various classes). It also suggests that this commonality springs from the origin of those objects in massive mergers. The difference between $z \\sim 1$ and $z \\sim 2$ samples, on the other hand, indicates that there is redshift evolution in the morphological character of IR-luminous objects. We argue in the next section that this difference is the result of higher gas fractions at higher redshifts. \\subsection{Comparison with theory and consequences for galaxy evolution \\label{sec:theory}} Hydrodynamical simulations of merging disk galaxies combined with radiative transfer calculations have provided us with a successful physical model of the origin of ULIRGs and QSOs \\citep{diMatteo05, Narayanan09}. These can be used, more generally, to study the outcome of any merger by varying the set of initial conditions given by parameters such as the mass ratio of the two galaxies, their gas fractions, their relative orientation and so on. Using the statistical results of such a suite of simulations, and embedding them into a cosmological framework using an observationally motivated halo occupation distribution, \\citeauthor{Hopkins08}, in a series of papers starting from (2008), have been able to calculate the role mergers play in the evolution of various global and observable properties of the galaxy population. Most recently, \\citet{Hopkins10a} calculated the IR-luminosity at which, according to their models, star formation transitions from occurring mostly in isolated galaxies to being primarily triggered by mergers. We plotted their results on top of our Figure~\\ref{fig:lum_vs_z} for comparison. Our data agree with their models on that we find nearly all of our isolated spirals below the predicted line, indicating that their models provide an accurate upper limit as to the luminosities at which quiescent star formation usually occurs. We also find good agreement at high redshift where nearly all of our objects fall above their line {\\em and} are, indeed, found to be mergers. We appear to disagree, on the other hand, at lower redshift where we find many merger-identified objects below their dividing line. This, however, is not necessarily a problem, the reason being that the $24\\mu m$ selection of our sample is very different from that of an $L_{IR}$ selection. In particular, we know that our selection highly favors objects whose SED contains an obscured AGN component \\citep{Sajina07,Dasyra09}. If we adopt the results of \\citet{Hopkins10b} which argue that the fueling of a strong AGN (i.e. $\\dot{M}> 0.1 \\mbox{ M}_{\\odot} \\mbox{ yr}^{-1}$ or $L_{IR} \\gtrsim 10^{11} \\mbox{ L}_{\\odot} \\mbox{ yr}^{-1}$) can only occur in mergers, then this means that we are preferentially selecting mergers, and that our sample is biased compared to the general infrared-luminous population. Our data, thus, do not lend themselves very well to a comparison with the results of \\citet{Hopkins10a} except on isolated objects, as discussed above. They do, however, seem to support another result, namely the one argued for in \\citet{Hopkins10b} which is that it is very hard to fuel a strong AGN other than through a merger event. Another aspect of those galaxy merger simulations is that the merging of two massive gas-rich disks almost always produces an extended IR-luminous phase dominated, at first, by star formation during early stages of the merging, and then by accretion onto the central black hole after coalescence. Our results, as well as those from observations of local ULIRGs \\citep{Veilleux09a,Farrah09}, show that the above scenario is, indeed, likely to be the evolutionary path of many low to medium redshift ($z < 1.5$) ULIRGs. However, both our observations and that of local ULIRGs also show that there is a great variety of objects at all stages of the merging process, implying many more possible evolutionary paths. This variety has not been fairly reproduced, so far, in simulations. At high redshifts (and high luminosities), our results indicate that mergers depart even more from that classical scenario, and that black hole accretion plays an important role, at least among our sources, already in the early stages of merging. CO observations have demonstrated that $24\\mu m$-selected galaxies at $z \\sim 2$ possess larger amounts of gas than do local ULIRGs \\citep{Yan10}. We speculate that these elevated gas fractions might be at the origin of the increased AGN activity at early stages. Another characteristic of bright $24\\mu m$-galaxies at $z \\sim 2$ that our observations, as well as that of others \\citep[e.g.][]{Dasyra08, Bussmann09}, have revealed is their preponderantly disk-dominated profiles. Here, we know, thanks to simulations by \\citet{Robertson06} and others, that high gas fractions in major mergers can act to stabilize the disk component and that major mergers can even result in rotationally supported disk remnants provided that the progenitors are extremely gas-rich ($f_{gas} \\ge 0.8$) \\citep[see also][]{Springel05,Barnes02}. The results of these simulations thus provide a natural explanation linking this large observed fraction of disk-like morphologies to the high gas content of these objects. \\citet{Robertson06} argue that the kind of galaxy assembly in which mergers result in a disk remnant is probably limited to high redshifts where gas reservoirs are much larger \\citep{Noterdaeme09}. Although it is impossible to know without detailed kinematics whether any of our $z \\sim 2$ mergers exhibit such behavior, we see, based on their profile, visual appearance and merger stage, two such candidates: MIPS16144 and MIPS16122. Both have clearly detected PAHs (EW$(11.3\\mu m) = 0.5$ and 0.19, respectively). These two tentative disk merger remnants at $z \\sim 2$, and the lack thereof at lower redshifts, combined with the overall large fraction of objects with disk profiles at those redshifts, suggest that this epoch might correspond to the end of the era of gas-rich mergers \\`{a} la \\citet{Robertson06}. Interestingly, that period in cosmos history also concurs with the onset of the rapid rise in mass density of elliptical and red sequence galaxies \\citep{Arnouts07,Ilbert09}. These two observations can naturally explain each other, and point to a transition, at $z \\sim 2$, from an epoch of disk-dominated galaxy evolution to an epoch of bulge formation." }, "1101/1101.2197_arXiv.txt": { "abstract": "We present a contemporary perspective on the String Landscape and the Multiverse of plausible string, M- and F-theory vacua. In contrast to traditional statistical classifications and capitulation to the anthropic principle, we seek only to demonstrate the existence of a non-zero probability for a universe matching our own observed physics within the solution ensemble. We argue for the importance of No-Scale Supergravity as an essential common underpinning for the spontaneous emergence of a cosmologically flat universe from the quantum ``nothingness''. Concretely, we continue to probe the phenomenology of a specific model which is testable at the LHC and Tevatron. Dubbed No-Scale ${\\cal F}$-$SU(5)$, it represents the intersection of the Flipped $SU(5)$ Grand Unified Theory (GUT) with extra TeV-Scale vector-like multiplets derived out of F-theory, and the dynamics of No-Scale Supergravity, which in turn imply a very restricted set of high energy boundary conditions. By secondarily minimizing the minimum of the scalar Higgs potential, we dynamically determine the ratio $\\tan \\beta \\simeq 15-20$ of up- to down-type Higgs vacuum expectation values (VEVs), the universal gaugino boundary mass $M_{1/2} \\simeq 450$~GeV, and consequently also the total magnitude of the GUT-scale Higgs VEVs, while constraining the low energy Standard Model gauge couplings. In particular, this local {\\it minimum minimorum} lies within the previously described ``golden strip'', satisfying all current experimental constraints. We emphasize, however, that the overarching goal is not to establish why our own particular universe possesses any number of specific characteristics, but rather to tease out what generic principles might govern the superset of all possible universes. ", "introduction": "The number of consistent, meta-stable vacua of string, M- or (predominantly) F-theory flux compactifications which exhibit broadly plausible phenomenology, including moduli stabilization and broken supersymmetry~\\cite{Bousso:2000xa, Giddings:2001yu, Kachru:2003aw, Susskind:2003kw, Denef:2004ze, Denef:2004cf}, is popularly estimated~\\cite{Denef:2004dm,Denef:2007pq} to be of order $10^{500}$. It is moreover currently in vogue to suggest that degeneracy of common features across these many ``universes'' might statistically isolate the physically realistic universe from the vast ``landscape'', much as the entropy function coaxes the singular order of macroscopic thermodynamics from the chaotic duplicity of the entangled quantum microstate. We argue here though the counter point that we are not obliged {\\it a priori} to live in the likeliest of all universes, but only in one which is possible. The existence merely of a non-zero probability for our existence is sufficient. We indulge for this effort the fanciful imagination that the ``Multiverse'' of string vacua might exhibit some literal realization beyond our own physical sphere. A single electron may be said to wander all histories through interfering apertures, though its arrival is ultimately registered at a localized point on the target. The journey to that destination is steered by the full dynamics of the theory, although the isolated spontaneous solution reflects only faintly the richness of the solution ensemble. Whether the Multiverse be reverie or reality, the conceptual superset of our own physics which it embodies must certainly represent the interference of all navigable universal histories. Surely many times afore has mankind's notion of the heavens expanded - the Earth dispatched from its central pedestal in our solar system and the Sun rendered one among some hundred billion stars of the Milky Way, itself reduced to one among some hundred billion galaxies. Finally perhaps, we come to the completion of our Odyssey, by realizing that our Universe is one of at least $10^{500}$ so possible, thus rendering the anthropic view of our position in the Universe (environmental coincidences explained away by the availability of $10^{11} \\times 10^{11}$ solar systems) functionally equivalent to the anthropic view of the origin of the Universe (coincidences in the form and content of physical laws explained away by the availability, through dynamical phase transitions, of $10^{500}$ universes). Nature's bounty has anyway invariably trumped our wildest anticipations, and though frugal and equanimous in law, she has spared no extravagance or whimsy in its manifestation. Our perspective should not be misconstrued, however, as complacent retreat into the tautology of the weak anthropic principle. It is indeed unassailable truism that an observed universe must afford and sustain the life of the observer, including requisite constraints, for example, on the cosmological constant~\\cite{Weinberg:1987dv} and gauge hierarchy. Our point of view, though, is sharply different; we should be able to resolve the cosmological constant and gauge hierarchy problems through investigation of the fundamental laws of our (or any single) Universe, its accidental and specific properties notwithstanding, without resorting to the existence of observers. In our view, the observer is the output of, not the {\\it raison d'\\^etre} of, our Universe. Thus, our attention is advance from this base camp of our own physics, as unlikely an appointment as it may be, to the summit goal of the master theory and symmetries which govern all possible universes. In so seeking, our first halting forage must be that of a concrete string model which can describe Nature locally. ", "conclusions": "The advancement of human scientific knowledge and technology is replete with instances of science fiction transitioning to scientific theory and eventually scientific fact. The conceptual notion of a ``Multiverse'' has long fascinated the human imagination, though this speculation has been largely devoid of a substantive underpinning in physical theory. The modern perspective presented here offers a tangible foundation upon which legitimate discussion and theoretical advancement of the Multiverse may commence, including the prescription of specific experimental tests which could either falsify or enhance the viability of our proposal. Our perspective diverges from the common appeals to statistics and the anthropic principle, suggesting instead that we may seek to establish the character of the master theory, of which our Universe is an isolated vacuum condensation, based on specific observed properties of our own physics which might be reasonably inferred to represent invariant common characteristics of all possible universes. We have focused on the discovery of a model universe consonant with our observable phenomenology, presenting it as confirmation of a non-zero probability of our own Universe transpiring within the larger String Landscape. The archetype model universe which we advance in this work implicates No-Scale Supergravity as the ubiquitous supporting structure which pervades the vacua of the Multiverse, being the crucial ingredient in the emanation of a cosmologically flat universe from the quantum ``nothingness''. In particular, the model dubbed No-Scale ${\\cal F}$-$SU(5)$ has demonstrated remarkable consistency between parameters determined dynamically (the top-down approach) and parameters determined through the application of current experimental constraints (the bottom-up approach). This enticing convergence of theory with experiment elevates No-Scale ${\\cal F}$-$SU(5)$, in our estimation, to a position as the current leading GUT candidate. The longer term viability of this suggestion is likely to be greatly clarified in the next few years, based upon the wealth of forthcoming experimental data. Building on the results presented in prior works~\\cite{Li:2010ws,Li:2010mi,Li:2010uu}, we have presented a dynamic determination of the penultimate Flipped $SU(5)$ unification scale $M_{32}$, or more fundamentally, the GUT Higgs VEV moduli. We have demonstrated that the $B_{\\mu}$ = 0 No-Scale boundary condition is again vital in dynamically determining the model parameters. Procedurally, we have fixed the unified gauge coupling, SM fermion Yukawa couplings, and Higgs bilinear term $\\mu\\simeq 460~{\\rm GeV}$ at the final unification scale $M_{\\cal F}$, while concurrently allowing the VEVs of the GUT Higgs fields $H$ and $\\overline{H}$ to float freely, as driven by $M_{32}$ and the low energy SM gauge couplings, via variation of the Weinberg angle. Employing the ``Super No-Scale'' condition to secondarily minimize the effective Higgs potential, we have obtained $M_{32}\\simeq 1.0 \\times 10^{16}$~GeV, $\\sin^2 (\\theta_{\\rm W}) \\simeq 0.236$, and $\\tan \\beta \\simeq 15-20$. The blueprints which we have outlined here, integrating precision phenomenology with prevailing experimental data and a fresh interpretation of the Multiverse and the Landscape of String vacua, offer a logically connected point of view from which additional investigation may be mounted. As we anticipate the impending stream of new experimental data which is likely to be revealed in ensuing years, we look forward to serious discussion and investigation of the perspective presented in this work. Though the mind boggles to contemplate the implications of this speculation, so it must also reel at even the undisputed realities of the Universe, these acknowledged facts alone being manifestly sufficient to humble our provincial notions of longevity, extent, and largess." }, "1101/1101.3223.txt": { "abstract": "Starting from the Cambridge Catalogues of radio sources, we have created a sample of 401 FRII radio sources that have counterparts in the main galaxy sample of the 7th Data release of the Sloan Digital Sky Survey and analyse their radio and optical properties. We find that the luminosity in the H$\\alpha$ line -- which we argue gives a better measure of the total emission-line flux than the widely used luminosity in \\oiii\\ -- is strongly correlated with the radio luminosity $P_{\\rm 1.4GHz}$. We show that the absence of emission lines in about one third of our sample is likely due to a detection threshold and not to a lack of optical activity. We also find a very strong correlation between the values of $L_{\\Ha}$ and $P_{\\rm 1.4GHz}$ when scaled by $``M_{\\rm BH}\"$, an estimate of the black hole mass. We find that the properties of FRII galaxies are mainly driven by the Eddington parameter $L_{\\Ha}$/$``M_{\\rm BH}\"$ or, equivalently, $P_{\\rm 1.4GHz}$/$``M_{\\rm BH}\"$. Radio galaxies with hot spots are found among the ones with the highest values of $P_{\\rm 1.4GHz}$/$``M_{\\rm BH}\"$. %We do not see any dichotomy between high and low excitation radio galaxies but rather a continuum of properties driven by the Eddington parameter $L_{\\Ha}$/$``M_{\\rm BH}\"$ or, equivalently, $P_{\\rm 1.4GHz}$/$``M_{\\rm BH}\"$. We note, however, that radio galaxies with hot spots are found among the ones with the highest values of $P_{\\rm 1.4GHz}$/$``M_{\\rm BH}\"$. Compared to classical AGN hosts in the main galaxy sample of the SDSS, our FRII galaxies show a larger proportion of objects with very hard ionizing radiation field and large ionization parameter. A few objects are, on the contrary, ionized by a softer radiation field. Two of them have double-peaked emission lines and deserve more attention. We find that the black hole masses and stellar masses in FRII galaxies are very closely related: $``M_{\\rm BH}\"$ $\\propto$ $M_{*}^{1.13}$ with very little scatter. A comparison sample of line-less galaxies in the SDSS follows exactly the same relation, although the masses are, on average, smaller. This suggests that the FRII radio phenomenon occurs in normal elliptical galaxies, preferentially in the most massive ones. Although most FRII galaxies are old, some contain traces of young stellar populations. Such young populations are not seen in normal line-less galaxies, suggesting that the radio (and optical) activity in some FRII galaxies may be triggered by recent star formation. The $``M_{\\rm BH}\"$ -- $M_{*}$ relation in a comparison sample of radio-quiet AGN hosts from the SDSS is very different, suggesting that galaxies which are still forming stars are also still building their central black holes. Globally, our study indicates that, while radio and optical activity are strongly related in FRII galaxies, the features of the optical activity in FRIIs are distinct from those of the bulk of radio-quiet active galaxies. An appendix gives the radio maps of our FRII galaxies, superimposed on the SDSS images, and the parameters derived for our analysis that were not publicly available. ", "introduction": "\\label{sec:Introduction} According to the Collins English Dictionary, the term \"radio galaxy\" (RG) refers to a galaxy that is a strong emitter of radio waves. However, only sources powered by accretion onto a super massive black hole can produce the extended structures called radio jets and lobes, and we will use the term radio galaxy in reference only to those sources. Accretion unto a massive black hole is what is considered to be the energy source of active galactic nuclei (AGN). It manifests itself in different ways and with different strengths in the entire electromagnetic spectrum. Most active galactic nuclei in the Sloan Digital Sky Survey (SDSS, York et al. 2000) are not radio active. Conversely, not all classical radio galaxies with extended radio lobes have emission lines in their optical spectra (Hine \\& Longair, 1979), suggesting that optical and radio activity are not necessarily concommittent (Best et al. 2005b). Radio galaxies can be divided into two classes according to the morphology of their radio structure (Fanaroff \\& Riley, 1974). FRI radio galaxies are core-dominated sources with radio jets fading and dissipating on a short distance from center, while FRII radio galaxies are edge-brightened sources with highly collimated jets. The sizes of both types of radio galaxies range from a few kiloparsecs for compact steep spectrum sources to a few megaparsecs for giant radio galaxies. FRII radio galaxies tend to be more luminous than FRI ones. Fanaroff \\& Riley suggested a dividing luminosity of $L_{178MHz} \\sim$ 2.5 $\\times$ 10$^{26}$W Hz$^{-1}$, but as shown by Ledlow \\& Owen (1996) the dividing luminosity between FRI and FRII radio sources is a function of the host galaxy optical luminosity. With the availability of observational data from large radio and optical surveys came the possibility to examine the optical properties of large samples of radio galaxies. Best et al. (2005a) cross-identified radio sources from the NVSS (Condon et al. 1998) and FIRST (Becker et al. 1995) radio surveys with the main galaxy sample of the second data release of the SDSS (Abazajian et al. 2004). The resulting sample, which is not morphology-specific is likely dominated by FRI or compact radio-sources, as can be judged by the low radio luminosities of the vast majority of their sample. Using this (or a similar) sample Kauffmann, Heckman \\& Best (2008) have argued that radio emission is likely due to the accretion of \\textit{hot} gas by massive black holes central to galaxies in a dense environment, while the optical AGN phenomenon, strongly favoured by the presence of a young stellar population, is likely due to the accretion of\\textit{ cold} gas. However, their sample lacks the brightest radio sources, the FRII ones, both because they are rare in the local Universe and because the very extended angular sizes of these sources does not allow easy cross-identification with optical galaxies. FRII radio galaxies constitute a much better defined class than FRI radio galaxies in terms of radio morphology (Fanaroff \\& Riley, 1974). Besides, because of the diversities of their optical properties, they constitute a prime target for understanding the relation between optical and radio activity. We have assembled a sample of morphologically selected FRII galaxies with optical counterparts available in the SDSS. Here, we present the result of our study comparing the radio and optical properties of FRII radio galaxies, including giant radio galaxies, in order to get more insight into the activity phenomenon in galaxies. Our data sample is, by necessity, much smaller than the sample used by Kauffmann et al. (2008), since FRII galaxies are much less common than FRI galaxies at low redshifts. Due to the selection process (see next section), our sample does not allow us to tackle such issues as luminosity distribution functions, which are examined by Kauffmann \\& Heckman (2009) using the same sample as Kauffmann et al. (2008). But we can look for the presence of correlations that might improve our understanding of radio loud AGN. The organization of the paper is as follows: In Section 2 we provide a brief description of the sample selection and data processing. In Section 3 we analyze the relation between the radio power of FRII sources and the strength of their optical activity. In Section 4 we focus on the special class of FRII galaxies that show hot spots. In Section 5 we discuss the emission line properties of FRII galaxies. In Section 6 we put FRII galaxies in the context of other groups of galaxies: line-less galaxies and radio-quiet AGN hosts. The main results of our investigation are summarized in Section 7. Throughout this paper we assume a $\\Lambda$ Cold Dark Matter cosmology with $H_{0}=71$km\\,s$^{-1}$Mpc$^{-1}$, $\\Omega_{\\rm m}=0.27$, and $\\Omega _{\\Lambda}=0.73$ (Spergel et al. 2003). %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "Using the Cambridge Catalogues of radio sources we have built a sample of FRII radio galaxies whose spectra are available in the main galaxy sample of SDSS DR7. From this sample, after inspection of the NVSS and FIRST radio maps, we extracted a sample of 401 FRII radio galaxies. In this paper, we examined the optical and radio properties of those objects, in order to find new clues about the relation between radio activity and the optical manifestations of the AGN in those objects. The stellar masses, emission line equivalent widths and fluxes were taken from the \\starlight\\ database (Cid Fernandes et al. 2009). We found that the luminosity in the H$\\alpha$ line -- which we argue gives a better measure of the total flux in the emission lines than the widely used luminosity in \\oiii\\ -- is strongly correlated with the radio luminosity $P_{\\rm 1.4GHz}$ over more than three orders of magnitude in $P_{\\rm 1.4GHz}$. A similar result was found by Zirbel \\& Baum (1995), using \\Ha\\ + \\nii, but they obtained $L_{\\Ha+\\nii}$ $\\propto$ $P_{\\rm 1.4GHz}^{0.75}$ while we obtain $L_{\\Ha}$ $\\propto$ $P_{\\rm 1.4GHz}^{1.13}$. In our sample, there is about one third of objects which do not have any line detected. We showed that, for those, the detection threshold is above the empirical relations between $L_{\\oiii}$ or $L_{\\Ha}$ and $P_{\\rm 1.4GHz}$. Therefore, there is nothing for the moment that indicates the existence of two classes of FRII radio galaxies with respect to the presence of emission lines. We also find a very strong correlation between the values of $L_{\\Ha}$ and $P_{\\rm 1.4GHz}$ when scaled by $``M_{\\rm BH}\"$, the black hole masses obtained from the observed stellar velocity dispersion suggesting that, in FRII radio galaxies, optical and radio activity have a common cause. Contrary to previous work (e.g. Buttiglione et al 2010, or Lin et al. 2010, however based on different samples) we see no sharp transition between high- and low-excitation radio galaxies or galaxies with $L_{\\oiii}$ smaller or larger than $10^6 L_\\odot$. We rather find that FRII galaxies present a continuum of properties driven by the Eddington parameter $L_{\\Ha}$/$``M_{\\rm BH}\"$ or, equivalently, $P_{\\rm 1.4GHz}$/$``M_{\\rm BH}\"$, %Contrary to previous studies (e.g. Buttiglione et al 2010, or Lin et al. 2010) we do not detect any dichotomy between high excitation and low excitation radio galaxies or galaxies with $L_{\\oiii}$ $> 10^6 L_\\odot$ but rather a continuum of properties where the main factor is the Eddington parameter $L_{\\Ha}$/$``M_{\\rm BH}\"$ or, equivalently, $P_{\\rm 1.4GHz}$/$``M_{\\rm BH}\"$. We note, however, that FRII galaxies with hot spots are found among the ones with the highest values of $P_{\\rm 1.4GHz}$/$``M_{\\rm BH}\"$. Those hot spots, which are the places where the relativistic jets from the active galactic nuclei interact with the environment, thus seem to require a radiatively efficient accretion to be produced. Those FRII galaxies that can be plotted in the classical BPT diagram or similar diagrams fall in the zone characterized by a hard ionizing spectrum, where the contribution of hot stars to the excitation is small or absent. This is expected, since it is known that radio galaxies are in general associated with massive, elliptical galaxies. Compared to classical AGN hosts found in the galaxy sample of the SDSS, there is a significantly larger proportion of objects with very hard ionizing radiation field and large ionization parameter. There are, however a few objects that lie close to the divisory line between pure star-forming galaxies and AGN hosts. This is a priori surprising, as there is no indication of present-day star formation in those objects. Two of them have double-peaked lines, as are found in some AGNs and attributed to binary black holes (among other possibilities). We suggest that those objects are ionized by a rather soft radiation field, as compared with the rest of the FRIIs. For the 346 FRII galaxies for which we could determine both the black hole mass and the stellar mass, we find that $``M_{\\rm BH}\"$ varies like $M_{*}^{1.13}$ with very little scatter. A comparison sample of line-less galaxies in the SDSS follows exactly the same relation, but with both masses shifted to lower values. This suggests that the FRII radio phenomenon occurs in normal elliptical galaxies, but is favoured by larger galaxy masses. The $D_{\\rm n}(4000)$ index indicates that, although most of the FRII galaxies are old, some contain traces of young stellar populations. Since such young populations are not seen in normal line-less galaxies, one can conjecture that the radio (and optical) activity is triggered by recent star formation. The $``M_{\\rm BH}\"$ -- $M_{*}$ relation in a comparison sample of radio-quiet AGN hosts from the SDSS is very different, suggesting that galaxies which are presently forming stars are still building their central black hole. The $D_{\\rm n}(4000)$ index in this sample indicates that the youngest galaxies have the smallers $``M_{\\rm BH}\"$/$M_{*}$ ratio, confirming this view. Overall, our study leads to the conclusion that, while radio and optical activity are strongly related in FRII galaxies, the features of the optical activity in those objects are distinct from those of radio quiet active galaxies." }, "1101/1101.6061_arXiv.txt": { "abstract": "We discuss constraints for the equation of state of hybrid star matter which can be obtained from next generation heavy-ion collisions at FAIR and NICA. Particular emphasis is on the planned NICA facility at JINR Dubna which shall provide fixed-target and collider experiments just in the relevant energy ranges. ", "introduction": " ", "conclusions": "" }, "1101/1101.4778_arXiv.txt": { "abstract": "We present the discovery of 5 millisecond pulsars found in the mid-Galactic latitude portion of the High Time Resolution Universe (HTRU) Survey. The pulsars have rotational periods from $\\sim2.3\\,\\rm{ms}$ to $\\sim7.5\\,\\rm{ms}$, and all are in binary systems with orbital periods ranging from $\\sim0.3$ to $\\sim150$~d. In four of these systems, the most likely companion is a white dwarf, with minimum masses of $\\sim 0.2\\,\\mathrm{M}_\\odot$. The other pulsar, J1731$-$1847, has a very low mass companion and exhibits eclipses, and is thus a member of the ``black widow'' class of pulsar binaries. These eclipses have been observed in bands centred near frequencies of 700, 1400 and 3000~MHz, from which measurements have been made of the electron density in the eclipse region. These measurements have been used to examine some possible eclipse mechanisms. The eclipse and other properties of this source are used to perform a comparison with the other known eclipsing and ``black widow'' pulsars. These new discoveries occupy a short-period and high-dispersion measure (DM) region of parameter space, which we demonstrate is a direct consequence of the high time and frequency resolution of the HTRU survey. The large implied distances to our new discoveries makes observation of their companions unlikely with both current optical telescopes and the Fermi Gamma-ray Space Telescope. The extremely circular orbits make any advance of periastron measurements highly unlikely. No relativistic Shapiro delays are obvious in any of the systems, although the low flux densities would make their detection difficult unless the orbits were fortuitously edge-on. ", "introduction": "Millisecond pulsars (MSPs) are neutron stars (NSs) with rapid rotation rates that are believed to be formed in binary systems when the NS accretes matter from the companion, causing a `spin-up' in the NS's rotation rate \\citep[e.g.\\,][]{alpar1982}. This process results in pulsars with spin periods between 1 and 100~ms and magnetic field strengths less than $10^{10}$~G. Those systems with spin periods less than about 20~ms are thought to have had low mass companions, while those with periods between 20 and 100~ms are thought to have been spun-up by a heavy white dwarf (WD) or NS companion. During the spin-up phase, the heating caused by accretion leads to the emission of X-ray radiation \\citep{do1973}. These accreting systems are known as high-mass and low-mass X-ray binaries depending upon the companion mass (HMXBs and LMXBs respectively, see \\citet{bvdh1991} for details of evolution). The link between LMXBs and MSPs was not confirmed until pulsations with a period of 2.4~ms were observed in the accreting X-ray binary, SAX~J1808.4$-$3658 \\citep{wvdk1998}. This link has more recently been reinforced by the discovery of PSR~J1023$+$0038 \\citep{asr2009}, where radio emission is likely to have only switched on recently after an LMXB phase. In contrast to this evolutionary model, however, around 20\\% of the MSP population appear to be isolated bodies, including the first MSP to be discovered, PSR~B1937+21 \\citep{backer1982}. Therefore, if this scenario is correct, and the isolated MSPs are the descendants of the eclipsing MSPs, we might expect some of them to eventually ablate their companions until nothing remains \\citep{rst1989}. The discovery of the so-called `black widow' pulsar, PSR~B1957$+$20 \\citep{fst1988} and other similar systems \\citep[e.g.\\,PSR~J2051$-$0827,\\,][]{bws1996}, provides some evidence of this process taking place. In these systems, the pulsar's companion is typically of very low mass ($\\sim 0.02\\,\\mathrm{M}_\\odot$), and the pulsar is eclipsed by the companion for at least 10\\% of the orbit. As the pulsar approaches eclipse, the pulses experience a delay due to the additional ionised gas through which the radiation has to pass, while outflowing material from the companion often gives rise to anomalous eclipses away from superior conjunction. However, the timescale over which the ablation process would lead to the companion being destroyed is far too long to explain the abundance of isolated MSPs \\citep{bws1996}. Clearly other mechanisms such as tidal disruption are needed to explain the existence of this group. The discovery of PSR~J1903$+$0327 \\citep{champion2008} further challenges the conventional formation scenarios. Not only is the pulsar's companion a main-sequence star, but the orbit is also highly eccentric ($e=0.44$), neither of which are predicted by the standard formation scenarios. Three alternative theories for the formation have been proposed; \\begin{inparaenum}[\\itshape a\\upshape)] \\item the pulsar was born in an eccentric orbit, spinning rapidly; \\item the pulsar was spun-up in a globular cluster before being ejected into the Galactic disk; and \\item the recycling of the pulsar in a triple system \\citep{champion2008}. \\end{inparaenum} \\citet{freire2010} have since confirmed that the main sequence star is the companion, and excluded the possibility that PSR~J1903+0327 is currently a member of a triple system, but maintain that the pulsar was born in a triple system, from which the component originally responsible for the spun-up of PSR~J1903+0327 has been ejected. However, if the alternatives are viable, we need to understand how many MSPs form via this channel and to do that a larger sample of MSPs is required. Beyond improving our understanding of binary star evolution, there are many areas of pulsar science which benefit from the discovery of recycled pulsars. The discovery of binaries with, for example, a neutron star or black hole companion offer the chance to test General Relativity (GR) and other theories of gravity in the strong field regime \\citep[e.g.\\,][]{kramer2004,kramer2006}. The case of a pulsar with a WD companion also offers the possibility of observing effects predicted by GR. For example, PSR~J1909$-$3744 is in a 1.5~day orbit with a companion of mass $0.2\\,\\mathrm{M}_\\odot$, where Shapiro delay has been observed, allowing the measurement of both the companion and pulsar masses \\citep{jacoby2003, jacoby2005}. The NS equation of state is currently poorly understood, and the predictions of NS masses and radii vary dramatically for different models \\citep{lp2004}. One way to place a limit on these parameters is to identify the limiting rotation frequency, beyond which the neutron star would break apart. Currently, the most rapidly-rotating known pulsar is PSR J1748$-$2446ad \\citep{hrs+06}, which rotates at 716~Hz; the discovery of an MSP rotating even faster than this could rule out some NS equation of state models \\citep{lp2007, lp2010}. By constraining the maximum mass of a NS, it should also be possible to eliminate many models of the equation of state \\citep{lp2004}. To do this, one can use binary effects such as Shapiro delay which allow the pulsar mass to be measured \\citep[e.g.\\ ][]{lyne2004}; an effect which is most measurable for those systems where the pulsar is found to have a massive companion or a near edge-on orbit \\citep[e.g.\\,][]{demorest2010}. Therefore, the discovery of more systems where the NS mass can be measured \\citep[e.g.\\,][]{freire2010} would contribute to the understanding of the NS equation of state. For some pulsars in binary systems, it has been possible to observe the companion optically \\citep[e.g.\\,][]{bassa2006} and measure its radial velocity as a function of orbital phase. However, for WD companions, often the faintness of the companion only allows the measurement of a temperature, from which a mass must be estimated using evolutionary models \\citep{vankerkwijk1996}. Currently, pulsar timing arrays \\citep{hobbs2009,ferdman2010,jenet2009} are attempting to detect gravitational waves by the correlation of arrival times \\citep[e.g.\\,][]{jenet2006} from pulsars. These arrays require long-term, high precision timing observations with high signal-to-noise ratio; MSPs, with their short spin periods and stable rotation rates, are well-suited for these timing arrays \\citep{verbiest2009}. However, not all MSPs have the narrow pulse profiles, regular rotation and high flux density required to obtain high timing precision, nor is the distribution of the known sources as uniform over the sky as desired \\citep[see, e.g.\\,][]{hd1983}. Hence, further discoveries from pulsar surveys may make important contributions to pulsar timing arrays and improve their sensitivity to the stochastic background of gravitational waves. The largest previous pulsar survey, that found many of the known MSPs, was the Parkes Multibeam pulsar survey (PMPS) by \\citet{mlc+01}, which surveyed the area of the Galactic plane bounded by $260\\degree \\leq l \\leq 50\\degree, |b| \\leq 5\\degree$, discovering 26 MSPs. At higher Galactic latitudes, the Swinburne Intermediate-latitude pulsar survey ($260\\degree \\leq l \\leq 50\\degree, 5\\degree \\leq |b| \\leq 15\\degree$), using a similar observing system except for a reduced sampling time, discovered a further 8 MSPs \\citep{ebsb01,eb2001}. However, these surveys were limited by the relatively broad frequency channels (3~MHz) and coarse time sampling (250~$\\mu$s and 125~$\\mu$s, respectively) compared with today's digital backends, which provide increased sensitivity to distant and rapidly-rotating pulsars. The ionised interstellar medium disperses the pulses emitted by radio pulsars as they traverse it. When removing this dispersion, increased frequency resolution allows this correction to be made with reduced smearing of the pulses. The increased time resolution provides greater sensitivity to short-period pulses, such as those from MSPs, and by sampling with 2 bits (compared to 1 bit per sample in the PMPS) sensitivity is further increased by reducing losses due to digitisation. The High Time Resolution Universe Survey \\citep[HTRU, ][]{keith2010} aims to make use of these improvements in backend technology to perform a survey of the entire southern sky with much-improved sensitivity to MSPs. The ongoing PALFA survey at Arecibo \\citep{cordes2006} and the GBT 350-MHz survey \\citep{boyles2010} have each discovered several MSPs which occupy the short-period and high-DM region of parameter space which we look to probe with HTRU. In particular, the discovery, as previously mentioned, of PSR~J1903$+$0327 \\citep{champion2008} as part of PALFA stands out with a rotation period of 2.15~ms and a DM of $297.5\\,\\mathrm{cm}^{-3}\\,\\mathrm{pc}$ as evidence that time resolution of 64~$\\mu$s allows this parameter space to be probed. In this paper we outline the spin and orbital parameters of five MSPs discovered as part of the HTRU survey, compare the properties of these pulsars to the previously-known population, and study in detail the eclipses that one of them displays. These discoveries have been made with $\\sim30\\%$ of the survey region observed, indicating that we might expect tens of MSPs to be discovered when the survey is completed. \\begin{table} \\begin{center} \\caption{Observational parameters for the mid-latitude portion of the HTRU survey.} \\begin{tabular}{lr} \\toprule Number of beams & 13 \\\\ Polarizations/beam & 2 \\\\ Centre Frequency & 1352 MHz\\\\ Frequency channels & 1024 $\\times$ 390.625 kHz* \\\\ \\midrule Galactic longitude range & $-120\\degree$ to $30\\degree$ \\\\ Galactic latitude range & $|b| \\leq 15\\degree$ \\\\ Sampling interval & 64 $\\mathrm{\\mu}$s \\\\ Bits/sample & 2 \\\\ Observation time/pointing & 540 s \\\\ \\bottomrule \\end{tabular} \\label{table:survey} \\end{center} *154 of these channels are then masked to remove interference \\end{table} ", "conclusions": "We present the discovery of 5 MSPs in the HTRU survey. These MSPs have short periods and DMs among the highest in the known population. One of the pulsars, PSR~J1731$-$1847, displays regular eclipses in its 0.3~d orbit and with a companion of minimum mass $0.04\\,\\mathrm{M}_\\odot$, appears to be a member of the `black widow' group of MSPs. The sensitivity to such objects can be attributed to the fast sampling rate and narrow filterbank channels of the hardware used for the HTRU survey." }, "1101/1101.4314_arXiv.txt": { "abstract": "Evidence is building that remnants of solar systems might orbit a large percentage of white dwarfs, as the polluted atmospheres of DAZ and DBZ white dwarfs indicate the very recent accretion of metal-rich material. \\citep{2010ApJ...722..725Z}. Some of these polluted white dwarfs are found to have large mid-infrared excesses from close-in debris disks that are thought to be reservoirs for the metal accretion. These systems are coined DAZd white dwarfs \\citep{2007ApJ...662..544V} Here we investigate the claims of \\citet{2008A&A...489..651B} that Sirius B, the nearest white dwarf to the Sun, might have an infrared excess from a dusty debris disk. Sirius B's companion, Sirius A is commonly observed as a mid-infrared photometric standard in the Southern hemisphere. We combine several years of Gemini/T-ReCS photometric standard observations to produce deep mid-infrared imaging in five $\\sim$10$\\micron$ filters (broad N + 4 narrowband), which reveal the presence of Sirius B. Our photometry is consistent with the expected photospheric emission such that we constrain any mid-infrared excess to $\\lesssim$10\\% of the photosphere. Thus we conclude that Sirius B does not have a large dusty disk, as seen in DAZd white dwarfs. ", "introduction": "Since its discovery in 1844 \\citep{1844MNRAS...6R.136B}, Sirius B has been a tantalizing object. While its close proximity to the Sun makes it ideal for detailed study, its binary companion, Sirius A, is the brightest star in the night-sky, complicating observations of Sirius B. With a $\\Delta$ mag of $\\sim$10 between Sirius A and B for optical and longer wavelengths, the most successful observations of Sirius B have been space-based. \\textit{Hubble Space Telescope (HST/STIS)} spectra of Sirius B have led to the accurate determination of its effective temperature ($T_{\\rm eff}=$25,193 K) and mass \\citep[0.978 $M_{\\sun}$;][]{2005MNRAS.362.1134B}, \\textit{Hipparcos} parallax measurements of Sirius A place the system 2.64 pc from the Sun \\citep{1997A&A...323L..49P}, and astrometric monitoring has determined Sirius B's orbital period (50.090 yrs), semi-major axis (7.500\" or 19.8 AU), and eccentricity \\citep[0.5923;][]{2001AJ....122.3472H}\\footnote{http://ad.usno.navy.mil/wds/orb6.html}. Increasingly, ground-based adaptive optics (AO) systems are becoming adept at high-contrast imaging, usually with the goal of discovering faint planets/brown dwarfs near their bright host stars. Imaging a faint white dwarf like Sirius B around a bright, main-sequence A star like Sirius A is a similar problem. Recently, \\citet{2008A&A...489..651B} used the ESO 3.6 meter telescope, along with the ADONIS AO system, to image Sirius B in the JHKs filters, which, before this work, were the longest wavelength photometric measurements of Sirius B. The \\citet{2008A&A...489..651B} measurements showed a small (1.7$\\sigma$) Ks-band excess when compared to the models of \\citet{2006AJ....132.1221H}. Similar near-IR excesses are present in the spectra of dusty white dwarfs, which are characterized by large excesses in the mid-infrared \\citep{2007ApJ...662..544V}. The first infrared excess around a white dwarf (G29-38) was found by \\citet{1987Natur.330..138Z}, and was initially thought to be a brown dwarf. This hypothesis was ruled out by optical/near-IR pulse monitoring that suggested a disk-geometry source of the excess \\citep{1990ApJ...357..216G,1991ApJ...374..330P}. Subsequent spectroscopy using \\textit{Spitzer}/IRS revealed the presence of small dust grains \\citep{2005ApJ...635L.161R}. A \\textit{Spitzer}/IRS survey found 4 dusty white dwarfs out of their sample of 124 \\citep{2007ApJS..171..206M,2007ApJ...662..544V}, and that each of the dusty white dwarfs were of type DAZ (metal-polluted, hydrogen atmosphere), leading \\citet{2007ApJ...662..544V} to coin DAZ stars with mid-infrared excesses, DAZd. The metals in DAZ white dwarfs are expected to settle below the white dwarf photosphere much faster than evolutionary timescales \\citep{2006A&A...453.1051K}, thus the presence of metals implies a recent accretion event. $\\sim$1/4 of DA (hydrogen atmosphere) white dwarfs are DAZ \\citep{1998ApJ...505L.143Z}, and $\\sim$1/3 of DB (helium atmosphere) white dwarfs are DBZ \\citep{2010ApJ...722..725Z}, demonstrating that accretion must be a common phenomenon for white dwarfs. The cause of this accretion, in many if not all cases, is thought to be tidally disrupted asteroid-sized objects from a remanent debris disk/planetary system \\citep{2008AJ....135.1785J,2010ApJ...722..725Z}. We note that Sirius B is a DA white dwarf, but determining if it is metal polluted (DAZ) is complicated by the difficulty of taking high-resolution spectra of Sirius B from the ground. This problem is exacerbated by the fact that Sirius B's high temperature, \\citep[25,193 K][]{2005MNRAS.362.1134B} impedes the detection and interpretation of photospheric metals \\citep{2006A&A...453.1051K,1995ApJ...454..429C}. Sirius B's high temperature also has implications for its potential to host a debris disk. While typical DAZd white dwarfs have temperatures ranging from T$\\approx$7,000-15,000 K \\citep{2009ApJ...694..805F}, Sirius B's high temperature would sublimate dust out to its tidal truncation radius \\citep{2003ApJ...584L..91J}, precluding tidal disruption of asteroids as the source of dust in the system. Instead of dust disks, hot white dwarfs can have metal vapor disks. The first system observed to display these features, SDSS 1228+1040 \\citep{2006Sci...314.1908G}, was subsequently found to have a dust disk at larger radii \\citep{2009ApJ...696.1402B}. Analogously, a debris disk around Sirius B would be at larger radii than is typical for DAZd white dwarfs, and the dust would be the result of collisions, rather than tidal disruption of asteroids. Since not all DAZ and DBZ white dwarfs are found to have debris disks, the source of their photospheric pollution might come from larger radii than are easily probed by current mid-infrared instruments (24$\\micron$ probes $\\sim$120 K dust, which is at $\\lesssim$1 AU, assuming radiative equilibrium). As a result, little is known about the outer regions of white dwarf debris disks. Spatially resolving a debris disk around a white dwarf would obviously be an important step in understanding the DAZd phenomenon. As the nearest white dwarf to the Sun, Sirius B would be a prime target for such a search, especially given the claims of \\citet{2008A&A...489..651B} that Sirius B might have a substantial mid-infrared excess. The importance of such a discovery would be magnified by the fact that dust particles $\\gtrsim$1 AU from Sirius B would be primarily heated by Sirius A (18 AU in 2005.0), which would potentially make it possible to probe the outer parts of the debris disk at much greater separations than is normally possible. The maximum radius of Sirius B's disk, based on tidal truncation from Sirius A, would be $\\sim1/4$ of the orbital semi-major axis \\citep{1994ApJ...421..651A}, which is 5 AU or 2\". In this work, we use Gemini/T-ReCS \\citep{1998SPIE.3354..534T} archival observations of Sirius to determine if Sirius B has a strong mid-infrared excess. As the primary Cohen standard for the Southern hemisphere \\citep{1992AJ....104.1650C}, Sirius A is commonly observed as a photometric calibration for a variety of T-ReCS programs. By co-adding these data, we obtained deep mid-infared images, allowing us to detect the very faint source, Sirius B, at high signal-to-noise, and with the redundancy of independent detections in 5 filters. ", "conclusions": "We used archival Gemini/T-ReCS data to directly image the nearest white dwarf, Sirius B, in 5 filters in the N-band (10$\\micron$) window. The data were taken over a several year period, where Sirius A was used as a photometric standard for many T-ReCS programs. Because of Sirius B's non-negligable orbital motion during that timespan, we shifted each image by the binary's calculated orbital ephemeris in order to stack the images on Sirius B. Although \\citet{2008A&A...489..651B} reported a slight excess in the Ks-band, we find no evidence of a large mid-infrared excess, as would be expected for a DAZd white dwarf with a dusty debris disk \\citep{2007ApJ...662..544V}. White dwarfs in Sirius-like binary systems might be good targets for observing the outer parts of these debris disks. Because of the low-luminosity of white dwarfs, circumstellar material at the separations typically observed in debris disks are too cold to emit significant radiation in the mid-infrared. However, in binary systems, regions of the white dwarf's disk that are not heated by the white dwarf can be heated by its more luminous companion." }, "1101/1101.3338_arXiv.txt": { "abstract": "The standard analysis of the CMB data assumes that the distance to the last scattering surface can be calculated using the distance-redshift relation as in the Friedmann model. However, in the inhomogeneous universe, even if $\\av{\\delta\\rho} =0$, the distance relation is not the same as in the unperturbed universe. This can be of serious consequences as a change of distance affects the mapping of CMB temperature fluctuations into the angular power spectrum $C_l$. In addition, if the change of distance is relatively uniform no new temperature fluctuations are generated. It is therefore a different effect than the lensing or ISW effects which introduce additional CMB anisotropies. This paper shows that the accuracy of the CMB analysis can be impaired by the accuracy of calculation of the distance within the cosmological models. Since this effect has not been fully explored before, to test how the inhomogeneities affect the distance-redshift relation, several methods are examined: the Dyer--Roeder relation, lensing approximation, and non-linear Swiss-Cheese model. In all cases, the distance to the last scattering surface is different than when homogeneity is assumed. The difference can be as low as 1\\% and as high as 80\\%. Excluding extreme cases, the distance changes by about 20 -- 30\\%. Since the distance to the last scattering surface is set by the position of the CMB peaks, in order to have a good fit, the distance needs to be adjusted. After correcting the distance, the cosmological parameters change. Therefore, a not properly estimated distance to the last scattering surface can be a major source of systematics. This paper shows that if inhomogeneities are taken into account when calculating the distance then models with positive spatial curvature and with $\\Omega_\\Lambda \\sim 0.8-0.9$ are preferred. The $\\Lambda$CDM model (i.e. a flat Friedmann solution with the cosmological constant), in most cases, is at odds with the current data. ", "introduction": "It has been said that we entered the era of precision cosmology. This is mainly due to observations of the cosmic microwave background radiation (CMB). As the CMB power spectrum is very sensitive to cosmological parameters it provides very tight constraints. Currently the errors are at the level of a few percent \\cite{WMAP7}. The CMB power spectrum is shaped by several processes that can be divided into two groups. The first group involves processes that occurred before and during the generation of the CMB. This part is well understood as the Universe at that times was very close to homogeneous and contribution from the spatial curvature and dark energy was negligible. The constraints on cosmological parameters coming just from this group of processes ware recently presented and discussed in \\cite{VoRD2010}. The second group of processes takes into account what happens with photons between the last scattering surface and the observer. This includes the distance to the last scattering surface, the ISW effect, the impact of reionization on the CMB, the Sunyaev-Zel'dovich and the lensing effects. The analysis of this second group of effects is a subject to systematics, as the geometry of the late time Universe does not have to be close to homogeneous and isotropic Robertson-Walker geometry (these processes are most often studied within the framework of linear perturbations around the homogeneous Friedmann models). On small scales (say $\\ell > 60$, i.e. $\\theta<3^\\circ$), expect for the distance, these processes should not significantly modify the overall shape of the CMB power spectrum. The distance to the last scattering surface affects the mapping of CMB temperature fluctuations into the angular power spectrum $C_l$. This effect can be modelled by introducing the shift parameter ${\\cal R}$, which is proportional to the distance to the last scattering surface \\cite{BoET1997,EfBo1999}. Therefore, there is a number of papers where the constraints from the CMB are just limited to the constraints on the shift parameter. This is usually done when testing different models of dark energy \\cite{MMOT2003,KuSz2008}, alternative cosmological models \\cite{LaMM2006,RyFG2007}, or Gpc-scale inhomogeneous models \\cite{AAG06,AlAm2007,BoWy2009,ClFZ2009,ABNV2009,ZaMS2008,ClRe2010,YoNS210,MZS10,BNV10}. As discussed in \\cite{WaMu2007,ElMu2007} this type of analysis can be improved but taking into account both the scale of the sound horizon at the last scattering and ${\\cal R}$. This paper studies the effect of the distance but does not consider any alternative cosmologies. Here we just focus on the effect of small-scale inhomogeneities on the distance to the last scattering. The presence of inhomogeneities is know to affects the distance-redshift relation \\cite{Sach1961,KrSa1966}. Some studies claim that the effect is large \\cite{KaVB1995,MKMR2007,MaKM2008,KaMa2009}, others that it is small \\cite{BrTT2007,BrTT2008,VaFW2008,WV09,Szyb2010}. Even if the change of the distance just at the level of a few percent \\cite{ClZu2009,ClFe2009,kbAA,kbMNRAS} this analysis is still important as a change of the distance by a few percent leads to a similar change of the inferred values of cosmological parameters. One may think that the effect of inhomogeneities is taken into account in the standard analysis of the CMB via the ISW or lensing effects. However, in both cases (in the standard approach) this is done by using the matter power spectrum. Thus, this analysis is insensitive to the change of the mean, i.e. the uniform change of the distance to the last scattering surface. For example the lensing analysis deals with the change of the distance but is only sensitive to the change of the variance, not the mean (cf. \\cite{Lcmb1,Lcmb2}). If the change of distance is relatively uniform (i.e. with a negligible variance) then no additional temperature anisotropies are generated. Since the change of distance affects the mapping of the physical position of the peaks to the peaks in the angular power spectrum $C_l$ this effect needs to be considered also at the level of the mean. Therefore, if inhomogeneities are not taken into account when calculating the distance, then they may become a major source of systematics. With the increasing precision of CMB experiments, in particular Planck, a proper understanding of systematics is important. Without taking into account all systematics the precision cosmology will not be an accurate cosmology. Therefore, this paper aims to study how the presence of inhomogeneities affects the distance-redshift relation, in particular the distance to the last scattering surface. Several models are considered and it is shown that the change of distance due to inhomogeneities is not negligible. The structure of this paper is as follows: Sec. \\ref{distance} discusses different methods for the distance calculation, Sec. \\ref{dissec} applies these methods to calculate the distance to the last scattering surface, Sec. \\ref{labdc} explores the implication of differences in the distance, and Sec. \\ref{conk} concludes the results. ", "conclusions": "\\label{conk} This paper investigated how the presence of inhomogeneities (small-scale inhomogeneities observed in the Universe) affects the distance to the surface of last scattering, and how via the change of the distance the inhomogeneities influence the CMB data. In the standard approach, the distance is calculated within a framework of homogeneous Friedmann models. There are several reasons for using this approach. First of all, there is an argument used by Weinberg \\cite{Wein1976} who pointed out that although for a single case the distance is modified by the inhomogeneities, but due to photon conservation, when averaged over large enough angular scales the overall effect is zero (however see \\cite{ElBD1998} for an argument why Weinberg's reasoning should not apply). Second of all, as seen from (\\ref{dDBa}) if the amount of voids is the same as amount of overdensities then $\\Delta_D$ should be zero and the distance should be exactly the same as in the Friedmann model. As density fluctuations should vanish after averaging over sufficiently large scales it seems reasonable to expect that the Friedmann distance-redshift relation correctly describes the reality. However, there is a difference between vanishing 3D average and vanishing of the average of the density fluctuations over the line of sight. If density fluctuations in the Universe are purely random then vanishing 3D average implies vanishing average over the line of sight. If density fluctuations are not purely random, i.e. if there is some degree of organization, then vanishing 3D average does not necessarily imply vanishing density fluctuations over the line of sight (for a discussion see \\cite{kbMNRAS}). As the large-scale matter distribution in the Universe has a form of the cosmic web, there is some kind of organization, and thus the result of (\\ref{dDBa}) does not have to be zero, even if 3D average of density fluctuations vanishes after averaging over sufficiently large scales. This is the case of model Len 2. Also, relation (\\ref{dDBa}) does not apply when structures are non-linear. The calculations within model SC were carried out within the non-linear regime. In addition matter distribution was chosen so that the density fluctuations along the line of sight vanish after averaging (i.e. $\\av{\\delta}_{1D} = 0$). In this case, again, the distance was not the same as within the Friedmann model. Even if the change of the distance is small, like in Len 1 model, still the effect is important, as we have already entered the era of precision cosmology -- as seen from Table \\ref{tab1} and Table \\ref{tab2} the change of the inferred cosmological parameters, even for Len 1 model, can be of order of a percent or more. If the change of the distance is larger, for example as in models Len 2 and SC, then the inferred parameters are significantly different. In this case spatially flat models are ruled out and the data favors positively curved models with with $\\Omega_\\Lambda \\approx 0.8-0.9$. As discussed in Sec. (\\ref{dissec}), within a chosen method, the distance to the last scattering surface is relatively uniform with only small variance. The presented, in this paper, results show that the change of the mean is important, and the proper handling of this effect is essential, otherwise the systematics may be larger than the precision of measurements. The only problem, however, is to have a good method for the calculation of the distance. All models presented in this paper have their limitations, and should rather be treated as toy models, i.e. as examples to show the significance of the problem. The results show that the accuracy of the CMB analysis strongly depends on the accuracy of the calculation of the distance corrections. Peebles described the importance of accuracy and precision in cosmology with the following example {\\it a digital scale may read out the weight of an object to many significant figures, in a precise measurement. But if the scale is not well calibrated the measurement may not be very accurate} \\cite{Peeb2010}. In our case, the `calibration' requires an inhomogeneous framework that will allow us to estimate the distance with inhomogeneities taken into account. As the non-linear corrections are important this cannot be done within the lensing approximation (\\ref{dDBa}) but requires a fully inhomogeneous framework. Only then we will be able to accurately predict the distance and say: `the distance is larger by $29.72\\%$' or `$36.14\\%$ than when homogeneity is assumed'. With the increasing precision of CMB experiments, in particular Planck, Atacama Cosmology Telescope, and South Pole Telescope this is important. Otherwise, what would be the meaning of a measurement with very high precision ($\\sim 1\\%$ statistical errors) if the accuracy is low ($\\sim 30\\%$ systematics)? \\ack I would like to thank Richard Bond, Chris Clarkson, Timothy Clifton, Ruth Durrer, Pedro Ferreira, Valerio Marra, Teppo Mattsson, Syksy R\\\"as\\\"anen, Marco Regis, David Wiltshire for useful discussions and suggestions. This research of was supported by the Marie Curie Fellowship (PIEF-GA-2009-252950)." }, "1101/1101.1567_arXiv.txt": { "abstract": "VDB0-B195D is a massive, blue star cluster in M31. It was observed as part of the Beijing-Arizona-Taiwan-Connecticut (BATC) Multicolor Sky Survey using 15 intermediate-band filters covering a wavelength range of 3000--10,000 \\AA. Based on aperture photometry, we obtain its spectral-energy distribution (SED) as defined by the 15 BATC filters. We apply previously established relations between the BATC intermediate-band and the Johnson-Cousins $UBVRI$ broad-band systems to convert our BATC photometry to the standard system. A detailed comparison shows that our newly derived $VRI$ magnitudes are fully consistent with previous results, while our new $B$ magnitude agrees to within $2\\sigma$. In addition, we determine the cluster's age and mass by comparing its SED (from 3000 to 20,000{\\AA}, comprising photometric data in the 15 BATC intermediate bands, optical broad-band $BVRI$, and 2MASS near-infrared $JHK_{\\rm s}$ data) with theoretical stellar population synthesis models, resulting in age and mass determinations of $60.0\\pm 8.0$~Myr and $(1.1-1.6) \\times 10^5 M_\\odot$, respectively. This age and mass confirms previous suggestions that VDB0-B195D is a young massive cluster in M31. ", "introduction": "Young massive star clusters (YMCs) are among the main objects resulting from violent star-forming episodes triggered by galaxy collisions, mergers, and close encounters \\citep[see][and references therein]{grijs07}. They are also referred to as `young populous clusters,' a term first coined by \\citet{hodge61}, who used it to describe 23 clusters containing bright, blue stars in the Large Magellanic Cloud. In \\citet{hodge61}, the `young' aspect is demonstrated by the fact that all clusters have main sequences that extend to absolute magnitudes brighter than $M_V=0$, while `populous' describes their richness (stellar membership). However, YMCs are also observed in quiescent galaxies \\citep{lr99} and in the disks of isolated spirals, although higher cluster-formation efficiencies are associated with environments exhibiting high star-formation rates \\citep[see][and references therein]{larsen04a,cw07}. It has become clear that, in many ways, YMCs resemble young versions of the old globular clusters (GCs) associated with all large galaxies \\citep[see][and references therein]{larsen04b}. YMCs are seemingly absent in the Milky Way; possibly the best example of a Galactic YMC is Westerlund 1, a heavily reddened cluster with an age and mass of 4--5 Myr \\citep{Crowther06} and $M_{\\rm cl} \\sim 10^5~M_\\odot$ \\citep{Clark05}, respectively. Since the pioneering work of \\citet{Tinsley68,Tinsley72} and \\citet{SSB73}, evolutionary population synthesis modeling has become a powerful tool to interpret integrated spectrophotometric observations of galaxies and their components, such as star clusters \\citep[e.g.,][]{Anders04}. The evolution of star clusters is usually modeled by means of the simple stellar population (SSP) approximation. An SSP is defined as a single generation of coeval stars formed from the same progenitor molecular cloud (thus implying a single metallicity), and governed by a given stellar initial mass function (IMF). Age and metallicity are two basic star cluster parameters. The most direct method to determine a cluster's age is by employing main-sequence photometry, since the absolute magnitude of the main-sequence turnoff is predominantly affected by age \\citep[see][and references therein]{puzia02}. However, until recently \\citep[cf.][]{perina09}, this method was only applied to the star clusters in the Milky Way and its satellites \\citep[e.g.,][]{rich01}, although \\citet{brown04} estimated the age of an M31 GC using extremely deep images observed with the {\\sl Hubble Space Telescope (HST)}'s Advanced Camera for Surveys. Generally, the ages of extragalactic star clusters are determined by comparing their observed spectral-energy distributions (SEDs) and/or spectroscopy with the predictions of SSP models \\citep{Williams01a,Williams01b,degrijs03a,degrijs03b,degrijs03c,bik03,jiang03,Beasley04,puzia05,ma06,fan06,ma07a,ma09,cald09, Wang10}. Nevertheless, SSP models assume that cluster IMFs are fully populated, i.e., that clusters contain infinite numbers of stars with a continuous distribution of stellar masses, and that all evolutionary stages are well sampled. Real clusters, however, contain a finite number of stars. Therefore, a disagreement between the observed cluster colors and theoretical colors derived from SSP models may become apparent \\citep[see][and references therein]{pisk09,ph10}. Other limitations inherent to SSP models arise from our poor understanding of some advanced stellar evolutionary stages, such as the supergiant and the asymptotic-giant-branch (AGB) phases \\citep[see][and references therein]{bc03}. Located at a distance of $785 \\pm 25$ kpc, corresponding to a distance modulus of $(m-M)_0 = 24.47 \\pm 0.07$ mag \\citep{McConnachie05}, M31 is the nearest and largest spiral galaxy in the Local Group of galaxies. It has been the subject of many GC studies and surveys, dating back to the early study of \\citet{hubble32}. Based on previous publications \\citep{hubble32,seynas45,hilt58,mayegg53,km60}, \\citet{vete62a} compiled the first large M31 GC catalog, containing $UBV$ photometric data of approximately 300 GC candidates. Over the past decades, several major catalogs of M31 GCs and GC candidates have been published, including major efforts by the Bologna group \\citep{batt80,batt87,batt93}, \\citet{bh00}, \\citet{gall04,gall05,gall06,gall07}, \\citet{kim07}, \\citet{cald09}, and \\citet{peacock09}. Following on from the first extensive spectroscopic survey of M31 GCs by \\citet{Sidney69}, a significant number of authors \\citep[e.g.,][and references therein]{hsv82,hbk91,DG,Federici93,jab,bh00,per02,gall06,lee08} have studied their spatial, kinematic, and chemical (metallicity) properties. M31 is known to host a large number of young star clusters \\citep[e.g.,] [and references therein]{Fusi05,cald09,Wang10}. \\citet{Fusi05} presented a comprehensive study of 67 very blue star clusters, which they referred to as `blue luminous compact clusters' (BLCCs). Since they are quite bright ($-6.5\\leq M_V \\leq -10.0$ mag) and very young ($<2$ Gyr), BLCCs may be equivalent to YMCs \\citep[see for details][]{perina09,perina09b}. To ascertain their properties, \\citet{perina09,perina09b} performed an imaging survey of 20 BLCCs in the disk of M31 using the {\\sl HST}'s Wide Field and Planetary Camera-2 (WFPC2). They obtained the reddening values, ages, and metallicities of their sample clusters by comparing the observed color-magnitude diagrams (CMDs) and luminosity functions with theoretical models. VDB0-B195D was first detected by \\citet{Sidney69}. Its color is extremely blue \\citep[e.g., $U-B=-0.48$ mag;][]{Sidney69} and it is very bright in blue bands \\citep[e.g., $U=14.66$ mag;][]{Sidney69}. As a consequence, \\citet{Sidney69} asserted that VDB0-B195D is the brightest open cluster in M31. He determined an integrated stellar spectral type equivalent to A0, which implies that the cluster contains massive stars. In addition, VDB0-B195D is particularly extended and most previous photometric studies did not include the full extent of the object's light distribution \\citep[see for details][]{perina09}. We will provide an overview of previous studies that included the cluster in \\S 2.1. It was observed as part of the galaxy calibration program of the Beijing-Arizona-Taiwan-Connecticut (BATC) Multicolor Sky Survey \\citep[e.g.,][]{fan96,zheng99} in 15 intermediate-band filters. Combined with photometry in optical broad-band $BVRI$ and near-infrared $JHK_{\\rm s}$ filters from the Two Micron All Sky Survey (2MASS) taken from \\citet{perina09}, we obtained the SED of VDB0-B195D in 22 filters, covering the wavelengh range from 3000 to 20,000 \\AA. In this paper, we describe the details of the observations and our approach to the data reduction in \\S 2. In \\S 3, we determine the age and mass of VDB0-B195D by comparing observational SEDs with population synthesis models. We discuss the implications of our results and provide a summary in \\S 4. ", "conclusions": "VDB0-B195D was previously shown to be a massive cluster based on {\\sl HST}/WFPC2 observations. Its color is extremely blue and it is very bright, particularly in blue bands. In addition, VDB0-B195D is an extended object, and most previous photometric measurements did not include its full flux distribution \\citep[see][for details]{perina09}. In this paper, we obtained the cluster's SED in the 15 BATC intermediate-band filters. We subsequently determined its age and mass by comparing our multicolor photometry with theoretical stellar population synthesis models. Our multicolor photometric data consist of 15 intermediate-band filters obtained in this paper, and broad-band $BVRI$ and 2MASS $JHK_{\\rm s}$ from \\citet{perina09}, covering a wavelength range from 3000 to 20,000~\\AA. Our results show that VDB0-B195D is a genuine YMC in M31. To understand the real nature of the BLCCs, \\citet{perina09,perina09b} performed an {\\sl HST} imaging survey of 20 BLCCs in M31's disk. As a test case, \\citet{perina09} presented details of the data-reduction pipeline that will be applied to all survey data and describe its application to VDB0-B195D. They estimated the object's age, by comparison of the observed CMD with theoretical isochrones from \\citet{Girardi02}, at $\\simeq 25$ Myr. In addition, they constrained realistic upper and lower limits to the cluster's age, independent of the adopted metallicity, within the relatively narrow range from 12 to 63 Myr. Using Maraston's SSP models of solar metallicity \\citep{Maraston98,Maraston05}, \\citet{salp55} and \\citet{Kroupa01} IMFs, and photometric values in the $V$ and 2MASS $J$, $H$, and $K_{\\rm s}$ bands, \\citet{perina09} concluded that the mass of VDB0-B195D is $>2.4\\times 10^4~M_\\odot$, with their best estimates in the range $\\simeq (4-9)\\times 10^4~M_\\odot$. \\citet{cald09} presented an updated catalog of 1300 objects in M31, including spectroscopic and imaging surveys, based on images from either the LGGS or the DSS and spectra taken with the Hectospec fiber positioner and spectrograph on the 6.5 m MMT. They derived ages and reddening values for 140 young clusters by comparing their observed spectra with model spectra from the Starburst99 SSP suite \\citep{Leitherer99}. The results show that these clusters are less than 2 Gyr old, while most have ages between $10^8$ and $10^9$ yr (the age of VDB0-B195D they derive is $\\log {\\rm age/yr} =7.6$). In addition, \\citet{cald09} also estimated the masses of these young clusters using $V$-band photometry and model mass-to-light ratios \\citep{Leitherer99} corresponding to the derived spectroscopic ages. This resulted in masses ranging from $2.5 \\times 10^2$ to $1.5 \\times 10^5~M_\\odot$. The mass of VDB0-B195D obtained by \\citet{cald09} is $\\log M_{\\rm cl}/M_\\odot =5.1$ (no uncertainty quoted). We compare the various age and mass estimates of VDB0-B195D in Table 6. Our newly obtained age is older than the estimates of both \\citet{perina09} and \\citet{cald09}, while the mass obtained in this paper is higher than the estimate of \\citet{perina09} and consistent with the determination of \\citet{cald09}. However, our results are in agreement with those of both \\citet{perina09} and \\citet{cald09} within $3\\sigma$. The age and mass obtained in this paper confirms that VDB0-B195D is genuinely a YMC in M31. As we know, SSP models describe a very special case of a continuous distribution of stellar mass (or light) along isochrones. This is well approximated by clusters with masses larger than $10^6~M_\\odot$. Also, for cluster masses of about $10^5~M_\\odot$, SSP models can probably still be applied since a systematic difference between SSP models and observations should, on average, be smaller than 0.05 mag for clusters older than 10 Myr (see Fig. 3 in Piskunov et al. 2009). However, from the results of this paper, we may conclude that, probably, a formal fitting of SSP models to observed SEDs cannot be used without caution even for relatively massive (or apparently massive) clusters, and it is highly doubtful that this approach can be applied in a routine work providing accurate cluster parameters. The relative accuracy of 10\\% for age and 20\\% found for the mass of VDB0-B195D seems to be rather formal and not very confident. In addition, observational star clusters' SEDs are affected by reddening, an effect that is also difficult to separate from the combined effects of age and metallicity (Calzetti 1997; Vazdekis et al. 1997; Origlia et al. 1999). Only the metallicity and reddening are derived accurately (and, ideally, independently), these degeneracies are largely (if not entirely) reduced, and ages can then also be estimated accurately based on a comparison of multicolor photometry spanning a significant wavelength range (de Grijs et al. 2003b; Anders et al. 2004) with theoretical stellar population synthesis models. It is true that the discrepancy between our observations and the best-fitting model is great, and the mass of VDB0-B195D obtained based on the magnitudes in different filters is very different. However, when we adopt a smaller reddening value, the results improve greatly. So, we conclude that the actual reddening value of VDB0-B195D may be smaller than $E(B-V)=0.2$." }, "1101/1101.1084_arXiv.txt": { "abstract": "This paper describes the image stacks and catalogs of the CFHT Legacy Survey produced using the MegaPipe data pipeline at the Canadian Astronomy Data Centre. The Legacy Survey is divided into to two parts: The Deep Survey consists of 4 fields each of 1 square degree, with magnitude limits (50\\% completeness for point sources) of $u=27.5$, $g=27.9$, $r=27.7$, $i=27.4$, and $z=26.2$, and contains $1.6\\times 10^6$ sources. The Wide Survey consists of 150 square degrees split over 4 fields, with magnitude limits of $u=26.0$, $g=26.5$, $r=25.9$, $i=25.7$, and $z=24.6$, and contains $3\\times10^7$ sources. This paper describes the calibration, image stacking and catalog generation process. The images and catalogs are available on the web through several interfaces: normal image and text file catalog downloads, a ``Google Sky'' interface, an image cutout service, and a catalog database query service. ", "introduction": "\\label{sec:intro} The Canada France Hawaii Telescope Legacy Survey (CFHTLS) consisted of three surveys: The Deep Survey imaged 4 square degrees; the primary goal was to measure cosmological parameters using Type Ia supernovae as standard candles. This portion of the CFHTLS is also known as the Super Nova Legacy Survey \\cite[SNLS]{SNLS1}. The Wide survey covered a wider area to shallower depth. It consists of 171 overlapping pointings. which (after accounting for overlaps) cover 150 square degrees. Weak lensing and large scale structure studies were the primary science goals. The Very Wide portion of the survey covered the ecliptic plane and was designed to study Kuiper Belt Objects. It was cancelled mid-way through the project but has become the Canada-France Ecliptic Plane Survey \\cite[CFEPS]{cfeps2009}. Interestingly, image stacks are secondary to the main science goals of the CFHTLS. The main reason for the Deep Fields is supernova cosmology. The supernovae are discovered using stacks generated on a nightly basis. Deeper stacks are generated but only as a reference for difference imaging. Similarly, the main science driver of the Wide fields is weak lensing. The actual weak lensing measurements are done on individual images in order to accurately quantify the PSF. The value of the Very Wide survey, which was looking for moving objects, is also in the individual images. Now that the primary science goals of the CFHTLS are well in hand, the archival data becomes important. Creating stacks and catalogs greatly increases the usefulness of archival MegaCam data. MegaPipe \\citep{gwyn2008} has been used to process over 2000 square degrees of MegaCam imaging over the last four years. Twice a year it goes through the CFHT archive at the CADC, determines which images are worth combining, computes accurate astrometric/photometric calibrations, stacks the images, and makes the stacked images to the astronomical community. In addition, it is used process PI data (on request) and to process data from large surveys such as the Next Generation Virgo Survey (Ferrarese et al. in preparation) This paper describes image stacks and catalogs of the CFHTLS Deep and Wide surveys generated by the MegaPipe image processing pipeline. It presents the astrometric and photometric calibration, the image stacking and catalog generation, and quality control. This paper outlines the various web applications that astronomers can use to download the data. This paper however does not describe the stacking of the Very Wide data; everything worth stacking was stacked as part of regular semi-annual MegaPipe processing. Some of the methods presented in this paper have already been described in the main MegaPipe paper \\citep{gwyn2008}. For completeness, this paper describes these methods again (usually in less detail), but more emphasis is placed on methods that are specific to the CFHTLS stack generation and quality control. The AB magnitude \\citep{abmag} system is used throughout this paper. ", "conclusions": "" }, "1101/1101.4945_arXiv.txt": { "abstract": "{ Recent studies of different X-ray binaries (XRBs) have shown a clear correlation between the radio and X-ray emission. We present evidence of a close relationship found between the radio and X-ray emission at different epochs for GRS\\,1915+105, using observations from the Ryle Telescope and Rossi X-ray Timing Explorer satellite. The strongest correlation was found during the hard state (also known as the `plateau' state), where a steady AU-scale jet is known to exist. Both the radio and X-ray emission were found to decay from the start of most plateau states, with the radio emission decaying faster. An empirical relationship of $S_{\\rm{radio }}\\propto S_{\\rm{X-ray}}^{\\xi}$ was then fitted to data taken only during the plateau state, resulting in a power-law index of $\\xi\\sim1.7\\pm0.3$, which is significantly higher than in other black hole XRBs in a similar state. An advection-flow model was then fitted to this relationship and compared to the universal XRB relationship as described by Gallo et al. (2003). We conclude that either (I) the accretion disk in this source is radiatively efficient, even during the continuous outflow of a compact jet, which could also suggest a universal turn-over from radiatively inefficient to efficient for all stellar-mass black holes at a critical mass accretion rate ($\\dot{m}_{\\rm{c}}\\approx10^{18.5}$~g/s); or (II) the X-rays in the plateau state are dominated by emission from the base of the jet and not the accretion disk (e.g. via inverse Compton scattering from the outflow). } {} ", "introduction": "It has been well established that the formation of astrophysical jets is associated with accretion disks. Powerful jets have been observed from sizes that range from the AGN of extra-galactic quasars, to the much smaller ejecta around proto-planetary disks. Whilst the theory of accretion has been the focus of intense study for many decades, there exists no satisfactory explanation for this scale-free, invariant relationship between accretion disks and jets. This paper presents a clear relationship between the accretion disk and jet formation in one of the most powerful black holes in the Galaxy, GRS\\,1915+105. Using the high-resolution results and conclusions presented in \\cite{2010MNRAS.401.2611R}, a confident assumption regarding the `state' of the jet can be made using radio and X-ray monitoring observations. As variations in the X-ray emission are associated with changes in the accretion disk, coupling these radio observations to simultaneous X-ray observations can directly relate to the inflow-outflow mechanisms around the black hole~\\cite[reviewed by][]{2004ARA&A..42..317F}. It has been clearly established that \\textit{changes} in the X-ray spectral state of GRS\\,1915+105 are associated with superluminal knots~\\citep{1999MNRAS.304..865F}; however, large-scale structures quickly become divorced from the central accretion region as they move away and typically last for many days or weeks. To study the direct coupling between the accretion disk and jet, we must select observations of the compact `steady jet' when it is close to the energy source. As VLBI observations have shown this emission region to be only a few light hours in size~\\citep{2000ApJ...543..373D}, changes in the accretion disk may therefore directly change the steady jet over the time-scale of hours. For the first time, observations taken over ten years have made it possible to clearly study the repetitive trends between the X-ray and radio. The results presented in this paper select radio and X-ray monitoring observations during the steady jet (`plateau') state. A direct relationship between the inflowing accretion and outflow jet can be made for GRS\\,1915+105. \\subsection{X-ray spectral states of GRS\\,1915+105} \\label{sec:X-ray_plateau} The X-ray spectra of GRS 1915+105 can be generalised into two distinct components: a soft blackbody disk ($kT\\sim1-2$~keV) and a hard non-thermal power-law extending to $\\geq100$~keV associated with the very inner region or `corona' of the accretion disk. \\cite{2000A&A...355..271B} initially identified 12 distinct X-ray classes with the RXTE-PCA based on the light-curve and colour-colour diagram of the source. From these classes, the overall X-ray spectra can be reduced as a transition between three basic states known as \\textit{state A}, \\textit{state B} and \\textit{state C}; these states represent the interchange between the multi-temperature disk and the power-law component. States A and B correspond to a soft energy spectrum, with the dominant X-ray emission coming from the inner region of a thermal accretion disk with temperatures of $\\sim1.8$~keV and $\\sim2.2$~keV respectively. In-contrast, state C represents the near absence of this inner region and exhibits a dominant power-law component in the X-ray emission. \\cite{2003A&A...404..283M} performed an analysis of the X-ray spectra with the RXTE-PCA, by isolating the observations into states A/B/C. They fitted a blackbody component to the energy spectrum and directly related this to the inner radius ($R_{\\rm{in}}=D\\sqrt{N\\cos\\theta})$, where $D$ is the distance of the source, $\\theta$ the inclination angle of the disk and N the normalisation of the thermal component. The inner radius was then compared to the overall temperature, $kT_{\\rm{in}}$. It was found at the start of state C that $R_{\\rm{in}}$ would rapidly increase in size and $kT_{\\rm{in}}$ would drop to a much cooler temperature. This was explained as the collapse of the inner accretion region, leaving a truncated thin disk. During the evolution of the $R_{\\rm{in}}-kT$ relationship in the state C, the temperature would slowly increase again and $R_{\\rm{in}}$ would reduce. This has been suggested to be the re-filling of the inner region of the thin disk, until the return of the thermally dominant state of the XRB. A detailed spectral observation of the state C by XMM-Newton was performed by \\cite{2006A&A...448..677M}. They found only a simple power-law of $\\Gamma\\sim1.7$ was needed to fit to the overall spectra, which then gave residuals of a 1~keV excess, small variations between $\\sim1.5-3$~keV and a deficit above 8~keV. The overall power-law was consistent with the RXTE-PCA observations, and was attributed to either a hot corona around the accretion disk or to Comptonized emission from the base of a jet \\citep{2004ApJ...615..416R}. The 8~keV deficit was explained as an optically thick reflector that gave evidence of the presence of a thin disk. The 1~keV excess was (tentatively) explained as the presence of an optically thin component (e.g. a wind/jet or a geometrically thick disk). \\cite{2006A&A...448..677M} noted that the relatively large amount of reflection components imply that the primary X-ray emitting region would have a size comparable to the inner disk radius. \\cite{2002MNRAS.331..745K} clearly established that radio emission is intimately related to the hard X-ray power-law component and implies a close physical connection between the inner accretion region and the outflowing synchrotron-emitting jet. They found a `one-to-one' relationship between radio oscillation events, originally discovered by \\cite{1997MNRAS.292..925P}, and quasi-periodic X-ray dips during certain epochs in state C; in all other cases the source showed either `low-level' or `high-level' radio emission, but no radio oscillation \\cite[see][for a detailed characterisation of the radio oscillation events]{2010ApJ...717.1222P}. Moreover, when the source stays in spectral hard state C for long periods of days to months, the `high-level' radio emission (i.e. non-oscillating emission) has been shown to originate from the compact flat-spectrum jet. This state is also known as the `plateau' state and is also referred to as class~$\\chi$ by \\cite{2000A&A...355..271B}. \\subsection{A universal X-ray binaries relationship and the fundamental plane} X-ray binaries (XRBs) are ideal objects to study the disk-jet coupling relationship, as they display bright X-ray and radio emission from an accretion disk and jet, respectively. \\cite{2000A&A...359..251C} found a clear relationship in the low-hard state of the XRB GX~339-4 between the X-ray and radio emission. The observed flat (or slightly inverted) spectrum was suggested to be a compact jet associated with the low-hard state. Discrete ejections of relativistic plasma are associated with transitional changes from this hard state to a softer state for most radio emitting black hole XRBs~\\citep{2004MNRAS.355.1105F,2009MNRAS.396.1370F}. An empirical non-linear relationship between the radio and X-ray luminosity was found by \\cite{2003A&A...400.1007C} as \\begin{equation} L_{\\rm{radio}}\\propto L_{\\rm{X-ray}}^{0.71\\pm0.1}, \\end{equation} \\noindent whilst the source remained in the low-hard state. \\cite{2003MNRAS.344...60G}, then compiled a large sample of quasi-simultaneous radio and X-ray observations of stellar-mass black hole binaries, confirming this relationship to be universal for most black hole XRBs. Following this, \\cite{2003MNRAS.345.1057M} and \\cite{2004A&A...414..895F} independently linked the existence of a `fundamental plane' associated with all accreting black holes, with the parameters $L_{\\rm{radio}}$, $L_{\\rm{X-ray}}$ and mass $M$. This relationship applied to a large range of masses, from super-massive black holes in extragalactic AGNs and the Galactic centre, Sgr~A$^{*}$, to stellar mass black holes in XRBs. This fundamental plane was found to take the form \\begin{equation} L_{\\rm{radio}}\\propto L_{\\rm{X-ray}}^{0.6}M^{0.8}. \\end{equation} The aim of this paper is to test the fundamental relationship between the inflow and outflowing emission, with the source variability found in GRS\\,1915+105. ", "conclusions": "This work has shown, for the first time, a direct relationship between the X-rays and radio in the steady jet state of GRS\\,1915+105. Previous attempts have failed to show this relationship~\\citep[e.g.][]{2001ApJ...556..515M}, as they have included X-ray/radio comparisons that include extended knots and or X-ray accretion in other states. \\cite{2003MNRAS.344...60G} observed a universal radio-X-ray correlation in low/hard state black holes, but only for radiatively inefficient accretion. Figure~\\ref{fig:universal_plane} shows the difference between the two models apply to other stellar-mass black holes and GRS\\,1915+105. The difference between the two models is likely to be due to the rate of accretion or spin of the black holes. GRS\\,1915+105 is known to be in a constant `soft'-like state as a large accretion rate is constantly present; however, it remains unclear if the X-ray emission, in the steady jet state, is produced from either the thin disk, an advection dominated flow or the compact jet. For other XRBs, the bright sources are likely to form only a transient soft-accretion disk and GRS\\,1915+105-type accretion may only occur in systems with higher accretion rates, like AGN. It is interesting that within each steady jet state, both the X-ray and radio luminosities fell with time. This suggests a cooling of the mechanisms or conditions that initially created the steady jet. Furthermore, whilst a coupling of the accretion process to the outflowing jet via advection is possible, the simplest explanation is that both the X-ray and radio emission originate directly from the jet as suggested by~\\cite{2005ApJ...635.1203M}." }, "1101/1101.1751_arXiv.txt": { "abstract": "{HC$_3$N is a molecule that is mainly associated with Galactic star-forming regions, but it has also been detected in extragalactic environments.} {To present the first extragalactic survey of HC$_3$N, when combining earlier data from the literature with six new single-dish detections, and to compare HC$_3$N with other molecular tracers (HCN, HNC), as well as other properties (silicate absorption strength, IR flux density ratios, \\ion{C}{ii} flux, and megamaser activity).} {We present mm IRAM~30~m, OSO~20~m, and SEST observations of HC$_3$N rotational lines (mainly the $J=$10\\nobrkhyph9 transition) and of the $J=$1\\nobrkhyph0 transitions of HCN and HNC. Our combined HC$_3$N data account for 13 galaxies (excluding the upper limits reported for the non-detections), while we have HCN and HNC data for more than 20 galaxies.} {A preliminary definition ``HC$_3$N-luminous galaxy'' is made based upon the HC$_3$N/HCN ratio. Most ($\\sim$80~$\\%$) HC$_3$N-luminous galaxies seem to be deeply obscured galaxies and (U)LIRGs. A majority ($\\sim$60~$\\%$ or more) of the HC$_3$N-luminous galaxies in the sample present OH mega- or strong kilomaser activity. A possible explanation is that both HC$_3$N and OH megamasers need warm dust for their excitation. Alternatively, the dust that excites the OH megamaser offers protection against UV destruction of HC$_3$N. A high silicate absorption strength is also found in several of the HC$_3$N-luminous objects, which may help the HC$_3$N to survive. Finally, we find that a high HC$_3$N/HCN ratio is related to a high dust temperature and a low \\ion{C}{ii} flux.} {} ", "introduction": "Finding useful tracers of the interaction between the activity in galaxy nuclei and surrounding interstellar medium (ISM) is an important and growing aspect of current extragalactic molecular astronomy. In this context, single dish surveys of polar molecules such as HCN, HNC, HCO$^+$, and CS have been used to investigate possible correlations between molecular line ratios and type/intensity of activity \\citep[e.g.][]{kohno01,aalto02,imanishi04,gracia06,krips08,baan08}. For example, it has been suggested that an elevated HCN/HCO$^+$~1\\nobrkhyph0 line intensity ratio indicates the presence of an AGN \\citep{gracia06}. Around an active galactic nucleus (AGN) the chemistry is supposedly dominated by hard X-rays in an X-ray dominated region (XDR), and some chemical models predict an abundance enhancement of HCN paired with selective destruction of HCO$^+$ \\citep{maloney96} -- which could lead to an elevated HCN/HCO$^+$ line ratio (under the circumstances that the line ratio directly reflects the abundance ratio). However, more recent chemical models instead suggest that HCO$^+$ is enhanced in XDRs \\citep{meijerink05,meijerink07}, and HCO$^+$ is also expected to be under-abundant in regions of very young star formation \\citep{aalto08}, so the line ratio is ambiguous. Other molecular tracers could help resolve the dichotomy of the HCN/HCO$^+$ line ratio. The serendipitous discovery of the $J$=10\\nobrkhyph9 line of HC$_3$N near the HNC~$J$=1\\nobrkhyph0 line in a survey by \\citet{aalto02} led us to look more closely at this molecule. HC$_3$N is the simplest of the cyanopolyynes (carbon chains with an attached CN group) and is a grain chemistry product, in contrast to molecules such as HCN and HCO$^+$. HC$_3$N thrives in warm, dense shielded regions such as hot cores where abundances can reach 10$^{-8}$ or even higher, since it is easily destroyed by photo-dissociation \\citep{rodriguez98} and C$^+$ ions \\citep{prasad80}. Therefore, HC$_3$N line emission could be used to identify galaxies where star formation is in the early, embedded stage of its evolution. Recently, HC$_3$N was found in high abundance in the highly obscured galaxy \\object{NGC 4418} \\citep{aalto07a}, as well as the ULIRG \\object{Arp 220} \\citep{aalto02}. We have searched for HC$_3$N line emission in a sample of galaxies in various stages and types of activity: AGNs, starbursts, and ultraluminous galaxies (ULIRGs). In some of the galaxies the nature of the activity is elusive since it is embedded in huge columns of dust absorbing emission at optical and infrared wavelengths. In some cases, the extinction is so strong that no emission emerges at optical or IR wavelengths requiring us to probe the nature of the activity at radio and mm wavelengths. HC$_3$N has a rich mm and sub-mm wavelength spectrum consisting of a multitude of rotational and vibrational lines often appearing close to each other in the same band. Through its vibrational transitions, HC$_3$N responds strongly to the IR field from dusty nuclei \\citep{costagliola10a}. Therefore, combining the rotational and vibrational line information of HC$_3$N allows us to study the abundance of HC$_3$N (comparing with chemical models of XDRs and starbursts) as well as the intensity and temperature structure of the buried IR source. Rotational lines of vibrationally excited HC$_3$N have recently been discovered in a few galaxies (\\object{NGC 4418} \\citep{costagliola10a}, \\object{Arp 220} \\citep{martin10}, and \\object{IC 860} \\citep{costagliola10b}), therefore showing that it is important to take both radiative and collisional excitation into consideration when interpreting HC$_3$N line emission from IR luminous galaxies. It can also be noted that absorption lines of HC$_3$N has been found in a $z\\sim0.89$ galaxy located in front of the quasar \\object{PKS 1830-211} \\citep{henkel09}. \\subsection{Outline} Here, the first survey of extragalactic HC$_3$N data is presented. We report new HC$_3$N observations in 19 galaxies (detections in six of them), mainly (U)LIRGs and starburst galaxies, and complete this sample with data from all earlier extragalactic HC$_3$N emission line single-dish detections found in the literature. The aim of the study is to compare the HC$_3$N luminosity with other molecular tracers as well as galaxy properties to see if the presence of HC$_3$N can be used to predict other galaxy properties, e.g. the source of activity in the galaxy. In Section~\\ref{sec:hc3ninspace}, the general properties of HC$_3$N in space are discussed. In Section~\\ref{sec:obs} we present the new observations and discuss the collection of data from the literature. In Section~\\ref{sec:res} we present the results in terms of line intensities and line ratios. In Section~\\ref{sec:dis} we discuss the interpretation of the HC$_3$N results and compare them with silicate absorption strength (Section~\\ref{sec:pah}), OH megamaser activity (Section~\\ref{sec:meg}), IR flux density ratios (Section~\\ref{sec:iras}), \\ion{C}{ii} flux (Section~\\ref{sec:cii}), and the HNC/HCN~1\\nobrkhyph0 line ratio (Section~\\ref{sec:hnc}). In Section~\\ref{sec:fut} future studies resulting from this project are discussed. ", "conclusions": "\\label{sec:con} We have presented the first survey of HC$_3$N observations in extragalactic objects. The main conclusions from this survey are as follows: \\begin{enumerate} \\item Bright HC$_3$N emission is rather uncommon in galaxies. It was only detected in 6 of the 19 galaxies which had not been investigated before, even though that sample was selected to find many HC$_3$N-luminous galaxies. \\item Most HC$_3$N-luminous galaxies are obscured galaxies. Starburst galaxies seem to be poor in HC$_3$N. There are too few AGN galaxies in the sample to tell if these normally are rich or poor in HC$_3$N. \\item Weak correlations can be seen between the HC$_3$N/HCN ratio and silicate 9.7~\\hbox{\\textmu}m absorption strength. \\item There is a strong correlation between OH megamaser activity and HC$_3$N luminosity. Most HC$_3$N-luminous galaxies have an OH megamaser. This could be related to a high dust obscuration in the HC$_3$N-luminous galaxies. \\item There is a connection between the HC$_3$N/HCN ratio and the IRAS 60~\\hbox{\\textmu}m/100~\\hbox{\\textmu}m flux density ratios, indicating a higher dust temperature in these galaxies, which could cause vibrational excitation of the HC$_3$N molecule. \\item There is a strong connection between a high HC$_3$N/HCN ratio and a low \\ion{C}{ii}/FIR flux ratio in the studied objects. This could be explained by C$^+$ ions being required to destroy the HC$_3$N molecule. \\item There is a correlation between the HC$_3$N/HCN and HNC/HCN line ratios. \\end{enumerate}" }, "1101/1101.4342_arXiv.txt": { "abstract": "We propose that at the beginning of the Maunder minimum the poloidal field or amplitude of meridional circulation or both fell abruptly to low values. With this proposition, a flux transport dynamo model is able to reproduce various important aspects of the historical records of the Maunder minimum remarkably well. ", "introduction": "One important aspect of the solar cycle is the Maunder minimum during 1645--1715 when the solar activity was strongly reduced (Ribes \\& Nesme-Ribes 1993). It was not an artifact of few observations, but a real phenomenon (Hoyt \\& Schatten 1996). From the study of historical data (Ribes \\& Nesme-Ribes 1993), it has been confirmed that the sunspot numbers in both the hemisphere fell abruptly to nearly zero value at the beginning of the Maunder minimum, whereas a few sunspots appeared in the southern hemisphere during the last phase. It is also established from the cosmogenic isotopes data (Beer et al. 1998; Miyahara et al. 2004) that the cyclic oscillations of solar activity continued in the heliosphere at a weaker level during the Maunder minimum, but with a period of 13--15 years instead of the regular 11-year period. The most promising model of studying solar cycle at present is the flux transport dynamo model (Choudhuri et al. 1995; Durney 1995; Dikpati \\& Charbonneau 1999; Chatterjee et al. 2004). The main sources of irregularities in this model are the stochastic fluctuations in the Babcock--Leighton process of poloidal field generation (Choudhuri 1992; Choudhuri et al. 2007) and the stochastic fluctuations of meridional circulation (hereafter MC) (Hathaway 1996). Therefore we propose that the polar field or amplitude of MC or both decreased at the beginning of Maunder minimum. With this proposition, we use a flux transport dynamo model to reproduce a Maunder minimum. The details of this work can be found in Choudhuri \\& Karak (2009) and Karak (2010). ", "conclusions": "We have shown that most of the important features of the Maunder minimum can be reproduced quite well by assuming a simple ansatz that the polar field or the amplitude of MC or both decreased significantly at the beginning of Maunder minimum. Because of our lack of knowledge about the physical conditions at the beginning of the Maunder minimum, we cannot say how exactly the Sun was driven to the Maunder minimum. However, we should mention that there are several independent studies (Wang \\& Sheeley 2003; Miyahara et al. 2004; Passos \\& Lopeas 2009) suggesting that the amplitude of MC was weaker during the Maunder minimum. If this happens to be correct, then this study along with several other studies (Chatterjee et al. 2004; Chatterjee \\& Choudhuri 2006; Jiang et al. 2007; Yeates et al. 2008; Goel \\& Choudhuri 2009; Karak \\& Choudhuri 2011) indicate that the solar dynamo actually is diffusion-dominated and not advection-dominated." }, "1101/1101.5202_arXiv.txt": { "abstract": "Millisecond pulsars are intrinsically very stable clocks and precise measurement of their observed pulse periods can be used to study a wide variety of astrophysical phenomena. In particular, observations of a large sample of millisecond pulsars at regular intervals, constituting a Pulsar Timing Array (PTA), can be used as a detector of low-frequency gravitational waves and to establish a standard of time independent of terrestrial atomic timescales. Three major timing array projects have been established: The European Pulsar Timing Array (EPTA), the North American pulsar timing array (NANOGrav) and the Parkes Pulsar Timing Array (PPTA). Results from the PPTA project are described in some detail and future prospects for PTA projects are discussed. ", "introduction": "Pulsars and especially millisecond pulsars (MSPs) are remarkably stable celestial clocks. This great period stability opens up a wide range of potential applications. Pulsar timing analysis is based on the measurement of precise pulse times of arrival (ToAs) at the telescope. These ToAs are then transformed to the Solar-System barycentre which approximates an inertial frame. A model for the pulsar, including its position, proper motion, period and period derivatives and binary parameters (if appropriate) can be used to predict the pulse ToAs. The difference between the observed and predicted ToAs are known as {\\em timing residuals}. These timing residuals contain information about errors in the model parameters and unmodelled phenomena affecting the observed pulse period and so are at the heart of all pulsar timing analyses; see Hobbs et al. \\cite{hem06} and Edwards et al. \\cite{ehm06} for more details. A pulsar timing array (PTA) consists of an array of pulsars, widely distributed on the celestial sphere, that are being timed with high precision and at frequent intervals over a long data span \\cite{hd83,rom89,fb90}. Such a PTA has the potential to detect low-frequency gravitational waves propagating in the Galaxy \\cite{saz78,det79}, to improve our knowledge of Solar-System parameters \\cite{fb90} and to establish a pulsar-based standard of time that is independent of terrestrial atomic timescales \\cite{pt96}. Only MSPs have sufficiently narrow pulses (in time units) and sufficiently stable periodicities to be useful for PTA applications. ", "conclusions": "The realisation of Pulsar Timing Arrays is an exciting new development in pulsar astrophysics. PTA data sets have many applications including the detection of low-frequency gravitational waves and the establishment of a pulsar-based standard of time. Combining of existing and future data sets to form an International Pulsar Timing Array will give improved results in all applications. Looking further into the future, the greatly increased sensitivity provided by the proposed Square Kilometer Array will make possible detailed studies of phenomena presently near or below the level of significant detection. \\begin{theacknowledgments} I thank my colleagues in the PPTA project and the staff of the Parkes Observatory for their efforts which have been vital to the realisation of the PPTA. The Parkes radio telescope is part of the Australia Telescope which is funded by the Commonwealth Government for operation as a National Facility managed by CSIRO. \\end{theacknowledgments}" }, "1101/1101.2614_arXiv.txt": { "abstract": "We have used the images from the ACS on HST in \\ha{}, and in the neighboring continuum, to produce flux calibrated images of the large spiral galaxy M51, and the dwarf irregular NGC 4449. From these images we have derived the absolute luminosities in \\ha{}, the areas, and the positions with respect to the galactic centers as reference points, of over 2600 \\hii{} regions in M51 and over 270 \\hii{} regions in NGC 4449. Using this database we have derived luminosity (L)--volume (V) relations for the regions in the two galaxies, showing that within the error limits these obey the equation $L \\sim V^{2/3}$, which differs from the linear relation expected for regions of constant uniform electron density. We discuss briefly possible models which would give rise to this behavior, notably models with strong density inhomogeneities within the regions. Plotting the luminosity functions for the two galaxies we find a break in the slope for M51 at $\\log(L)$ = 38.5 dex (units in erg s$^{-1}$) for M51 in good agreement with the previous ground-based study by Rand, and above this luminosity NGC 4449 also shows a sharp decline in its luminosity function, although the number of regions is too small to plot the function well at higher luminosities. The cumulative diameter distribution for the \\hii{} regions of M51 shows dual behaviour, with a break at a radius close to 100 pc, the radius of regions with the break luminosity. Here too we indicate the possible physical implications. ", "introduction": "\\label{sec:intro} There is a well grounded tradition of studying populations of \\hii{} regions in spiral and irregular galaxies, which goes back to the work of Hodge, Kennicutt, and their collaborators \\citep{kennicutt80, hodge89, kennicutt89}. It has become customary to measure their \\ha{} luminosities and their radii, and prepare systematic catalogs. As techniques have improved and the precision of the measurements has increased the scope of the measurements has been gradually extended. The methods for discriminating an \\hii{} region from its diffuse surrounding emission have been improved, and the luminosity functions have been extended to lower limiting values. Although the ultimate aim should be to improve our understanding of the star formation process and the interaction of, notably massive, stars with the ISM, working on the improvement of the intrinsic data base is an objective intrinsically worth pursuing. Considerable worthwhile ground based data have been reported (a representative but in no way comprehensive list of references is: \\citet{rand92, knapen93, rozas96, rozas99, rozas00, gonzalez97, feinstein97, youngblood99, mackenty00, thilker02, buckalew06, hakobyan08}) and, in the case of M33, \\citet{cardwell00} presented a list of over 9,000 detected regions in a single galaxy, but it is clearly of interest to analyze data of the quality and angular resolution provided by HST for this purpose. However relatively little work has been presented of this type, because the required narrow band filters at arbitrary recession velocities are not included in the optics of the HST instruments. A notable previous study of M51 based on WFPC2 narrow band imaging was presented by \\citet{scoville01}, and a parallel study of the same galaxy based on NICMOS observations in Pa$\\alpha$ was published by \\citet{alonso01}. In this article we use ACS data for the first time to examine the statistical properties of \\hii{} regions in whole galaxies, presenting results for M51, and also for the dwarf irregular NGC 4449. In section 2 we discuss how the archival data were processed in order to prepare the \\ha{} maps, and in section 3 we describe the critically important processes of continuum and background subtraction, followed by the derivation of the luminosities, areas, and positions of the regions. In section 4 we derive the luminosity-radius relations and the luminosity functions for both galaxies and the diameter distribution for M51, and in section 5 we offer a discussion and our conclusions. ", "conclusions": "\\label{sec:conclusion} We have used observations of two low redshift galaxies, M51 and NGC 4449, with the ACS on HST to quantify some of the principal physical characteristics of their populations of \\hii{} regions. For both galaxies we have produced catalogs of the luminosities in \\ha{}, the projected areas, and center positions of the \\hii{} regions. For M51 the catalog contains 2659 regions, as the resolution and image quality of HST have allowed us to make a significant advance over previous ground based studies. Our catalog goes down 0.6 dex in luminosity compared to the detailed ground based study by \\citet{rand92}, and the number of regions measured is greater by a factor 4. For NGC 4449 we analyzed the emission from 273 regions, which is not in fact a striking quantitative advance on ground-based catalogs \\citep[see e.g.][]{sabbadin79} in this much less massive object with far fewer regions. From our databases we have derived the luminosity--volume (or luminosity--radius) relations for the regions of each galaxy, finding that in both cases the luminosity varies as the 2/3 power of the volume, (i.e. as the square of the radius) of the region. For a set of \\hii{} regions with constant uniform density the relation between luminosity and volume would have been linear. We can account for our result if the \\hii{} regions are globally in quasi pressure equilibrium with their surroundings, so that the mean electron density of a large and luminous region is less than that in a smaller less luminous region, because the mean temperature of the gas is larger in the former case. We have presented evidence for pressure equilibrium previously \\citep{letter2010}, as the measured electron density in the \\hii{} regions of our two measured galaxies varies with galactocentric distance with the same scale length as the neutral hydrogen, which would be the case for systems in pressure equilibrium. So there is a coherent qualitative explanation for the observed luminosity-volume relation, but a quantitative explanation for the simple relation found must await a more complete modelling exercise. We also present luminosity functions for the \\hii{} regions of both galaxies. Above the completeness limit, both functions show, to zero order approximation, a classical falling exponential. For M51 there is clear evidence for a break in the exponential, at a luminosity of $\\log(L)$ = 38.51 (units in erg s$^{-1}$) with a steeper fall to higher luminosities confirming the previous finding by \\citet{rand92}. The luminosity function for the interarm regions alone barely shows evidence for this break, but the number of highly luminous regions in this function is so small that it is not possible to state that the function is really distinct in form from the overall function for the galaxy, or from the function for the regions in the arms. For NGC 4449 there is also a steepening of the luminosity function above 38.5 dex, but here again the numbers of regions are too small to show whether this is a physical effect or merely a statistical effect. Explanations for the broken exponential, which has been reported previously in a number of papers, notably in \\citet{kennicutt89}, have included the blending of \\hii{} regions such that the largest regions absorb the smaller surrounding population \\citep{pleuss00} and the increasing escape of ionizing photons from the larger regions \\citep{beckman00}. The first hypothesis does yield a broken exponential but when developed in due detail predicts a shallower outer slope rather than the steeper slope observed. The evidence presented here that the more luminous a region the lower is its electron density does give some indirect support to the second hypothesis, but further work, notably incorporating the effects of dust extinction on the ionizing photon budget, will be needed to clarify the situation." }, "1101/1101.2108_arXiv.txt": { "abstract": "{I discuss three different topics in Galactic chemical evolution: the \"puzzling\" absence of any observational signature of secondary elements ; the building of the Galactic halo in the framework of hierarchical galaxy formation, as evidenced from its metallicity distribution ; and the potentially important role that radial migration may play in the evolution of galactic disks, according to recent studies. } \\FullConference{11th Symposium on Nuclei in the Cosmos, NIC XI\\\\ July 19-23, 2010\\\\ Heidelberg, Germany} \\begin{document} ", "introduction": " ", "conclusions": "" }, "1101/1101.5491.txt": { "abstract": "The cosmic dark ages ended a few hundred million years after the Big Bang, when the first stars began to fill the universe with new light. It has generally been argued that these stars formed in isolation and were extremely massive -- perhaps $100$ times as massive as the Sun. In a recent study, Clark and collaborators showed that this picture requires revision. They demonstrated that the accretion disks that build up around Population~III stars are strongly susceptible to fragmentation and that the first stars should therefore form in clusters rather than in isolation. We here use a series of high-resolution hydrodynamical simulations performed with the moving mesh code {\\small AREPO} to follow up on this proposal and to study the influence of environmental parameters on the level of fragmentation. We model the collapse of five independent minihalos from cosmological initial conditions, through the runaway condensation of their central gas clouds, to the formation of the first protostar, and beyond for a further $1000$~years. During this latter accretion phase, we represent the optically thick regions of protostars by sink particles. Gas accumulates rapidly in the circumstellar disk around the first protostar, fragmenting vigorously to produce a small group of protostars. After an initial burst, gravitational instability recurs periodically, forming additional protostars with masses ranging from $\\sim 0.1$ to $10\\,{\\rm M}_\\odot$. Although the shape, multiplicity, and normalization of the protostellar mass function depend on the details of the sink-particle algorithm, fragmentation into protostars with diverse masses occurs in all cases, confirming earlier reports of Population~III stars forming in clusters. Depending on the efficiency of later accretion and merging, Population~III stars may enter the main sequence in clusters and with much more diverse masses than are commonly assumed. ", "introduction": "Over the past decade, a consensus has emerged on how the first, so-called Population~III (Pop~III), stars may have formed out of the hydrogen and helium forged in the Big Bang \\citep{bl01,bl04a,bromm09}. Non-linear structure formation from an initially nearly featureless universe began when dark matter minihalos with masses of order $10^6\\,M_\\odot$ collapsed at redshifts $z\\sim 20-50$, confining gas within their gravitational potential wells \\citep{htl96,tegmark97}. Cooling through ro-vibrational transitions of newly formed molecular hydrogen then triggered runaway collapse at the center of the halo, resulting in the formation of a protostar with density more than twenty orders of magnitude higher than the cosmic mean \\citep{yoh08}. The envelope of the accreting protostar remained hot due to the lack of cooling by metals and dust, so that accretion rates were on average higher than in star-formation regions today \\citep{mk04,mo07,zy07}. With a few exceptions \\citep{tao09}, simulations of this initial collapse phase have shown no fragmentation \\citep{abn02,bcl02,bl04b,yoshida06b,on07,yoh08}, leading to the conclusion that the first stars formed in isolation and were extremely massive. \\begin{table*} \\begin{center} \\caption{Simulation parameters} \\begin{tabular}{ccccccccccc} \\hline \\hline Simulation & Size [kpc] & Particles & $M_{\\rm dm}~[{\\rm M}_\\odot]$ & $M_{\\rm dm, ref}~[{\\rm M}_\\odot]$ & $M_{\\rm gas}~[{\\rm M}_\\odot]$ & $\\sigma_{8}$ & $M_{\\rm vir}~[{\\rm M}_\\odot]$ & $r_{\\rm vir}~[{\\rm pc}]$ & $\\lambda$ & $z_{\\rm coll}$ \\\\ \\hline MH-1 & $1000$ & $512^3$ & $272$ & $3.53$ & $0.72$ & $0.81$ & $5.8\\times 10^{5}$ & $150$ & $0.059$ & $18.6$ \\\\ MH-2 & $500$ & $256^3$ & $272$ & $3.53$ & $0.72$ & $0.9$ & $3.0\\times 10^{5}$ & $110$ & $0.055$ & $19.5$ \\\\ MH-3 & $250$ & $128^3$ & $272$ & $3.53$ & $0.72$ & $1.2$ & $2.3\\times 10^{5}$ & $94$ & $0.073$ & $20.9$ \\\\ MH-4 & $500$ & $256^3$ & $272$ & $3.53$ & $0.72$ & $1.1$ & $3.1\\times 10^{5}$ & $97$ & $0.044$ & $22.6$ \\\\ MH-5 & $500$ \u00a0& $256^3$ & $272$ & $3.53$ & $0.72$ & $1.3$ & $1.8\\times 10^{5}$ & $58$ & $0.038 $ & $31.7$ \\\\ \\hline \\multicolumn{11}{l}{} \\\\ \\multicolumn{11}{l}{\\parbox{16.5cm}{The comoving box sizes, particle numbers, initial DM masses, refined DM masses, gas masses, and normalizations used in the simulations, as well as the viral masses, virial radii, spin parameters, and collapse redshifts of the first minihalos that form. The halo properties agree well with the results of previous studies \\citep[e.g.,][]{mba01,yoshida03a,gao07, on07}.}}\\\\ \\end{tabular} \\end{center} \\end{table*} In contrast, studies of present-day star formation have generally found fragmentation to occur shortly after the formation of the first protostar \\citep{kb00,bbb03,krumholz09,peters10}. Following up on this result, sink particles were recently used in studies of the fragmentation of primordial gas in minihalos \\citep{cgk08,sgb10,clark11a}. They found that the metal-free gas clouds fragment strongly, with the details of the process depending on the degree of turbulence in the halo. Focusing on the dynamical evolution of the high-density gas in the central regions of a minihalo, \\citet{clark11b} demonstrated that the protostellar disks around primordial stars accrete from the infalling envelope faster than they can transfer their mass onto the central object. As a result, they rapidly become gravitationally unstable and fragment to build-up binary or higher-order multiple stellar systems. These results challenge the idea that the first stars formed in isolation and give rise to a number of new questions: What is the mass spectrum of Pop~III stars in groups? Is it different from the previously preferred single mode of star formation? How does it depend on the dynamical characteristics of the host halo? How does it influence subsequent cosmic evolution? To address these issues, we here report on a series of hydrodynamical simulations performed with the quasi-Lagrangian moving mesh code {\\small AREPO} \\citep{springel10a}. Its high accuracy, flexibility, and efficiency allow us to follow the evolution of the gas in five statistically independent minihalos from fully cosmological initial conditions down to the optically thick regime where primordial protostars are born. We use a sink-particle approach that is similar to previously employed techniques to capture the subsequent accretion phase \\citep[e.g.,][]{bbp95,kmk04, jappsen05, federrath10}, which allows us to reach well beyond the formation of the first protostar, and circumvent the limitations posed in previous high-resolution {\\it ab initio} calculations of primordial star formation \\citep{yoh08}. The five realizations give us an indication of how common fragmentation is, and of the shape of the mass function of Pop~III protostars during the early stages of accretion. The structure of our work is as follows: In Section~2, we describe the numerical setup and physical ingredients of the simulations. In Section~3, we present the results of the simulations, followed by a discussion of radiation feedback, a resolution study, and the caveats of the sink particle treatment. Finally, in Section~4 we summarize our results and draw conclusions. All distances quoted in this paper are in proper units, unless noted otherwise. \\begin{figure*} \\begin{center} \\resizebox{12cm}{13.5cm} {\\unitlength1cm \\begin{picture}(12,13.5) \\put(0,7.5){\\includegraphics[width=6cm,height=6cm]{fig1a.eps}} \\put(6,7.5){\\includegraphics[width=6cm,height=6cm]{fig1b.eps}} \\put(0,1.5){\\includegraphics[width=6cm,height=6cm]{fig1c.eps}} \\put(6,1.5){\\includegraphics[width=6cm,height=6cm]{fig1d.eps}} \\put(0,0){\\includegraphics[width=12cm,height=1.5cm]{fig1e.eps}} \\end{picture}} \\caption{A test of the classical \\citet{truelove98} criterion in {\\scriptsize AREPO} for the \\citet{bb79} isothermal collapse problem. The individual panels compare simulations with on-the-fly mesh refinement with $1$, $2$, $4$, and $8$ cells per Jeans length, respectively. We show the density-squared weighted mass density projected along the line of sight in a box with $200\\,{\\rm AU}$ on a side, centered on one of the two main clumps. Similar to the result found in AMR codes, the amount of artificial fragmentation increases dramatically once the local Jeans length is resolved by less than four cells.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\resizebox{13.5cm}{16cm} {\\unitlength1cm \\begin{picture}(13.5,16) \\put(0,10.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig2a.eps}} \\put(4.5,10.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig2b.eps}} \\put(9,10.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig2c.eps}} \\put(0,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig2d.eps}} \\put(4.5,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig2e.eps}} \\put(9,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig2f.eps}} \\put(0,1.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig2g.eps}} \\put(4.5,1.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig2h.eps}} \\put(9,1.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig2i.eps}} \\put(0,0){\\includegraphics[width=4.5cm,height=1.5cm]{fig2j.eps}} \\put(4.5,0){\\includegraphics[width=4.5cm,height=1.5cm]{fig2k.eps}} \\put(9,0){\\includegraphics[width=4.5cm,height=1.5cm]{fig2l.eps}} \\end{picture}} \\caption{A sequential zoom-in on the gas in each of the five minihalos. The panels show the density-squared weighted number density of hydrogen nuclei projected along the line of sight. The gas virializes on a scale of $\\simeq 5\\,{\\rm kpc}$ (comoving), followed by the runaway collapse of the central $\\simeq 1\\,{\\rm pc}$, where the gas becomes self-gravitating and decouples from the dark matter. In the final stages of the collapse, a fully molecular core on a scale of a few hundred AU forms. The visible turbulence induced by the virialization of the dark matter halo survives down to the smallest scales and later influences the fragmentation of the gas.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\resizebox{13.5cm}{10.5cm} {\\unitlength1cm \\begin{picture}(13.5,10.5) \\put(0,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig3a.eps}} \\put(4.5,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig3b.eps}} \\put(9,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig3c.eps}} \\put(0,1.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig3d.eps}} \\put(4.5,1.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig3e.eps}} \\put(9,1.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig3f.eps}} \\put(0,0){\\includegraphics[width=4.5cm,height=1.5cm]{fig3g.eps}} \\put(4.5,0){\\includegraphics[width=4.5cm,height=1.5cm]{fig3h.eps}} \\put(9,0){\\includegraphics[width=4.5cm,height=1.5cm]{fig3i.eps}} \\end{picture}} \\caption{See Figure~2 for caption.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=12cm]{fig4.eps} \\caption{Temperature versus number density of hydrogen nuclei in the fully cosmological simulations. The mass in each bin is color-coded from blue (minimum) to yellow (maximum). The evolution of the gas is similar to the results of previous studies, with the exception that HD cooling becomes important in simulations MH-2 and MH-3 and leads to a prolonged cooling phase at $n_{\\rm H}\\ga 10^4\\,{\\rm cm}^{-3}$.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=12cm]{fig5.eps} \\caption{Same as Figure~4, but for the follow-up high resolution simulations. In simulations MH-2 and MH-3, the gas heats up more violently after the prolonged cooling phase due to HD cooling, which creates shocks that are visible as `fingers' of hot, underdense gas. At densities greater than $\\sim 10^{16}\\,{\\rm cm}^{-3}$, the gas becomes optically thick to cooling radiation and shortly thereafter forms a protostar.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=14cm]{fig6.eps} \\caption{The panels on the left show the spherically averaged number density of hydrogen nuclei and enclosed gas mass as a function of radius just before the formation of the first protostar. The panels on the right show the spherically averaged specific angular momentum, temperature, and velocity dispersion in units of the sound speed as a function of enclosed gas mass. The density, enclosed gas mass, and angular momentum profiles are very similar overall, while the thermal evolution of the gas displays some scatter due to the activation of HD cooling in simulations MH-2 and MH-3. In these two minihalos, the gas heats up later but more violently after becoming gravitationally unstable. The velocity dispersion shows no convincing correlation with the thermal history of the gas and is always close to Mach numbers $M\\simeq 1$, indicating transsonic turbulence. An interesting feature is the simultaneous collapse of a second gas cloud in simulation MH-2, which is visible as a density peak at about $1000\\,{\\rm AU}$ from the primary cloud. Similar behavior was found in another recent study \\citep{tao09}.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\resizebox{16cm}{21.5cm} {\\unitlength1cm \\begin{picture}(16,21.5) \\put(0,17.5){\\includegraphics[width=4cm,height=4cm]{fig7a.eps}} \\put(4,17.5){\\includegraphics[width=4cm,height=4cm]{fig7b.eps}} \\put(8,17.5){\\includegraphics[width=4cm,height=4cm]{fig7c.eps}} \\put(12,17.5){\\includegraphics[width=4cm,height=4cm]{fig7d.eps}} \\put(0,13.5){\\includegraphics[width=4cm,height=4cm]{fig7e.eps}} \\put(4,13.5){\\includegraphics[width=4cm,height=4cm]{fig7f.eps}} \\put(8,13.5){\\includegraphics[width=4cm,height=4cm]{fig7g.eps}} \\put(12,13.5){\\includegraphics[width=4cm,height=4cm]{fig7h.eps}} \\put(0,9.5){\\includegraphics[width=4cm,height=4cm]{fig7i.eps}} \\put(4,9.5){\\includegraphics[width=4cm,height=4cm]{fig7j.eps}} \\put(8,9.5){\\includegraphics[width=4cm,height=4cm]{fig7k.eps}} \\put(12,9.5){\\includegraphics[width=4cm,height=4cm]{fig7l.eps}} \\put(0,5.5){\\includegraphics[width=4cm,height=4cm]{fig7m.eps}} \\put(4,5.5){\\includegraphics[width=4cm,height=4cm]{fig7n.eps}} \\put(8,5.5){\\includegraphics[width=4cm,height=4cm]{fig7o.eps}} \\put(12,5.5){\\includegraphics[width=4cm,height=4cm]{fig7p.eps}} \\put(0,1.5){\\includegraphics[width=4cm,height=4cm]{fig7q.eps}} \\put(4,1.5){\\includegraphics[width=4cm,height=4cm]{fig7r.eps}} \\put(8,1.5){\\includegraphics[width=4cm,height=4cm]{fig7s.eps}} \\put(12,1.5){\\includegraphics[width=4cm,height=4cm]{fig7t.eps}} \\put(0,0){\\includegraphics[width=16cm,height=1.5cm]{fig7u.eps}} \\end{picture}} \\caption{The formation of a protostellar cluster at the center of the minihalos using standard sink particles. The panels show the density-squared weighted number density of hydrogen nuclei projected along the line of sight. Black dots and crosses denote protostars with masses below and above $1\\,{\\rm M}_\\odot$, respectively. The process of initial disk formation and fragmentation is remarkably similar in all minihalos \\citep[see also][]{clark11b}, after which N-body effects become important and lead to relatively unique configurations. For example, in simulation MH-4 dynamical interactions have led to the ejection of a low-mass protostar after only $\\simeq 50\\,{\\rm yr}$. This occurs significantly later in the other four minihalos.} \\end{center} \\end{figure*} % \\begin{figure*} % \\begin{center} % \\includegraphics[width=13.5cm]{fig8.eps} % \\caption{The central $2000\\,{\\rm AU}$ after $1000\\,{\\rm yr}$ of continued fragmentation and accretion. Black dots, crosses and stars denote protostars with masses below $1\\,{\\rm M}_\\odot$, between $1\\,{\\rm M}_\\odot$ and $3\\,{\\rm M}_\\odot$, and above $3\\,{\\rm M}_\\odot$. A relatively rich protostellar cluster with a range of masses has survived in each case. In a few minihalos, low-mass protostars have been ejected out of the central gas cloud, such that they are no longer visible here. In simulation MH-2, two independent clumps have collapsed almost simultaneously and formed their own clusters before eventually merging (see also Figure~6).} % \\end{center} % \\end{figure*} \\begin{figure*} \\begin{center} \\resizebox{13.5cm}{10.5cm} {\\unitlength1cm \\begin{picture}(13.5,10.5) \\put(0,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig8a.eps}} \\put(4.5,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig8b.eps}} \\put(9,6){\\includegraphics[width=4.5cm,height=4.5cm]{fig8c.eps}} \\put(2.25,1.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig8d.eps}} \\put(6.75,1.5){\\includegraphics[width=4.5cm,height=4.5cm]{fig8e.eps}} \\put(0,0){\\includegraphics[width=13.5cm,height=1.5cm]{fig8f.eps}} \\end{picture}} \\caption{The central $2000\\,{\\rm AU}$ after $1000\\,{\\rm yr}$ of continued fragmentation and accretion. Black dots, crosses and stars denote protostars with masses below $1\\,{\\rm M}_\\odot$, between $1\\,{\\rm M}_\\odot$ and $3\\,{\\rm M}_\\odot$, and above $3\\,{\\rm M}_\\odot$. A relatively rich protostellar cluster with a range of masses has survived in each case. In a few minihalos, low-mass protostars have been ejected out of the central gas cloud, such that they are no longer visible here. In simulation MH-2, two independent clumps have collapsed almost simultaneously and formed their own clusters before eventually merging (see also Figure~6).} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=14cm]{fig9.eps} \\caption{The mass accretion histories of the protostars (left panels), and the average accretion rates of the entire ensemble of protostars (right panels) using standard sink particles. In the left panels, each line denotes the evolution of an individual protostar. The thick dashed lines in both panels denote the cumulative values. Mergers are indicated by instantaneous jumps in mass. We find that star formation is not restricted to a single burst, but occurs continually for the entire simulated timespan. In every minihalo, between $5$ and $15$ protostars with masses ranging from $0.1$ to nearly $10\\,{\\rm M}_\\odot$ are formed. The total accretion rates are nearly constant over time at a few $0.01\\,{\\rm M}_{\\odot}\\,{\\rm yr}^{-1}$.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=14cm]{fig10.eps} \\caption{Same as Figure~9, but using adhesive sink particles. In this case fewer protostars with systematically higher masses are formed, although the total amount of gas within protostars is nearly identical.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=14cm]{fig11.eps} \\caption{The protostellar mass function after $1000$ years of continued fragmentation and accretion. The dark and light shadings distinguish the mass functions obtained for standard and adhesive sink particles, respectively. Despite very aggressive merging, a small cluster of protostars with a range of masses is formed even in the latter case. In the bottom right panel, we also show the cumulative mass functions obtained by summing up the contributions from the individual minihalos, and renormalized for better visibility. The resulting distribution is relatively flat between $\\sim 0.1$ and $\\sim 10\\,{\\rm M}_\\odot$, indicating that most of the mass is locked up in high-mass protostars.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\includegraphics[width=14cm]{fig12.eps} \\caption{The radial velocity and ratio of the radial velocity to the escape velocity of all protostars as a function of distance to the center of mass, shown for standard (left panels) and adhesive (right panels) sink particles. The critical ratio of unity, where protostars are assumed to escape from the central gas cloud, is denoted by the dashed line. The escape velocity is determined by using the total mass enclosed within the current distance of each protostar from the center of mass. Black dots, crosses and stars denote protostars with masses below $1\\,{\\rm M}_\\odot$, between $1\\,{\\rm M}_\\odot$ and $3\\,{\\rm M}_\\odot$, and above $3\\,{\\rm M}_\\odot$. In our standard implementation of sink particles, a number of low-mass protostars obtain high radial velocities and escape from the central gas cloud. They stop accreting after they receive substantial radial velocities during close encounters with other protostars. For adhesive sink particles, this occurs significantly less often. In this case no protostars reside within the central $\\simeq 100\\,{\\rm AU}$, since momentum conservation acts to move merged protostars to larger annuli.} \\end{center} \\end{figure*} \\begin{figure*} \\begin{center} \\resizebox{12cm}{7.5cm} {\\unitlength1cm \\begin{picture}(12,7.5) \\put(0.0,1.5){\\includegraphics[width=6cm,height=6cm]{fig13a.eps}} \\put(6.0,1.5){\\includegraphics[width=6cm,height=6cm]{fig13b.eps}} \\put(0.0,0.0){\\includegraphics[width=12cm,height=1.5cm]{fig13c.eps}} \\end{picture}} \\caption{A comparison of the fragmentation of the gas in simulation MH-1 with and without infrared radiation emitted by accreting protostars. The panels show the density-squared weighted temperature projected along the line of sight. The radiation slightly heats and puffs up the disk, such that the fragments form at somewhat larger distances from the center. The qualitative nature of radiation feedback is therefore very similar to what studies of present-day star formation have shown \\citep{krumholz09,peters10}. However, the effect is significantly reduced here due to the high temperature of the gas and its very efficient cooling by molecular hydrogen lines.} \\end{center} \\end{figure*} ", "conclusions": "We have used the moving mesh code {\\small AREPO} to follow the runaway collapse of the gas in five statistically independent minihalos from cosmological to protostellar scales -- over more than twenty orders of magnitude in density. We have captured the subsequent evolution of the newborn protostellar cloud for more than $100$ dynamical times with a sink-particle algorithm that resolves the accretion of the gas down to scales of $100\\,{\\rm R}_\\odot$ or less, comparable to the maximum photospheric size of Pop~III protostars. As proposed by \\citet{clark11b}, in all five minihalos a circumstellar disk forms that fragments vigorously into a small cluster of protostars with a range of masses. The gas becomes gravitationally unstable multiple times and continues to form protostars for the entire duration of the simulation. After only $1000\\,{\\rm yr}$ of accretion, of order $10$ protostars per minihalo have formed and display a relatively flat protostellar mass function ranging from $\\sim 0.1$ to nearly $10\\,{\\rm M}_\\odot$. Although the sink-particle technique employed in this study allows a continuation of the calculations to much later times than was possible in previous {\\it ab initio} simulations of primordial star formation \\citep{abn02,yoh08}, a number of uncertainties remain. One problem is the artificially reduced density and pressure created by the removal of mesh-generating points around sink particles. However, we have found that using systematically smaller accretion radii results in more fragmentation and a decrease of the typical fragment mass, such that a correct treatment of the boundary region between sink particles and the gas should strengthen our conclusions. A more important caveat is the inability of the sink-particle algorithm to model the gasdynamical friction between physically extended protostars during close encounters. In an attempt to maximize this effect, we have implemented adhesive sink particles in addition to a more standard formulation of sink particles. This leads to an increased merger rate, so that fewer protostars survive. However, even this extreme assumption does not prevent the formation of a relatively rich cluster of protostars. A final caveat is that we do not self-consistently model the interaction of the protostars with the surrounding gas. It is unclear how important the resulting neglect of gasdynamical friction and torques is, since simulations that model the protostellar surface as well as the parent gas cloud are not yet feasible. However, it appears unlikely that this caveat will qualitatively affect our conclusions, since fragmentation typically occurs on scales larger than the minimum resolution length. Modulo the uncertainties mentioned above, the simulations presented here portray a very different picture of primordial star formation than is commonly assumed. Instead of forming a single object, the gas in minihalos fragments vigorously into a number of protostars with a range of masses. It is an open question as to how this early mass function will be mapped into the final mass function of Pop~III, after accretion, fragmentation and merging have finally stopped. However, it is interesting to speculate how this nonlinear mapping will play out. If a flat, broad mass function persists, a number of lower-mass Pop III stars will have formed which could survive to the present day if their mass remains below $\\sim 0.8\\,M_\\odot$. Although this possibility is speculative because of the above uncertainties and the fact that we follow the protostellar accretion only for the first $1000$ out of $10^5$ or $10^6\\,{\\rm yr}$, it is worth looking for such Pop~III fossils in ongoing and planned large surveys of metal-poor stars in the Milky Way \\citep{bc05}. Specifically, such surveys should be focused on the Galactic bulge, where the Pop~III survivors should preferentially reside due to the biasing of the minihalo formation sites \\citep{dmm05,gao10}. The planned Apache Point Observatory Galactic Evolution Experiment (APOGEE) with its near-IR capability may be well suited for this search \\citep{majewski10}. Because of interstellar pollution, even true Pop~III fossils would appear as extreme Pop~II stars, but upper limits would be significantly lower than the currently probed values \\citep{fjb09}. Furthermore, the presence and mutual competition of multiple accretors in a given minihalo will act to limit the growth of the most massive objects. Pop~III stars are therefore less likely to reach masses in excess of $\\sim 140\\,{\\rm M}_\\odot$, the threshold for triggering extremely energetic pair-instability supernovae (PISNe) during stellar death \\citep{hw02}. A reduced PISN rate is more easily compatible with the absence of their distinct nucleosynthetic signatures in any of the extremely metal-poor halo stars observed so far \\citep{iwamoto05}. Pop~III stars could still have given rise to numerous extremely luminous supernova explosions, if they had masses of a few $10\\,{\\rm M}_\\odot$ and if they were rapid rotators, as suggested by recent studies \\citep{sbl11}. They would then have exploded as core-collapse hypernovae, with explosion energies that are similar to PISNe \\citep{un02}. Finally, our results may challenge models of so-called `dark stars', which are Pop~III stars powered by DM self-annihilation heating \\citep{freese08,iocco08}. These models invoke an increased DM interaction rate at the center of the Pop~III star which itself lies at rest at the center of its minihalo. On top of recent claims that the initial collapse is not substantially delayed by DM annihilations \\citep{ripamonti10}, the complex dynamics of a protostellar cluster at the center of the minihalo may further preclude efficient DM capture and heating." }, "1101/1101.4048_arXiv.txt": { "abstract": "Many efforts have been made to model the mass distribution and dynamical evolution of the circumnuclear gas in active galactic nuclei (AGNs). However, chemical evolution is not included in detail in three-dimensional (3-D) hydrodynamic simulations. The X-ray radiation from the AGN can drive the gas chemistry and affect the thermodynamics, as well as the excitation of the interstellar medium (ISM). Therefore, we estimate the effects (on chemical abundances and excitation) of X-ray irradiation by the AGN, for atomic and molecular gas in a 3-D hydrodynamic model of an AGN torus. We obtain the abundances of various species from an X-ray chemical model. A 3-D radiative transfer code estimates the level populations, which result in line intensity maps. Predictions for the CO $J=1\\rightarrow0$ to $J=9\\rightarrow8$ lines indicate that mid-$J$ CO lines are excellent probes of density and dynamics in the central ($\\lesssim60~\\rm pc$) region of the AGN, in contrast to the low-$J$ CO lines. Analysis of the $X_{\\rm CO}/\\alpha$ conversion factors shows that only the higher-$J$ CO lines can be used for gas mass determination in AGN tori. The \\cii\\ $158~\\mum$ emission traces mostly the hot ($T_k>1000~\\rm K$) central region of the AGN torus. The \\cii\\ $158~\\mum$ line will be useful for ALMA observations of high redshift ($z\\gtrsim1$) AGNs. The spatial scales ($\\ge0.25~\\rm pc$) probed with our simulations match the size of the structures that ALMA will resolve in nearby ($\\le45~\\rm Mpc$ at $0.01''$) galaxies. ", "introduction": "The formation and growth of a central black hole and its interaction with intense star-forming regions is one of the topics most debated in the context of galaxy evolution. There is observational evidence for a common physical process from which most active galactic nuclei (AGNs) and starbursts originate \\citep[e.g.,][]{soltan82, magorrian98, ferrarese00, graham01, haring04}. A plausible scenario considers that starbursts, super-massive black hole growth, and the formation of red elliptical and submillimeter galaxies, are connected through an evolutionary sequence caused by mergers between gas-rich galaxies \\citep{hopkins06, hopkins08, tacconi08, narayanan09, narayanan10}. In this scenario, the starbursts and (X-ray producing) AGNs seem to be co-eval, and the interaction processes between them (phase d and e in Figure~1 of \\citealt{hopkins08}), that dominate the formation and emission of molecular gas, is one of the long-standing issues concerning active galaxies. Numerous molecules tracing different (AGN and starburst driven) gas chemistry have been detected in Galactic and (active) extragalactic environments. Studies have shown that chemical differentiation observed within Galactic molecular clouds is also seen at larger ($\\sim$100 pc) scales in nearby galaxies \\citep[e.g.,][]{henkel87, nguyen91, martin03, usero04, tacconi08, pb07, pb09, pb10, baan10, vdwerf10}. The evolution of the ISM in the inner $100~\\rm pc$ region around a $10^8~M_{\\sun}$ supermassive black hole (SMBH) was investigated by \\citet{wada02} (hereafter \\citetalias{wada02}) using three-dimensional (3-D) Euler-grid hydrodynamic simulations. They took into account self-gravity of the gas, radiative cooling and heating due to supernovae (SNe). A clumpy and filamentary torus-like structure was found to be reproduced on a scale of tens of pc around the SMBH, with highly inhomogeneous ambient density and temperature, and turbulent velocity field. Their results indicated that AGNs could be obscured by the circumnuclear material. This represents theoretical support for observational evidence showing that some AGNs are obscured by nuclear starbursts \\citep[e.g.,][and references therein]{levenson01, levenson07, ballantyne08}. Several efforts have been made to estimate the molecular line emission from the nuclear region in these 3-D hydrodynamic simulations, and to compare the results with observational data. For instance, \\citet{wada05} (hereafter \\citetalias{wada05}) derived molecular line intensities emitted from the nuclear starburst region around a SMBH in an AGN. They used the 3-D hydrodynamic simulations (density, temperature, and velocity field data) of the multi-phase gas modeled by \\citetalias{wada02} as input for 3-D non-LTE radiative transfer calculations of \\twco\\ and \\thco\\ lines. They found that the CO-to-H$_2$ conversion factor ($X$-factor) is not uniformly distributed in the central 100 pc and the $X$-factor for \\twco\\ $J=1\\rightarrow0$ is not constant with density, in contrast with the \\twco\\ $J=3\\rightarrow2$ line. Similarly, the role of the \\hcn\\ and \\hcop\\ high-density tracers in the inhomogeneous molecular torus of \\citetalias{wada02} was studied by \\citet{yamada07} (hereafter \\citetalias{yamada07}). These non-LTE radiative transfer calculations suggested a complicated excitation state of the rotational lines of \\hcn\\ (with maser action) and \\hcop, regardless of the spatially uniform chemical abundance assumed. However, all these previous efforts to estimate the molecular line emissions from the central 100 pc of an AGN leave room for improvements. First of all, the radiative cooling in the simulations by \\citetalias{wada02} are not consistent with the chemical abundances in the cold and dense gas because (collisional) formation and (radiative) destruction of H$_2$ by far ultraviolet radiation (FUV) was not included. Therefore, the cold and dense gas in the simulations by \\citetalias{wada02} does not necessarily represent the dusty molecular gas phase around an AGN. Hence, in order to study the distribution and structures of the various density regimes of the H$_2$ gas, the 3-D hydrodynamic simulations of \\citetalias{wada02} were extended by \\citet*{wada09} (hereafter \\citetalias{wada09}) to solve the nonequilibrium chemistry of hydrogen molecules along with the hydrodynamics. The formation of H$_2$ on dust and its radiative destruction by far ultraviolet radiation (FUV) from massive stars are also included in the model by \\citetalias{wada09}. This allows to track the evolution of molecular hydrogen and its interplay with the \\hi\\ phase in the central $64 \\times 64 \\times 32~\\rm pc$ region. Thus, the radiative cooling in the model by \\citetalias{wada09} is more consistent with the chemical abundances expected in the cold ISM, in comparison with the models by \\citetalias{wada02}. Different SN rates and strengths of the uniform FUV field were also explored in order to study their effects on the structures of molecular gas. On the other hand, the inhomogeneous density and temperature structures observed in the 3-D hydrodynamic models are not the only factors that drive molecular abundances and excitation conditions of molecular lines. There is observational and theoretical evidence in the literature that supports different chemical evolution scenarios due to X-ray and UV radiation from the central AGN and circumnuclear starburst, as well as mechanical heating produced by turbulence and supernovae \\citep[e.g.,][]{kohno01, kohno07, kohno05, imanishi04, imanishi06a, imanishi06b, imanishi06, aalto07a, meijerink07, garcia07, loenen08, garcia08, pb09}. The strong UV and X-ray radiation from the AGN and accretion disk could affect both the dynamics and excitation of the molecular gas \\citep[e.g.,][]{ohsuga01a, ohsuga01b, meijerink07, vdwerf10}. However, the radiation field from the AGN itself was not taken into account in the earlier estimates of molecular line emissions from hydrodynamical simulations. A preliminary estimate of the potential effects of hard X-rays ($E>1$ keV) on the molecular gas was done by \\citetalias{wada09} using the X-ray Dissociated Region (XDR) models of \\citet{meijerink05}. It was found that XDR chemistry may change the distribution of H$_2$ around an AGN, if X-ray effects are explicitly included in the hydrodynamic model. The X-ray chemistry depends mainly on $H_X/n$, where $H_X$ is the X-ray energy deposition rate and $n$ is the number density of the gas \\citep{maloney96}. Although the H$_2$ abundance is robust in a clumpy medium like the one found in the hydrodynamical models, the temperature of the gas affected by an X-ray flux is expected to be a factor of $\\sim5$ higher than that found in standard models of a Photon Dominated Region (PDR; e.g., \\citealt{hollenbach99}), for $log(H_X)/n > \u221226$ \\citep{meijerink07}. This is because the ionization heating by X-rays is more efficient than photo-electric emission by dust grains. Other molecular and atomic lines have also been suggested as tracers of the AGN and starburst activity in nearby galaxies ($z<1$) as well as in high ($z\\ge1$) redshift galaxies. \\citet{spaans08} studied the possibility of using \\twco\\ and H$_2$ emission lines to trace a young population of accreting massive ($\\ge10^6~M_{\\sun}$) black holes at redshifts $z=5-20$ and radiating close to the Eddington limit. An enhancement in the intensities of various \\twco\\ transitions up the rotational ladder, as well as other molecular and atomic lines like \\thco, \\hcn, \\hcop, \\ci, \\cii, \\oi\\ and \\nii, is also expected to be observed when X-ray irradiation dominates the local gas chemistry \\citep{meijerink07, spaans08}. Simulations of quasars at $z\\sim6$ with massive ($10^{12}-10^{13}~M_{\\sun}$) halos and different merging histories showed that mid-$J$ \\twco\\ lines are highly excited by a starburst, while high velocity peaks are expected to be produced by AGN-driven winds \\citep{narayanan08a, narayanan08b}. It was further found by \\citet{narayanan09} that the compact \\twco\\ spatial extents, broad linewidths and high excitation conditions observed in Submillimetre Galaxies (SMGs) at $z\\sim2$ can be explained if SMGs are a transition phase of major merging events. In this work we use the XDR/PDR chemical model by \\citet{meijerink05} to estimate the abundances of more than 100 species (atoms and molecules) at each grid point in the computational box of the extended 3-D hydrodynamical models of an AGN torus by \\citetalias{wada09}. We also estimate the actual X-ray flux emerging from the AGN, derived from the central black hole mass. Flux attenuation by photo-absorption of X-rays along the ray path and the distance from the central black hole is included. Thus, we estimate non-homogeneous abundances at each grid point that depend on the local density and impinging X-ray flux. An extended version of the non-LTE 3-D radiative transfer code $\\beta$3D by \\citet{poelman05} is used to compute the level populations of any molecule or atom for which collision data exist in the LAMDA\\footnote{http://www.strw.leidenuniv.nl/$\\sim$moldata/} database \\citep{schoier05}. Molecular and atomic line intensities and profiles are calculated with a line tracing approach for an arbitrary viewing angle. Model predictions for future ALMA observations of CO lines and the \\cii\\ $158\\mum$ fine structure line are presented. The organization of this article is as follows. In Sec.~\\ref{sec:3Dmodel} we describe the numerical method. The results and analysis are presented in Sec.~\\ref{sec:results}. The final remarks are presented in Sec.~\\ref{sec:remarks}. ", "conclusions": "" }, "1101/1101.3342_arXiv.txt": { "abstract": "We use quasi-simultaneous near-infrared (near-IR) and optical spectroscopy from four observing runs to study the continuum around 1 $\\mu$m in 23 well-known broad-emission line active galactic nuclei (AGN). We show that, after correcting the optical spectra for host galaxy light, the AGN continuum around this wavelength can be approximated by the sum of mainly two emission components, a hot dust blackbody and an accretion disc. The accretion disc spectrum appears to dominate the flux at $\\sim 1$~$\\mu$m, which allows us to derive a relation for estimating AGN black hole masses based on the near-IR virial product. This result also means that a near-IR reverberation programme can determine the AGN state independent of simultaneous optical spectroscopy. On average we derive hot dust blackbody temperatures of $\\sim 1400$~K, a value close to the sublimation temperature of silicate dust grains, and relatively low hot dust covering factors of $\\sim 7\\%$. Our preliminary variability studies indicate that in most sources the hot dust emission responds to changes in the accretion disc flux with the expected time lag, however, a few sources show a behaviour that can be attributed to dust destruction. ", "introduction": "The broad emission line region (BELR) of active galactic nuclei (AGN) is one of the most direct tracers of the immediate environment of supermassive black holes. However, despite decades of intensive optical and ultraviolet (UV) spectrophotometric studies its geometry and kinematics remain ill-defined \\citep[see, e.g., review by][]{Sul00}. Our current, limited knowledge of its physical condition and scale was gained primarily through the application of photoionisation models \\citep[see, e.g., review by][]{Ferl03} and through reverberation mapping studies \\citep[see, e.g., review by][]{Pet93}. We have started to extend these studies to near-infrared (near-IR) wavelengths. In \\citet[][hereafter Paper I]{L08a} we outlined the rationale of our programme, presented the observations of the first three epochs and addressed briefly some of the important issues regarding the physics of the most prominent broad emission lines. Here we present the fourth epoch of observation and investigate the continuum around the 1 $\\mu$m inflection point. The AGN spectral continuum region around the rest-frame wavelength of $\\sim 1$ $\\mu$m is believed to sample simultaneously two important emission components, namely, the accretion disc \\citep[e.g.,][]{Mal82, Mal83} and the hottest part of the putative dusty torus \\citep[e.g.,][]{Barv87, Neu87}. However, although it is assumed to be understood, it has not yet been sampled spectroscopically in its entirety. By probing the long-wavelength end of the accretion disc spectrum \\citep{Kish05, Kish08} such an investigation has the potential to solve the discrepancy often found between theoretical models and observations \\citep[see, e.g., review by][]{Korat99}. Furthermore, since the accretion disc most likely illuminates directly the inner parts of the dust structure, monitoring the change in spectral flux and slope of these two emission components relative to each other can constrain the location and geometry of the obscurer. So far, only few sources have been observed in dust reverberation programmes \\citep[e.g.,][]{Glass92, Nel96, Okn99, Glass04, Min04, Sug06, Kosh09}. The paper is organised as follows. In Section 2 we briefly introduce the sample and discuss the observations. In Section 3 we derive pure AGN continuum spectral energy distributions (SEDs), based on which we constrain the individual continuum components (Section 4). The variability of these components is disussed in Section 5. Finally, in Section 6 we summarize our main results and present our conclusions. Throughout this paper we have assumed cosmological parameters $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_{\\Lambda}=0.7$. ", "conclusions": "We have used four epochs of quasi-simultaneous (within two months) near-IR and optical spectroscopy of 23 broad-line AGN to study the continuum spectral shape around 1 $\\mu$m. Our main results can be summarized as follows. \\vspace*{0.2cm} (i) The accretion disc spectrum appears to dominate the flux at $\\sim 1$~$\\mu$m, which allows us to derive a new relation that can be used to estimate AGN black hole masses. It is based on the near-IR virial product, defined here as the product between the width of the Pa$\\beta$ broad emission line and the integrated 1~$\\mu$m continuum luminosity. The dominance of the accretion disc spectrum at such long wavelengths means that the AGN state can be determined directly from the near-IR spectrum, making simultaneous optical spectroscopy for a reverberation programme unnecessary. (ii) An adequate subtraction of (in particular optical) host galaxy light reveals that the AGN continuum in the rest-frame frequency range of $\\nu \\sim 10^{14} - 10^{15}$ Hz can be approximated by the sum of mainly two emission components, a hot dust blackbody and an accretion disc spectrum. (iii) For the hot dust component we derive temperatures in the range of $T_{\\rm hot} \\sim 1100 - 1700$~K, which are typical values of the dust sublimation temperature, with a mean of $\\langle T_{\\rm hot} \\rangle = 1365\\pm18$ K. This mean value is close to the sublimation temperature of silicate dust grains, indicating that either carbonaceous dust is rare in AGN or that some other mechanism than dust sublimation sets the maximum dust temperature. The resulting hot dust covering factors are relatively low and in the range of $C \\sim 0.01 - 0.6$, with a mean of $\\langle C \\rangle = 0.07\\pm0.02$. (iv) Our preliminary variability studies have identified promising candidates for a future near-IR reverberation programme. Our three most variable sources in the near-IR are 3C~273, HE~1228$+$013 and NGC~4151. Furthermore, we have studied the response of the hot dust emission to changes in the accretion disc flux. Most sources show the expected time lag, but a few sources have a deficit of hot dust in the high state, which indicates dust destruction. \\vspace*{0.2cm} In our future work we will study the variability of the near-IR broad emission lines and constrain their physical conditions using detailed photoionisation models. In the longer term we plan to image our sample with current and future near-IR interferometers \\citep{Elvis02}." }, "1101/1101.2344_arXiv.txt": { "abstract": "We present results of Mt. Maidanak Observatory astroclimate study. Our data based on AZT--22 1.5m telescope observations in 1996--2005. ", "introduction": "\\label{sec:intro} Mount Maidanak Observatory at Ulugh Beg Astronomical Institute (UBAI) of Uzbek Academy of Sciences located in the South Uzbekistan, on the slopes of the Baisun range, at an altitude of 2600~m. {\\bf Brief history.} The Sternberg Astronomical Institute (SAI) of Moscow State University began first studies of the astronomical conditions at the most promising sites in the Central Asia in the end of the 1960s. One of the principal criteria used to select the telescope site was that the peak be isolated. In 1975, the SAI performed measurements of the astronomical climate at Mt. Maidanak with a two-beam instrument. Visual estimations showed the seeing $\\epsilon = 0.6\\arcsec$. Later SAI installed 3 telescopes (1.5m AZT--22 and two smaller instruments) on the observatory. Currently, all telescopes belong to the Mt. Maidanak Observatory of the UBAI. {\\bf 1.5m AZT--22 telescope} (Fig.~\\ref{fig:f1}) was designed by Leningrad Optics and Mechanics Amalgamation (LOMO). It was installed at Mt. Maidanak at the end of the 1980s, and made its first light detection in 1991. A two mirror quasi-Ritchey--Chretien system with a relative aperture of 1:8 is the principal optical system of the telescope. {\\bf Astroclimate.} The results of extensive studies of the observing conditions at Mt. Maidanak performed in the 1960s and 1970s were confirmed in the last decade with a series of observations of the astronomical climate using modern techniques. The four-year observations (1996-99) using the Differential Image Motion Monitor (DIMM) revealed a mean seeing of $0.69\\arcsec$ (Ehgamberdiev et al.\\cite{Ehgamberdiev00}). Later studies of the atmosphere at Mt. Maidanak with the MASS instrument (Kornilov et al.\\cite{Kornilov09}) detected a free--atmosphere seeing of $0.47\\arcsec$ at heights of 0.5 km and more above the level of Mt. Maidanak. Here we are presenting the study of the realistic seeing of CCD images obtained on the AZT--22 telescope. ", "conclusions": "" }, "1101/1101.2458_arXiv.txt": { "abstract": "The jet opening angle of gamma-ray bursts (GRBs) is an important parameter for determining the characteristics of the progenitor, and the information contained in the opening angle gives insight into the relativistic outflow and the total energy that is contained in the burst. Unfortunately, a confident inference of the jet opening angle usually requires broadband measurement of the afterglow of the GRB, from the X-ray down to the radio and from minutes to days after the prompt gamma-ray emission, which may be difficult to obtain. For this reason, very few of all detected GRBs have constrained jet angles. We present an alternative approach to derive jet opening angles from the prompt emission of the GRB, given that the GRB has a measurable $E_{peak}$ and fluence, and which does not require any afterglow measurements. We present the distribution of derived jet opening angles for the first two years of the Fermi Gamma-ray Burst Monitor (GBM) operation, and we compare a number of our derived opening angles to the reported opening angles using the traditional afterglow method. We derive the collimation-corrected gamma-ray energy, $E_{\\gamma}$, for GRBs with redshift and find that some of the GRBs in our sample are inconsistent with a proto-magnetar progenitor. Finally, we show that the use of the derived jet opening angles results in a tighter correlation between the rest-frame $E_{peak}$ and $E_{\\gamma}$ than has previously been presented, which places long GRBs and short GRBs onto one empirical power law. ", "introduction": "The Gamma-Ray Burst Monitor (GBM) onboard the Fermi Gamma-Ray Space Telescope has detected over 500 GRBs in its first 2 years of operation. A forthcoming catalog \\citep{GoldsteinGBM} contains time-integrated and time-resolved spectra for nearly all bursts during this time frame. With 12 sodium iodide (NaI) detectors and two bismuth germanate (BGO) detectors, GBM covers a wide energy band from 8 keV up to 40 MeV with roughly 2000 square centimeters of total detector surface area \\citep{Meegan}. This energy range effectively samples the prompt emission of GRBs and allows for rapid all-sky triggering and monitoring. To gather information about the GRB afterglow properties, redshift, and jet opening angle, other instruments are required. Since GBM can only localize a burst to 4 degree accuracy including systematic uncertainties \\citep{Briggs}, a simultaneous detection with \\Swift is usually required to derive a precise location for follow-up observations. \\Swift comprises a Burst Alert Telescope (BAT), an X-Ray Telescope (XRT) and an Ultraviolet-Optical Telescope (UVOT) \\citep{Barthelmy}. The \\Swift prompt energy coverage extends from 20 -- 150 keV, which does not allow for a comprehensive study of the higher energy prompt emission of GRBs which normally peaks at a few hundred keV. Therefore, it is obvious that the synergy between GBM and \\Swift results in a better understanding of this phenomenon. However, since there are relatively few GRBs that have been simultaneously detected by \\Swift and GBM, we are motivated to extend the ability to determine intrinsic properties of GRBs, such as jet opening angles and energetics, to the GBM observations of the prompt emission alone, using \\Swift afterglow studies for calibration. Current GRB theories assume that the explosions are collimated rather than isotropic, because otherwise the extreme energy outflow in gamma-rays during the prompt emission would eliminate many of the viable stellar mass progenitor models \\citep{Rhoads, Sari}. In fact, \\citet{Rhoads} and \\citet{Sari} proposed observable, achromatic `jet' breaks in the broadband afterglow light curves of GRBs to distinguish the jet collimation opening angle before such breaks were discovered \\citep{Granot}. It is now widely accepted that GRBs are collimated, yet few well-constrained jet angles have been unambiguously identified. Further, nearly all estimated jet opening angles have been for the long soft class of GRBs, while there are no constrained estimates for the short hard class \\citep{Kouveliotou}, most likely because short GRB afterglows are fainter and are thus less likely to be monitored long enough to detect a break \\citep{Gehrels, Kann}. In addition, precise locations are required and time must be requested of various observatories over a broad spectral range to study the late-time afterglow from minutes to days after the prompt emission. Several observational constraints and effects can hamper the identification of jet breaks such as gaps in temporal and spectral coverage and the presence of optical bumps in the light curve and X-ray flares. Even if a jet break is detected, an assumption of the density of the circumburst medium is required \\citep{Sari, Chevalier}, and so a certain amount of uncertainty is inherent in the calculation of the jet opening angle. Once determined, jet opening angles can lead to an estimate of the total energy release in gamma-rays. If the redshift and the prompt emission fluence of the GRB are also known, the collimation-corrected energy release at the source will provide the total energy budget of the GRB. These results constrain progenitor models and provide an estimate of the bulk Lorentz factor of the ejecta \\citep{Granot}. In addition, a more reliable and robust study of cosmology would be possible if a large number of collimation-corrected energies were known \\citep{Bloom}. Such studies are currently incomplete due to small statistics. For example, out of nearly 500 GRBs detected by GBM through July 2010, only 30 have observed redshifts from the simultaneous detection by \\Swift, and of these 30 GRBs only 8 have inferred jet opening angles from afterglow studies. The importance of GRBs to cosmology has recently become very clear with the detection of bursts out to a $z$ of 8.2 \\citep{090423}, making these events among the farthest detected in the observable Universe. Many authors have studied the spectral properties such as the peak luminosity; isotropic energy release in gamma-rays, $E_ {iso}$; the peak energy of the GRB power density spectrum, $E_{peak}$ \\citep{Mallozzi, Koshut, Lloyd, Amati, Bloom, Ghirlanda, Yonetoku, Firmani}. The purpose of this paper is to show that we can provide an alternative to the afterglow lightcurve monitoring method by deriving jet opening angles from the prompt emission of GRBs. We then utilize the opening angles to estimate the collimation-corrected energy release in gamma-rays, $E_{\\gamma}$, and show that there is a tight correlation between $E_{peak}$ and $E_{\\gamma}$. Section 2 describes the data sample and section 3 consists of our data analysis methods and results. Finally, we discuss the cosmological consequences and implications for progenitors of our study in section 4. ", "conclusions": "We have confirmed the $E_{peak}$/Fluence energy ratio results for GBM bursts, and have shown how they can be related to two different classes of GRBs. Most likely the larger of the two distributions belongs primarily to the long class of bursts typically associated with low-metallicity core-collapse supernovae, while the smaller mode belongs to the short class which is believed to be associated with neutron star-neutron star and neutron star-black hole mergers. The distribution of the energy ratio provides information for the determination of the rest-frame energetics, as can be seen in more detail when plotted in the $E_{peak}$--Fluence plane. The hard spectral cutoff in this plane is apparent for all BATSE and GBM bursts, despite the fact that these were detected with two different instruments with different sensitivities and different band passes. From this we can infer the cutoff is most likely not detector dependent, and using only physical observables, we can derive one of the most important parameters of GRBs, the jet opening angle, and given an observed redshift, calculate $E_ {\\gamma}$. Previously, the jet opening angle was only inferred for a handful of bursts for which X-ray, optical, and radio measurements of the afterglow were available \\citep{Sari}. Even if these measurements were available, very few bursts have discernible jet breaks that denote the moment at which the relativistic ejecta slow down to the point that the observed relativistic beaming angle is the same as the actual beaming angle of the outflow \\citep{Granot}. \\citet{Ghirlanda05} derived jet opening angles without afterglow measurements, although their derivation required the bursts to be described by the Ghirlanda relation and relied on the lag-luminosity relationship \\citep{Norris} to derive pseudo-redshifts, yielding an ensemble distribution of jet opening angles. From this, they postulated that bursts with softer $E_{peak}$ have larger opening angles than bursts with harder $E_{peak}$. By traditional classification of GRBs \\citep{Kouveliotou}, the interpretation is that short hard GRBs would have smaller opening angles than long soft GRBs. This is, in general, contradictory to our findings. To contrast, our results show that we can reproduce jet opening angles for individual GRBs that can be spectrally analyzed with prompt emission alone and does not require the estimation of pseudo-redshifts. Our resulting distribution of opening angles also agrees with a current theoretical model that the beamed outflow at the rotational poles of the progenitor is produced by the rotational angular momentum, as well as the configuration of the magnetic field \\citep{Livio}. In this model, the degree of collimation is related to the ratio of radius of the compact object to the radius of the accretion disk. In the case of a collapsar, the radius of the central object is much smaller than the radius of the disk, resulting in a tightly collimated beam, while a merger model results in much less collimation since the radius of the accretion disk is on the order of the radius of the central object. Due to the much larger opening angles, observations of jet breaks for most short hard GRBs as well as some from the long soft class are unlikely since the estimated jet break times will be on the order of several months \\citep{Sari}. Applying the jet opening angle to the redshift sample, we can estimate the entire energy outflow in gamma-rays. From our results, we can rule out a proto-magnetar progenitor for a few long GRBs (080916C, 080810, 090323, 090519, and 090902B) as well as a short GRB (090510). When we compare this energy budget for each GRB to the $E_{peak}$, we find a clear correlation between the energy of peak power production and the total energy output in gamma-rays. One major distinction of our relation shows that both classes of GRBs can be described by the same power law fit, which has not been previously shown. The uncertainties of our calculations include geometric effects, as GRBs are likely seen slightly off- axis from the center of the jet, and this error propagates through to the calculation of $E_\\gamma$. These uncertainties will likely be small for long bursts with small opening angles, due to the small possible displacement of the viewing angle relative to the jet opening angle. This relationship may potentially be exploited to extrapolate a rough redshift distribution for GRBs without redshift estimates, although a larger sample of GRBs with measured redshift is desired to confirm the existence of such a relation between $E_{peak}$ and $E_{\\gamma}$." }, "1101/1101.0207_arXiv.txt": { "abstract": "The radiative neutron capture on lithium-7 is calculated model independently using a low energy halo effective field theory. The cross section is expressed in terms of scattering parameters directly related to the $S$-matrix element. The cross section depends on the poorly known $p$-wave effective range parameter $r_1$. This constitutes the leading order uncertainty in traditional model calculations. It is explicitly demonstrated by comparing with potential model calculations. A single parameter fit describes the low energy data extremely well and yields $r_1\\approx -1.47$ fm$^{-1}$. ", "introduction": "} Low energy nuclear reactions play a crucial role in Big Bang Nucleosynthesis (BBN), stellar burning and element synthesis at supernova sites~\\cite{Burles:1999zt,Rolfs:1988,Barwick:2004ep}. Besides placing constraints on our understanding of element formation, these low energy reactions play an important role in testing astrophysical models and physics beyond the Standard Model of particle physics. Often the key nuclear reactions occur at energies that are not directly accessible in terrestrial laboratories. Radiative proton capture on beryllium \\pBe~ is one of them ---it is important for boron-8 production in the sun, whose weak decay results in the high energy neutrinos that are detected at terrestrial laboratories looking for physics beyond the Standard Model. The relevant solar energy, the Gamow peak, for this reaction is around $20$ keV~\\cite{Adelberger:1998qm}. This necessitates extrapolation to solar energies of known experimental capture cross sections from above around $100$ keV. Theoretical input becomes necessary for this extrapolation. Effective field theory (EFT) is an ideal formalism for this as it provides a model-independent calculation with reliable error estimates. In an EFT, one identifies the relevant low energy degrees of freedom and constructs the most general interactions allowed by symmetry without modeling the short distance physics. The interactions are organized in a low momentum expansion. At a given order in the expansion, a finite number of interactions has to be considered and an {\\em a priori} estimate of the theoretical error can be made. Establishing theoretical errors is crucial due to astrophysical demands~\\cite{Burles:1999zt,Rolfs:1988,Adelberger:1998qm}. A systematic expansion of interactions is important because many processes involve external currents, and any prescription used in phenomenological models involve some uncertainty. As an example, the cross section for $n(p,\\gamma)d$ at BBN energies was calculated within EFT to an accuracy of about $1\\%$~\\cite{Rupak:1999rk}. Systematic treatment of two-body currents was necessary to achieve this level of precision, and it addressed a critical need~\\cite{Burles:1999zt} for nuclear theory input in astrophysics. While applications of EFT to systems with $A\\lesssim 4$ nucleons is well developed, for $A\\gtrsim 5$ it is still in its infancy. However, some loosely bound systems, like halo nuclei open new possibilities. The small separation energy of the valence nucleons in halo nuclei provides a small expansion parameter for constructing a halo EFT~\\cite{Bertulani:2002sz,*Bedaque:2003wa}. The $^8$B nucleus with a proton weakly bound to the $^7$Be core by $0.1375$ MeV is a halo system. Current extrapolation of the \\pBe~ cross section to solar energies introduce errors in the $5-20\\%$ range~\\cite{Adelberger:1998qm,PhysRevC.68.045802,PhysRevC.70.065802}. A model-independent EFT calculation would be very useful to estimate the errors in the extrapolation. In addition, this would be an important step in developing EFT techniques for weakly-bound nuclei as has been accomplished in the few nucleon systems. Experiments such as those planned at the future FRIB~\\cite{FRIB} would explore exotic nuclei near the drip lines where halo systems abound. Structures and reactions with halo EFT can serve as benchmark for phenomenological models of nuclei near the drip lines. In this paper we consider the low energy reaction \\nLi, which is a isospin mirror to \\pBe. The $n$-$^7{\\rm Li}$ system allows formulating the EFT for the nuclear interactions without the added complication of the Coulomb force. Traditionally \\nLi~ has been calculated in a single-particle approximation as a $^7$Li core plus a valence neutron interacting via a Woods-Saxon potential~\\cite{Tombrello:1965}. This approximation breaks down at higher energies when the internal structure of the $^7$Li core is probed, for example, near the threshold for $^7$Li$(\\gamma, ^3$He$)\\alpha$ which is about $0.5$ MeV above the binding energy $B\\approx 2.03$ MeV of the $^8$Li core. We treat the $^7$Li nucleus as point-like since we work at very low energies. Once the nuclear piece is calculated in EFT for the $n$-$^{7}{\\rm Li}$ system the Coulomb interaction in $p$-$^{7}{\\rm Be}$ can be incorporated systematically as have been done for proton fusion in EFT~\\cite{Kong:1999tw,*Kong:1999sf,*Kong:1999mp,*Kong:1998sx,*Kong:2000px}. The reaction \\nLi, besides being a check on the mirror \\pBe~reaction, is important in inhomogeneous BBN. It impacts the production of carbon-oxygen-nitrogen in the early universe, and constrains astrophysical models~\\cite{kawano:1991ApJ372}. We calculate the \\nLi~ reaction analytically and express the result in terms of parameters directly related to observables, thus quantifying the dominant theoretical uncertainty in the single particle approximation. ", "conclusions": "} We considered radiative capture reactions for halo nuclei. The low energy \\nLi~ cross section was calculated at leading order using EFT. In the single particle approximation, the cross section was derived in terms of scattering parameters that are directly related to $S$-matrix elements. Using a model-independent formalism we demonstrated and quantified the theoretical uncertainty associated with phenomenological potentials in the single particle approximation. The leading order result depends on the $p$-wave effective range parameter $r_1$ that is poorly known. Without detailed knowledge about this parameter, model calculations deviate from data at low energy. We extract the effective range $r_1$ by fitting our analytic form to data. At higher order in the EFT expansion, the cross section would get corrections from two sources: higher order initial and final state interactions, and two-body currents. The initial and final state interactions can be related to the ERE. At the very low energy, it is the final state interactions, which modify the wave function renormalization constants, that are important. At next-to-next-to-leading order the shape parameter associated with $p$-wave interaction contribute~\\cite{Bertulani:2002sz,*Bedaque:2003wa, HigaRupak}. In addition, at higher order two-body currents such as $E_i (N F_j C)^\\dagger [NF_x( \\stackrel{\\rightarrow}{\\nabla}/M_C - \\stackrel{\\leftarrow}{\\nabla}/M_N)_y C] R_{ijxy}$, where $E_i$ is the electric field, contribute. These operators are not constrained by elastic scattering. A higher order EFT calculation would reduce theoretical errors though at the expense of additional parameters. This is not necessarily a drawback as what we gain is a model-independent understanding of the sources of higher order contributions, and a more detailed knowledge about the kind of experimental input that is required to better constrain the low energy theory. Coulomb interactions in $p+{}^7$Be scattering and \\pBe~reaction is being considered where the current formulation plays a crucial role~\\cite{HigaRupak}. The power counting of electromagnetic currents beyond leading order is being considered as well." }, "1101/1101.0814_arXiv.txt": { "abstract": "In this paper we show how a self-consistent treatment of hydrogen and helium emission line fluxes of the hosts of long gamma-ray bursts can result in improved understanding of the dust properties in these galaxies. In particular, we find that even with modest signal to noise spectroscopy we can differentiate different values for $R_V$, the ratio of total to selective extinction. The inclusion of Paschen and Brackett lines, even at low signal to noise, greatly increase the accuracy of the derived reddening. This method is often associated with strong systematic errors, caused by the need for multiple instruments to cover the wide wavelength range, the requirement to separate stellar hydrogen absorption from the nebular emission, and because of the dependancy of the predicted line fluxes on the electron temperature. We show how these three systematic errors can be negated, by using suitable instrumentation (in particular X-shooter on the Very Large Telescope) and wide wavelength coverage. We demonstrate this method using an extensive optical and near-infrared spectroscopic campaign of the host galaxy of gamma-ray burst 060218 (SN 2006aj), obtained with FORS1, UVES and ISAAC on the VLT, covering a broad wavelength range with both high and low spectral resolution. We contrast our findings of this source with X-shooter data of a star forming region in the host of GRB 100316D, and show the improvement over existing published fluxes of long GRB hosts. ", "introduction": "} As long gamma-ray bursts (hereafter referred to simply as GRBs) are produced by massive stars, as evidenced by their associated supernovae (e.g. Galama et al. 1998; Hjorth et al 2003; Pian et al.~2006; Starling et al. 2010; but see also Fynbo et al.~2006), there is a clear link between high mass star formation and GRBs. This link is apparent in the properties of their host galaxies: the hosts show bright nebular emission lines associated with active star formation, and show spectral energy distributions and morphologies generally consistent with blue, metal-poor, star forming dwarf galaxies (e.g. Savaglio, Glazebrook \\& Le Borgne 2009). From high resolution imaging it is evident that the positions of long GRBs in their hosts are strongly correlated with the location of starformation within these hosts (e.g. Fruchter et al 2006). It is therefore a reasonable assumption to make that the properties of the star forming regions in the host also lead to insight in the properties of the progenitor star: the progenitor only lives for a short time. A constraining parameter on the evolution of the progenitor object is its metallicity, or more specifically the abundance of iron, as it is the main driver of stellar winds. Measuring the gas phase iron abundance from emission lines of galaxies is complicated, as these emission lines are faint and iron is easily and rapidly depleted into dust grains. Oxygen is a good alternative: its abundance can be easily measured (the lines are bright) either directly (using oxygen lines at various ionisation stages in combination with electron temperature and density sensitive lines); or indirectly through the strong forbidden lines coupled with radiation transport simulations. In order to derive accurate abundances, one needs to correct for the extinction experienced by the hot gas that emits the different emission lines. Besides deriving element abundances from emission lines, there are further reasons why we are interested in the extinction in the star forming region(s) in the host galaxies. Particularly interesting is the combination of afterglow spectroscopy and photometry, which probes the line of sight dust and gas properties within the host, with host galaxy spectroscopy in emission. This combination may constrain the possible destruction of dust by the GRB, it may show how special the region in which the GRB took place is with respect to other star forming regions, and may shed light on the (depletion) chemistry in the ISM in GRB host galaxies. In this paper we focus on the extinction in star forming regions within GRB hosts (perhaps better described as the attenuation of emission lines by dust within the hosts). The host galaxy of GRB\\,060218 is a particularly attractive target to do this. The supernova accompanying this GRB is well studied, and therefore a connection with host galaxy properties (spectroscopic, spectral energy distribution; see e.g. Savaglio, Glazebrook \\& Le Borgne 2009) may provide tight constraints on the progenitor evolution. In a previous paper (Wiersema et al. 2007; hereafter W07) we performed a study of the metal abundances in this host and the ISM velocity structure, both obtained through emission and absorption lines detected in a VLT UVES spectrum of the supernova (2006aj). We obtained three more datasets of this source, which we will use to probe the extinction properties of this host in this paper. These data are part of a dedicated VLT program (program 381.D-0723, PI Wiersema) aimed at a thorough understanding of the stellar population(s) of this host galaxy, and consist of optical spectroscopy at medium and high resolution (using the FORS1 and UVES instruments) and low resolution K band spectroscopy (using the ISAAC instrument). This paper is organized as follows: in Section 2 we describe the way we use H and He\\,I line flux ratios to fit the extinction properties; in Section 3 we describe the observations and the way we measure emission line fluxes from the spectra; in Section 4 we fit the data using our models and in Section 5 we discuss our results. Throughout the paper we adopt a cosmology with $H_0 = 71$ km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m = 0.27$, $\\Omega_\\Lambda = 0.73$. ", "conclusions": "\\label{sec:conclusions} We have presented a spectroscopic dataset of the host of GRB\\,060218, which, together with already published data for a star forming tegion in the host of GRB\\,100316D, we use to get further insight into the attenuation of the emission lines by dust. Specifically we show how hydrogen and helium fluxes of many recombination transitions can be fit simultaneously to get insight into the electron temperature and the $R_V$ of the dust. X-shooter is the ideal instrument to do this for a large sample of GRB hosts, as long as accurate calibration can be achieved. The resulting dust properties from recombination lines can be combined with information from absorption lines of afterglow spectra; extinction curves from afterglow spectral energy distributions and reddening fits on integrated host galaxy spectral energy distributions to study dust formation and destruction in low metallicity dwarf galaxies." }, "1101/1101.2020_arXiv.txt": { "abstract": "We present the results of the ``Cosmogrid'' cosmological $N$-body simulation suites based on the concordance LCDM model. The Cosmogrid simulation was performed in a 30Mpc box with $2048^3$ particles. The mass of each particle is $1.28 \\times 10^5 M_{\\odot}$, which is sufficient to resolve ultra-faint dwarfs. We found that the halo mass function shows good agreement with the \\citet{Sheth1999} fitting function down to $\\sim 10^7 M_{\\odot}$. We have analyzed the spherically averaged density profiles of the three most massive halos which are of galaxy group size and contain at least 170 million particles. The slopes of these density profiles become shallower than $-1$ at the inner most radius. We also find a clear correlation of halo concentration with mass. The mass dependence of the concentration parameter cannot be expressed by a single power law, however a simple model based on the Press-Schechter theory gives reasonable agreement with this dependence. The spin parameter does not show a correlation with the halo mass. The probability distribution functions for both concentration and spin are well fitted by the log-normal distribution for halos with the masses larger than $\\sim 10^8 M_{\\odot}$. ", "introduction": "According to the present standard LCDM model, the Universe is thought to be composed primarily of cold dark matter (CDM) and dark energy \\citep{White1978, Peacock1999}. Structure formation of the Universe proceeds hierarchically in this model. Smaller-scale structures collapse first, and then merge into larger-scale structures. There is serious discrepancy between the distribution of subhalos in galaxy-sized halos obtained by numerical simulations and the observed number of dwarf galaxies in the Local Group \\citep{Klypin1999, Moore1999a}. This ``missing dwarf problem'' is still considered to be one of the most serious problem in the CDM paradigm \\citep[e.g.][]{Kroupa2010}. In order to understand the origin of this discrepancy, it is necessary to perform high-resolution cosmological $N$-body simulations and obtain unbiased sample of galaxy-sized halos with resolution high enough to obtain reliable statistics of subhalos since the subhalo abundance shows large halo-to-halo variations \\citep{Ishiyama2009}. Cosmological $N$-body simulations have been widely used to study the nonlinear structure formation of the Universe and have been an important tool for a better understanding of our Universe. In order to study the spatial correlation of galaxies, the first cosmological $N$-body simulations were performed in the 1970s using approximately 1000 particles \\citep[e.g.][]{Miyoshi1975, Fall1978, Aarseth1979, Efstathiou1979}. Since then, the development of better simulation algorithms and improvements in the performance of computers allow us to use much larger numbers of particles and have drastically increased the resolution of cosmological simulations. Today, it is not uncommon that the number of particles exceeds $10^9$ in high resolution simulations. In these works, the size of the simulation volumes is typically $[O(\\rm Gpc)]^3$ and populations of galaxy clusters, gravitational lensing, and the baryon acoustic oscillation are studied \\citep[e.g.][]{Evrard2002, Wambsganss2004, Teyssier2009, Kim2009, Crocce2010}. The simulation results are also used to construct mock halo catalogues for next generation large volume surveys. Others use simulations of $[O(\\rm 100Mpc)]^3$ volumes to study the internal properties of galaxy-sized dark matter halos, their formation, evolution, and statistical properties \\citep[e.g.][]{Springel2005, Klypin2010, White2010}. Using the results of high-resolution simulations of small-scale structures, we can study the fine structures of galactic halos, the distribution of subhalos, their structures, and their dependence on the nature of dark matter. This information has a strong impact on the indirect search for dark matter since gamma-ray flux by self-annihilation is proportional to local density if we consider neutralino as the candidate of dark matter. Thus, we can restrict the nature of dark matter using the results of high-resolution simulations of small-scale structures and indirect searches of dark matter. In addition, galaxies are considered to form in dark matter halos with a mass larger than a critical value \\citep{Strigari2008, Li2009, Maccio2009, Okamoto2009}. The structures of smallest halos which can host galaxies is important for the understanding of the galaxy formation processes. The simulation of smaller-scale structures of dark matter halos is not a trivial task since a very wide dynamic ranges of space, mass, and time must be covered. In particular, the number of time steps of such simulations is significantly larger than that of larger-scale simulations since the dynamical time scale is proportional to $1.0/\\sqrt{G\\bar{\\rho}}$, where $\\bar{\\rho}$ is the local density. Structures of smaller-scales form earlier, and thus have higher densities, therefore, simulations of smaller scales are computationally more expensive. Recently, simulations with galactic halos of very high resolution have been performed \\citep{Diemand2008, Springel2008, Stadel2009}. These works used the re-simulation method, where one selects one or a few halos at $z=0$ from a simulation which covers a large volume [typically a cube of size O(100Mpc)] with a relatively low-resolution. The corresponding regions of these halos are then identified in the initial particle distribution, and the particles in these regions are replaced by a larger number of smaller particles. After this is done, the entire volume is simulated to $z=0$ again. With this re-simulation method, we can resolve the structures of selected halos with extremely high resolution \\citep{Diemand2008, Springel2008, Stadel2009}. However, this method cannot be used for the study of halo-to-halo variations. Different halos are born in different environments and grow differently. The difference in the environment and growth history must be the cause of halo-to-halo variations. Therefore, in order to study variations, we need a bias-free set of a large number of halos. Clearly one cannot obtain a large number of halos with re-simulation method in practical time. In principle, one can improve the statistics by increasing the number of halos selected for re-simulations. In order to avoid the selection bias, we need to apply random, bias-free selection, and the most reliable bias-free selection is to select all halos, in other words, to simulate the entire simulation box with uniformly high mass resolution. \\citet{Ishiyama2009} performed the first bias-free high resolution simulation of small-scale structures. They analyzed the statistics of the subhalo abundance using the complete set of halos in the simulation box. The number of particles was $1600^3$ in a 46.5Mpc cubic box and the mass of a particle was $10^6 M_{\\odot}$. The subhalo abundance showed large halo-to-halo variations [see also \\citet{Ishiyama2008, Boylan2010}]. The concentration parameter and the radius at the moment of the maximum expansion showed fairly tight correlation with the subhalo abundance. Halos formed earlier have smaller number of subhalos at present. This correlation suggests that the difference in the formation history is the origin of the variation of the subhalo abundance [see also \\citet{Gao2004, Bosch2005, Zentner2005}]. The Millennium-II simulation \\citep{Boylan2009} used a 137Mpc cubic box and the particle mass of $\\sim 9.45 \\times 10^6 M_{\\odot}$. Its result is suitable for the analysis of the statistics of galaxy-sized dark matter halos, because the number of halos is larger than that of \\citet{Ishiyama2009}. However, due to the lack of the mass resolution, it cannot be used to study the statistics of dwarf-galaxy-sized halos and the statistics of subhalos with the size larger than faint dwarf galaxy. In this paper, we describe the first result of our Cosmogrid simulation. We simulated the evolution of halos in a 30Mpc cubic box using $2048^3$ particles. The mass of one particle is $1.28 \\times 10^5 M_{\\odot}$. The resolution reaches down to ultra-faint dwarf-galaxy-sized halos ($\\sim 10^7 M_{\\odot}$) and is more than 8 times better than that of our previous simulation \\citep{Ishiyama2009}. We focus on the halo mass function with the mass down to $10^7 M_{\\odot}$, the structures of most massive halos, and statistics of the internal properties of dwarf-galaxy-sized halos. We describe our initial conditions and numerical settings in Section \\ref{sec:method}, and results in Section \\ref{sec:result}. We discuss and summarize our results in Section \\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} We present the first scientific results of the Cosmogrid simulation. Because of unprecedentedly high resolution and powerful statistics, the simulation is suitable to resolve internal properties of halos with the mass larger than dwarf galaxy and subhalos whose scales are comparable to ultra-faint dwarf galaxies. We summarize the main results of this paper as follows: \\begin{itemize} \\item The halo mass function is well described by the \\citet{Sheth1999} fitting function down to $\\sim 10^7 M_{\\odot}$ from $1.0 \\times 10^{13} M_{\\odot}$. The differences are less than 10\\% at $z=0$ from $M=5.0 \\times 10^7 M_{\\odot}$ to $M=2.0 \\times 10^{12} M_{\\odot}$. \\item We analyzed the spherically averaged density profiles of the three most massive halos which contain more than 170 million particles. Their mass are 5.24, 3.58, and 2.25 $\\times 10^{13} M_{\\odot}$. We confirmed that the slopes of density profiles of these halos become shallower than $-1$ at the inner most radius. The results are consistent with the recent studies based on high resolution simulations for galactic halos. \\item We studied internal properties of halos at $z=0$ with the mass more than $\\sim 10^8 M_{\\odot}$. The concentration parameter measured by the maximum rotational velocity radius is weakly correlated with the halo mass. We found that the dependence of the concentration parameter with halo mass cannot be expressed by a single power law, but levels off at small mass. The slope of the mass-concentration relation is around $-0.07$ for halos with the mass $10^{10} M_{\\odot}$, and $-0.06$ for halos with the mass $10^{9} M_{\\odot}$. The shallowing slope naturally emerges from the nature of the power spectrum of initial density fluctuations. A simple model based on the Press-Schechter theory gives reasonable agreement with the simulation result. The spin parameter does not show a correlation with the halo mass. The probability distribution functions of concentration and spin are well fitted by the log-normal distribution for halos with the mass larger than $\\sim 10^8 M_{\\odot}$. \\end{itemize} We have shown here a first analysis of the Cosmogrid data and we plan to extend our analysis in future publications. Some of the topics that we want to address are: the variation of density profiles and its impact on the dark matter detectability, the statistics of subhalo abundance of the mass scale down to ultra faint dwarf in dwarf- and galaxy-sized halos, the assembly histories of halos, and the evolution of internal properties of halos which are presented in this paper at only $z=0$." }, "1101/1101.1679_arXiv.txt": { "abstract": "A non-parametric smoothing method is presented that reduces noise in multi-wavelength imaging data sets. Using Principle Component Analysis (hereafter PCA) to associate pixels according to their $ugriz$-band colors, smoothing is done over pixels with a similar location in PCA space. This method smoothes over pixels with similar color, which reduces the amount of mixing of different colors within the smoothing region. The method is tested using a mock galaxy with signal-to-noise levels and color characteristics of SDSS data. When comparing this method to smoothing methods using a fixed radial profile or an adaptive radial profile, the $\\chi^2$-like statistic for the method presented here is smaller. The method shows a small dependence on input parameters. Running this method on SDSS data and fitting theoretical stellar population models to the smoothed data of the mock galaxy and SDSS data, shows that the method reduces scatter in the best-fit stellar population analysis parameters, when compared to cases where no smoothing is done. For an area centered on the star forming region of the mock galaxy, the median and standard deviation of the PCA-smoothed data is 7 Myr ($\\pm$ 3 Myr), as compared to 10 Myr ($\\pm$ 1 Myr) for a simple radial average, where the noise-free true value is 7.5 Myr ($\\pm$ 3.7 Myr). ", "introduction": "Galaxy formation theories predict that baryons cool in dark matter halos in such a way as to provide connections between galaxy observables and dark matter properties \\citep{whi78,col89}. For example, the Tully-Fisher relation \\citep{tul77} shows the connection between galaxy luminosity and circular velocity, where the circular velocity depends on the dark matter and baryonic mass profiles. A key tool in studying galaxy evolution is the semi-analytic model approach to relating the observable properties of galaxies to the underlying formation physics \\citep{eis96,mo98}. Semi-analytic models of galaxy formation predict observables such as the luminosity function, radii, rotation curves, clustering statistics, colors, and the stellar mass of galaxies. \\cite{gne07} and \\cite{dut07} have shown that semi-analytic models can predict the joint distribution of galaxy observables, with model parameters that depend on the baryonic mass profile, where the baryonic mass profile includes the stellar mass and gas mass. Budgeting baryonic mass into stellar mass and gas mass is an important tunable parameter in the modeling \\citep{mcg05}. One major difficulty lies in converting multi-wavelength imaging data into the stellar mass, where the multi-wavelength imaging data is inherently noisy. Large surveys produce multi-wavelength maps of galaxies, which are used to measure the baryonic properties of the galaxy population. Large surveys produce large data sets, which are comparable in size to modern simulations \\citep{del06,bow06}. Since models of galaxy formation predict stellar mass and surface density, the multi-wavelength maps of galaxies must be converted into stellar populations using synthetic stellar population models \\citep{mar05,bru03}. In order to analyze the stellar populations in multi-wavelength images of galaxies, noise-reduction techniques must be employed. Large surveys are key to answering these questions because they provide data uniformity over a large area of the sky. SDSS provides $ugriz$-band data over 25\\% of the sky with a photometric calibration accuracy good to 2\\% \\citep{ive04}. The existence of these noisy, but uniform and large, data sets, along with the need for stellar population modeling, requires a smoothing technique for multi-wavelength data sets. A study similar to this one is (\\cite{lan07}; hereafter L07), which studies the pixel color magnitude relation for nearby galaxies. L07 noted the distinct difference of pixel color magnitude diagrams with different Hubble types, where Early-type galaxies have redder pCDMs. L07 noted how morphological features were related to distinct features in the pCDM. Scatter in a pCDM was caused by extinction, showing the need for accurate ISM extinction models when modeling observed colors. The study by L07 shows how pixel maps of galaxies are correlated with galaxy type, and might reveal hidden features. L07 does not employ noise-reduction techniques, such as the ones presented in this paper, which may affect the structure of the pixel diagrams. Another study similar to this one is \\cite{wel08}. \\cite{wel08} used the pixel-z technique, which combines stellar population synthesis models with multi-wavelength pixel photometry of galaxies to study the stellar population content of SDSS galaxies. \\cite{wel08} showed how the star formation rate varies with local galaxy density, varies with position in a galaxy, and studied the mean star formation rate. The pixel-z method does not include any smoothing techniques. This current work will be complementary to that work, by providing a smoothing technique that will minimize the effect data noise will have on the best-fit stellar population parameters. Adaptsmooth \\citep{zib09} is a multi-wavelength smoothing algorithm that is similar to this work. Adaptsmooth uses a circularly symmetric radial median to reduce noise. The radius of the circle is defined so that median smoothed data has a signal-to-noise of 20. All of the imaging wavelength data (i.e. $ugriz$-bands) are smoothed to the same radius, which is determined from the maximum radius of the multi-wavelength data in question, which is usually the $u$-band or $z$-band in SDSS data as they have the lowest signal-to-noise. Adaptsmooth is adaptive, in the sense that the radius varies with position in the galaxy, as the signal-to-noise varies. However, the radial median filter is still azimuthally symmetric. This means that blue star formation regions can get median filtered together with redder disk. The PCA-smoothing method presented in this paper, median filters over pixels that are associated in PCA space according to their color. We focus on SDSS because it is a large and uniform data-set. SDSS has coverage over 25\\% of the sky, where the imaging data covers a large range of optical wavelengths, and does so in a uniform manner. The distribution of galaxy properties has been well studied for SDSS data sets \\citep{bla03}. Many semi-analytic models have been constrained using SDSS data \\citep{gne07,li07}. In the following paper the technique is described in Section 2, a comparison to other methods is done in Section 3, case studies are presented in Section 4, and the conclusion is presented in Section 5. ", "conclusions": "This paper presents a method for smoothing SDSS data using a variation of Principal Component Analysis. The method is performed by running PCA simultaneously on multi-wavelength images of galaxies, and then smoothing over pixels that have similar locations in PCA space and spatial location within the galaxy. The advantages of the method are 1) no mixing of colors, 2) the method is geared towards stellar population analysis, 3) the parameters are tunable, and 4) the results are not extremely sensitive to the input parameters. The disadvantages of the method are 1) requiring initial analysis to identify the galaxy, 2) running PCA which may take computational time, and 3) requires well understood and uniform noise characteristic across different wavelengths. The smoothing parameters can be tuned to adjust the tradeoff between more smoothing and more color mixing versus less smoothing and more color purity. Increasing the $SNRenhanced$ constant, results in an increased signal-to-noise of the smoothed pixel, at the cost of mixing over different colors. Lowering the $SNRenhanced$ constant, results in a more pure color with less smoothing over different colors, at the cost of a lower smoothed signal-to-noise. The method was tested and demonstrated using a mock galaxy with $ugriz$-band images having SNRs similar to that seen in typical SDSS data. Figures ~\\ref{fig:IMSTATall} and ~\\ref{fig:ANALYZEHII} show that the $FOM$ for the PCA-smoothing method is always better (lower), when compared to azimuthally symmetric smoothing routines. Considering Figures ~\\ref{fig:ANALYZEHII} and Figure ~\\ref{fig:IMSTATall}, the best-fit parameters are $SNRmax=30.0$, $Radius=4.0$, $SNRenhanced=60.0$ and $SNRmin=3.0$. The lack of extreme peaks in the $FOM$ shows the robustness of the method. Figures ~\\ref{fig:IMSTATall} and ~\\ref{fig:ANALYZEHII} imply that as long as the user doesn't use extreme smoothing parameters, a reliable result will be obtained. Analysis of a region located on the boundary between an HII region and the red disk (Figure ~\\ref{fig:Deltagr}), shows that PCA smoothing is better at predicting a ($g$-$r$) color to 0.2 mag, when compared to simple radial smoothing or Adaptsmooth. The PCA-smoothing algorithm can be run on the SDSS data set, with the parameters described in this paper. The galaxies in the low-redshift NYU-VAGC \\citep{bla05} would be perfect for analysis, as it includes galaxies within a comoving distance range of $10 < d < 150$ Mpc/h. These nearby galaxies are spatially resolved, and perfect for this type of analysis. This method is geared towards large area surveys having multi-wavelength data, over a large part of the sky, and having uniform noise characteristics (i.e. COSMOS, DEEP, SDSS, 2MASS, DES). The method can also be applied to Galactic Nebulae as well, which are also asymmetrical extended objects with multi-wavelength data. \\clearpage" }, "1101/1101.1026.txt": { "abstract": " ", "introduction": "%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The simplest and most popular model of dynamical dark energy is quintessence, a single scalar field whose vacuum energy dominates the Universe driving its acceleration. Quintessence energy density varies with time and a way to distinguish it against a cosmological constant is to observe the effect of the different expansion history on dark matter structure formation \\cite{Wang:1998gt}. In its standard version, quintessence is described by a minimally-coupled canonical field \\cite{Zlatev:1998tr}. In this case scalar fluctuations propagate at the speed of light maintaining quintessence homogeneous even in the presence of dark matter clumps \\cite{Ferreira:1997au}. Quintessence can cluster only on scales larger than the horizon, where fluctuations have no time to propagate. However, observations on such large scales are strongly limited by cosmic variance and this effect is difficult to observe. A model of quintessence that can cluster on all observable scales has been recently proposed in \\cite{Creminelli:2008wc,Creminelli:2009mu}. It is based on a single scalar degree of freedom with fluctuations characterized by a practically zero speed of sound. As explained in \\cite{Creminelli:2008wc}, there are several theoretical motivations to consider this case. In the limit of zero sound speed one recovers the Ghost Condensate theory \\cite{ArkaniHamed:2003uy}, which is invariant under shift symmetry. Thus, there is no fine tuning in assuming that the speed of sound is very small: quintessence models with vanishing speed of sound should be thought of as deformations of this particular limit where shift symmetry is recovered \\cite{Creminelli:2006xe,Senatore:2004rj}. Moreover, using the tools developed in \\cite{Creminelli:2006xe,Cheung:2007st}, formulated in the context of an effective field theory, it has been shown that quintessence with an equation of state $w<-1$ can be free from ghosts and gradient instabilities only if the speed of sound is very tiny, $|c_s| \\lesssim 10^{-15}$ \\cite{Creminelli:2008wc}. Stability can be guaranteed by the presence of higher derivative operators \\cite{Creminelli:2006xe,ArkaniHamed:2003uy}, although their effect is absent on cosmologically relevant scales \\cite{Creminelli:2008wc}. Apart from these theoretical considerations, a very important motivation to consider this model is that a series of galaxy and cosmic shear surveys are currently planned with the aim of understanding the nature of dark energy through its role in the structure formation. In this context the clustering scenario represents a phenomenologically interesting counterpart to the case of a smooth quintessence component. Indeed, quintessence with vanishing speed of sound actively participates to the formation of structures together with the dark matter and gives distinct modifications to the standard picture that can be strongly constrained by future data. In the past, several articles have investigated the observational consequences of a clustering quintessence in the linear regime, in particular, on the cosmic microwave background \\cite{DeDeo:2003te,Weller:2003hw,BeanDore,Hannestad:2005ak,Sapone:2009mb}, galaxy redshift surveys \\cite{Takada:2006xs}, large neutral hydrogen surveys \\cite{TorresRodriguez:2007mk}, the cross-correlation of the integrated Sachs-Wolfe effect in the cosmic microwave background with the large-scale structures \\cite{Hu:2004yd,Corasaniti:2005pq}, or on weak lensing \\cite{Sapone:2010uy}. Theoretical investigations of the effect of dark energy on the nonlinear evolution of structures are particularly crucial. First of all, from linear theory alone it is difficult to distinguish the effects of quintessence on structure formation through its modification of the expansion history from those genuinely due to its perturbations. As we will see, the nonlinear evolution breaks this degeneracy. Furthermore, numerical simulations taking into account the gravitationally coupled evolution of dark matter particles and a clustering scalar field are still under construction and for the clustering scenario considered here they are totally missing. On the other hand, future redshift and weak lensing surveys will require very accurate predictions, both for the dark matter density and galaxy correlators, particularly on nonlinear scales where the signal is larger. Finally, to conclude this series of motivations we remind that the study of nonlinearities is receiving a lot of attention in the context of primordial non-Gaussianities \\cite{Komatsu:2009kd}. It is pertinent to ask whether a second clustering component could mimic the effect of primordial non-Gaussianities on the nonlinear evolution. A first description of clustering quintessence in the nonlinear regime was given in \\cite{Creminelli:2009mu}. There it was shown that in the limit of zero sound speed pressure gradients are negligible and, as long as the fluid approximation is valid, quintessence follows geodesics remaining comoving with the dark matter (see also \\cite{Lim:2010yk} for a more recent model with identical phenomenology). In particular, reference \\cite{Creminelli:2009mu} studied the effect of quintessence with vanishing sound speed on the structure formation in the nonlinear regime, in the context of the spherical collapse model (see \\cite{Bjaelde:2010qp} for a study of the spherical collapse when $c_s^2$ of quintessence is small but finite). Due to the absence of pressure gradients, comoving regions behave as closed FRW universes and the spherical collapse can be solved exactly. The modifications to the critical threshold of collapse are small and the effects on the dark matter mass function are dominated by the modification on the linear dark matter growth function, which are also small. Today they are of the order of few per cent for realistic values of $w$. A larger effect occurs when one considers the {\\em total} mass function, which includes the contribution of quintessence overdensities to the virialized halos. Indeed, quintessence contributes to the total halo mass by a fraction which increases at lower redshifts and is proportional to the ratio between quintessence and dark matter energy densities, {\\em i.e.}~$\\sim (1 + w)\\,\\Omega_Q/ \\Omega_m$. In this paper we study the nonlinear regime of clustering quintessence using Eulerian Perturbation Theory (EPT). In particular, we extend the standard EPT approach for dark matter \\cite{Peebles,Fry:1983cj,Bernardeau:2001qr} to the presence of a second fluid, a clustering quintessence, comoving and coupled only gravitationally to dark matter. This is the first natural step to the study of nonlinear perturbations beyond the spherical approximation. In contrast to the spherical collapse model, this approach is perturbative and solutions can be found order by order. Notice that on small scales the EPT perturbative expansion for density correlators is not well defined, because it presents large cancellations between contributions of the same order. However, in the standard case it has been shown that classes of higher-order corrections can be resummed, leading to a well established perturbative scheme known, in its first formulation, as Renormalized Perturbation Theory (RPT) \\cite{Crocce:2005xy,Crocce:2005xz,Crocce:2007dt,Bernardeau:2008fa,Bernardeau:2010md}. Complementary approaches can be found in \\cite{Matarrese:2007wc, Matarrese:2007aj,Pietroni:2008jx, Anselmi:2010fs, Saracco:2009df}. In this work, we begin by considering the continuity equations for the dark matter and quintessence density contrasts and the Euler equation for their common velocity. Since gravitational observables are sensitive only to the sum of dark matter and quintessence fluctuations, we derive the continuity equation for the {\\em total density contrast}. As both fluids are comoving, this equation and the Euler equation (together with the Poisson equation relating the total density to the gravitational potential) form a closed system in Fourier space, which can be solved perturbatively, as in the standard EPT approach. As in the standard case, the nonlinear couplings in the continuity and Euler equations are at most quadratic and the vertices are the same as those for a single dark matter fluid. The only difference is that the velocity divergence in the continuity equation is enhanced by the factor $(1 + w)\\, \\Omega_Q/ \\Omega_m$. At linear order, this term is responsible for a rapid evolution of the growth rate at low redshifts, and changes the standard relation between the velocity divergence and the growth factor. Due to the absence of pressure gradients, the solutions for the linear growth and the linear growth rate of the total fluid can be written in integral form, as in the standard $\\Lambda$CDM case. Using these solutions we are able to find simple fitting functions for these quantities, which generalize those currently employed in $\\Lambda$CDM cosmologies \\cite{Lahav:1991wc,Carroll:1991mt}. At higher order in the perturbative expansion clustering dark energy is responsible for an additional time-dependence of the kernels $F_n$ and $G_n$ defining the $n$-th order nonlinear corrections. The effect on $F_2$ and $G_2$ is of the order of the ratio between quintessence and total density perturbations, $\\sim {\\delta \\rho_Q}/({\\delta \\rho_m +\\delta \\rho_Q})$, and gives distinctive signatures in the higher-order correlation functions such as the bispectrum. In particular, the {\\em reduced} bispectrum, whose expression at leading order in EPT is independent of the linear power spectrum normalization, presents corrections {\\em only} in the clustering case. Analogous corrections have been found in the halo mass function from the contribution of the quintessence mass to collapsed objects \\cite{Creminelli:2009mu}. It is not the first time that EPT is generalized to the presence of several components. For instance, in \\cite{Somogyi:2009mh} EPT has been applied to the problem of following the nonlinear evolution of baryon and cold dark matter perturbations evolving from distinct initial conditions and in \\cite{Wong:2008ws,Shoji:2009gg,Saito:2008bp,Saito:2009ah,Lesgourgues:2009am,Brouzakis:2010md} to the study of nonlinear perturbations in the presence of massive neutrinos. For modified gravity models it has been used in \\cite{Koyama:2009me,Scoccimarro:2009eu, Chan:2009ew} to calculate the nonlinear power spectrum and, in particular, in \\cite{Scoccimarro:2009eu, Chan:2009ew, BernardeauBrax2010} to compute the matter bispectrum. Higher-order observables, such as the normalized skewness $S_3 \\equiv \\langle \\delta^3 \\rangle / \\langle \\delta^2 \\rangle^2$, have been also studied in \\cite{Multamaki:2003vs,Lue:2003ky,Amendola:2004wa} in the context of modified gravity models, where variations up to $\\sim 10\\%$ have been found. This paper is organized as follows. In section~\\ref{sec:eom} we present the equations of motion describing the coupled evolution of matter and quintessence perturbations. In section~\\ref{sec:linear} we solve the linearized equations for the density growth factor and the density growth rate and we discuss the solutions and fitting formulae. In section~\\ref{sec:nonlinear} we discuss the perturbative solutions in EPT. In particular, we derive the second-order solutions for the density and velocity fields, while in section~\\ref{sec:bispectrum} we derive the lowest-order observables: the density tree-level power spectrum and bispectrum. As a practical illustration of these results, in section~\\ref{sec:StoN} we compare the signal-to-noise expected for the effect of clustering and smooth quintessence with respect to the $\\Lambda$CDM case in ideal measurements of the density large-scale power spectrum, bispectrum and reduced bispectrum in a box of $1\\cGpc$ at redshift $z=0.5$. Finally, we present our conclusions in section~\\ref{sec:conclusions}. In addition, we present in appendix \\ref{app:scalar} a discussion on the analogy between the scalar field and the perfect fluid, with a derivation of the continuity, Euler and Poisson equations in the regime considered in this paper. In appendix \\ref{app:vertices} we derive evolution equations for the vertices in the spherical collapse approximation at all orders and in appendix \\ref{app:zdist} we discuss the redshift-space distortion effects in the clustering quintessence case. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ", "conclusions": "\\label{sec:conclusions} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In this work we considered the case of a quintessence characterized by a vanishing speed of sound. In this case quintessence perturbations grow on all observables scales, inducing relevant effects on the evolution of structures when the dark energy component comes to dominate the energy density of the Universe. At late time, both dark matter and quintessence perturbations act as a source for the gravitational potential and they are practically indistinguishable by gravitational observations. Based on this, we introduce a {\\em total density perturbation} as a weighted sum of matter and quintessence perturbations. Since quintessence is {\\em comoving} with dark matter, the evolution of such a quantity is determined by a closed set of equations: the continuity equation for the total density field, the Euler equation for the common velocity field of the two components, and the Poisson equation relating the gravitational potential to the total density perturbation. This allows us to study the evolution of the total density fluctuation in Eulerian Perturbation Theory, in complete analogy with the usual treatment of matter fluctuations in a $\\Lambda$CDM or smooth quintessence cosmology. The equations of motion for the total perturbations are equivalent to those for the matter perturbation alone, with a simple correction: the linear term in the velocity divergence of the continuity equation is proportional to the function $C= 1 + (1+w) \\Omega_Q/\\Omega_m$. At early times, when quintessence is negligible, $C=1$ and we recover the standard evolution. At late times clustering quintessence increases the growth rate of fluctuations. At linear order it is possible to obtain an exact integral expression for the growth of the total perturbation, eq.~\\eqref{int_sol}. This solution relies on the fact that for a quintessence with vanishing speed of sound comoving regions behave as independent FRW universes. This integral expression allowed us to derive simple fitting functions for both the growth function and the growth rate of the total fluctuations. Beyond linear order, our set-up allows to straightforwardly apply to the clustering quintessence scenario standard EPT, but also more efficient resummation techniques such as Renormalized Perturbation Theory \\cite{Crocce:2005xy, Crocce:2005xz, Crocce:2007dt, Bernardeau:2008fa,Bernardeau:2010md} and the Renormalization Group approach \\cite{Matarrese:2007wc, Matarrese:2007aj,Pietroni:2008jx, Anselmi:2010fs, Saracco:2009df}. We showed that linear theory does not fully describe the rich phenomenology of a quintessence with zero speed of sound. Indeed, we found significant effects on the late-time evolution of higher-order perturbations. These can affect the total power spectrum over a wide range of observable scales, where the evolution of perturbations becomes nonlinear. Since they directly depend on nonlinear corrections over the linear density field, also higher-order correlation functions, such as the bispectrum, are affected by the clustering of quintessence at low redshift. In particular, we studied second-order solutions in EPT for the density contrast and velocity fields. In terms of these solutions we derived the leading-order (or tree-level) contribution in EPT to the total bispectrum. On large scales, this is expected to be a good approximation to the fully nonlinear bispectrum. In particular, we showed that the {\\em reduced} bispectrum, which is normalized in such a way as to be independent of the linear evolution, receives significant corrections {\\em only} in the clustering case. These corrections are of the order of $\\delta \\rho_Q/(\\delta \\rho_m + \\delta \\rho_Q)$, \\ie~the ratio between the quintessence and total density perturbations, which at $z=0$ amounts to $5\\%$ for $|1+w| =0.1$. These signatures offer a practical way of distinguishing the clustering scenario from the smooth one with the next generation of redshift and weak lensing surveys such as BOSS or Euclid. Corrections of the same magnitude to the reduced matter bispectrum are expected as well from non-Gaussian initial conditions. However, notice that, at least on large scales, such corrections present a different (in fact, opposite) redshift evolution as well as different dependences on scales and shapes (see, for instance, \\cite{Liguori:2010hx}). In Section~\\ref{sec:StoN} we provided a simple estimate of the signal-to-noise ratio expected for the {\\em effect} of quintessence on the power spectrum and bispectrum of the total density field. We limited our analysis to an ideal box of $1\\cGpc$ at fixed redshift $z=0.5$. In particular, for the linear power spectrum we considered the signal expected for the {\\em difference} between the predictions of smooth quintessence and $\\Lambda$CDM, $P_{Q_s} - P_\\Lambda$, of clustering quintessence and $\\Lambda$CDM, $P_{Q_c} - P_\\Lambda$, and of clustering and smooth quintessence, $P_{Q_c} - P_{Q_s}$. We performed the same analysis for the tree-level bispectrum and the reduced bispectrum. In order to provide a fair comparison between power spectrum and bispectrum, we included all measurable triangular configurations down to a given $k_{\\max}$. Below a given scale, the signal-to-noise ratio of the bispectrum becomes more important than the one of the power spectrum. Interestingly, this takes place at a larger scale for the clustering case. This preliminary analysis is clearly very limited. We discussed simply the density perturbations without considering a particular observable. Moreover, we did not include higher-order nonlinear corrections. These are expected to be relevant for wavenumbers close to the maximum value considered here, \\ie~$k_{\\rm max} = 0.2\\kMpc$. Nevertheless, our results emphasize the importance of a joint analysis of the power spectrum {\\em and} bispectrum in future redshift and weak lensing surveys, possibly extending over a large redshift range to take full advantage of the rich time-dependence in the clustering quintessence scenario. Furthermore, they suggest that smaller scales, where the evolution of perturbations is nonlinear, are likely to be affected significantly by quintessence clustering. Thus, extending this analysis to the nonlinear regime---somehow a necessary task---can significantly improve the constraints on dark energy with respect to those forecasted assuming only linear theory predictions. Given the significantly different phenomenology between the smooth and clustering cases, it is crucial to develop accurate theoretical predictions to compare with cosmological observations. As shown, the inclusion of higher-order perturbations, either as corrections to the power spectrum or in higher-order statistics, can be very important in this process. In this work we presented a first step in this direction using perturbative techniques based on Eulerian Perturbation Theory. We leave to future study the extension to more efficient resummation schemes. \\\\ \\bigskip {\\bf Acknowledgments}\\\\ We are grateful to Francis Bernardeau, Pierstefano Corasaniti, Paolo Creminelli, Guido D'Amico, Enrique Gazta\\~ naga, Christian Marinoni, Rom\\'an Scoccimarro and Atsushi Taruya for useful discussions. ES acknowledges support by the European Commission under the Marie Curie Inter European Fellowship and he is grateful to the Center for Cosmology and Particle Physics of New York University for kind hospitality during the completion of this project. \\appendix %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%" }, "1101/1101.1486_arXiv.txt": { "abstract": "We investigate the distribution of bright main sequence stars near the northern edge of the M33 disk. Clustering on sub-kpc scales is seen among stars with ages $\\sim 10$ Myr, and two large star-forming complexes are identified. Similar large-scale grouping is not evident among stars with ages 100 Myr. These stars are also distributed over a much larger area than those with younger ages, and it is argued that random stellar motions alone, as opposed to orderly motions of the type spurred by large scale secular effects, can re-distribute stars out to distances of at least 2 kpc (i.e. one disk scale length) from their birth places on 100 Myr timescales. Such random motions may thus play a significant role in populating the outer regions of the M33 disk. Finally, it is suggested that -- to the extent that the ambient properties of the outer disk mirror those in the main body of the disk -- stars in this part of M33 may have formed in star clusters with masses 50 -- 250 M$_\\odot$, which is substantially lower than the peak of the solar neighborhood initial cluster mass function. ", "introduction": "The stellar disks of spiral galaxies can extend to many scale lengths (e.g. Davidge 2006; Pohlen \\& Trujillo 2006), and the stars that populate the peripheral regions of disks likely have a range of origins. Some of the stars at large radii probably formed {\\it in situ}. Ultraviolet light concentrations that trace young stellar regions are seen at large radii in some nearby spiral galaxies (e.g. Gil de Paz et al. 2008; Zaritsky \\& Christlein 2007). While the density of interstellar material at large radii tends to be too low to trigger large-scale star formation, localized density enhancements may occur as a result of compression from spiral density waves (e.g. Bush et al. 2008). The presence of a dark baryonic component in the disk plane could also enable star formation in areas where the gas density may otherwise appear to be too low (Revaz et al. 2009). Some fraction of the stars in the outer disk are probably migrants from smaller radii, as recent studies have shown that secular processes (e.g. Roskar et al. 2008) and radial mixing induced by interactions (e.g. Quillen et al. 2009) can contribute significantly to the stellar contents of the outer regions of disks. The distribution of stars in the nearest spiral galaxies provide clues into the processes that populate the outermost regions of disks. Star-forming activity at large radii produces distinct signatures in galactic light and color profiles (e.g. Sanchez-Blazquez et al. 2009). As for secular effects, the processes that re-distribute stars throughout disks act in a cumulative manner on stellar orbits, with the result that stars that have moved the furthest from their places of birth will tend to be the oldest -- stars with progressively older ages may thus be found at progressively larger galactocentric distances (e.g. Roskar et al. 2008), in contradiction to what might niavely be expected due to inside-out disk formation. The present letter is part of a larger study of young and intermediate age stars throughout the disk of M33 (Davidge et al. 2011, in preparation). The entire dataset consists of five MegaCam pointings, and here we examine the distribution of stars in an area that includes two of the most remote star-forming complexes in M33. A distance modulus of 24.93 (Bonanos et al. 2006) is adopted. Recent distance modulus estimates for M33 show a spread of a few tenths of a dex, and the Bonanos et al. value was selected because it is based on eclipsing binaries, which are a primary distance indicator. ", "conclusions": "The star-star separation function (S3F) has been used to investigate the projected distribution of main sequence stars in the northern disk of M33. Two young stellar complexes produce significant signal in the S3F of stars with ages $\\sim 10$ Myr at separations $r < 150$ arcsec ($d < 0.7$ kpc). However, signatures of clustering are greatly diminished among stars with ages $\\sim 40$ Myr, and the smooth S3F of stars with ages $\\sim 100$ Myr suggests that there is little if any large-scale clustering among these stars. Thus, large scale stellar structures in this part of M33 evidently dissipate over time scales $\\leq 100$ Myr. Stellar complexes in the outer regions of disks may be subjected to disruption mechanisms that differ from those in the main body of the disk. There is evidence for heating by halo structures in the outer regions of nearby galaxies (e.g. Martin \\& Kennicutt 2001), and dynamical measurements suggest that halo bombardment becomes a significant source of heating at 4 disk scale lengths in nearby spirals (Herrmann et al. 2009), and this is the part of the M33 disk that we examine here. The broad, evenly distributed signal in the 100 Myr S3F between 150 and 450 arcsec (0.7 -- 2.1 kpc) results from random motions on the order of $\\sim 20$ km sec$^{-1}$, and this is comparable to the outer disk extraplanar motions measured by Herrmann et al. (2009). Putman et al. (2009) find that HI in the outer regions of M33 has a velocity dispersion of 18.5 km sec$^{-1}$, and suggest that this may be a relic of an interaction within the past few Gyr between M31 and M33. Newly formed stellar systems will be more prone to disruption if they have a low star formation efficiency (SFE), as feedback will remove gas early-on, thereby reducing -- perhaps catastrophically -- the gravitational field of the nascent system (Lada \\& Lada 2003). A general trend for the SFE to diminish towards larger radii is seen in nearby galaxies (Leroy et al. 2008). This result is based on measurements made over kpc spatial scales, which is comparable to the sizes of the large structures probed here, Star clusters are sub-structures within the large-scale complexes that are investigated here. The largest disk star clusters in M33 subtend $\\leq 2$ arcsec (San Roman et al. 2010), and so fall in the smallest bin in the S3F. The signal in the S3F of the 100 Myr sample in the 0 - 20 arcsec bin is markedly smaller than in the 10 Myr sample, and if this trend extends to sub-arcsec sizes then this will be consistent with stellar clusters dissipating over $\\sim 0.1$ Gyr timescales. In fact, the spatial distribution of star clusters with ages $< 0.1 - 0.3$ Gyr in M33 is more compact than that of stars with the same age (Sarajedini \\& Mancone 2007; Roman et al. 2010), suggesting that young star clusters in M33 dissipate over time spans that are less than a few tenths of a Gyr. The disruption timescale of star clusters depends on a number of factors, including the rate at which remnant gas is removed from the cluster, the local environment, and two-body relaxation (e.g. summary by Elmegreen \\& Hunter 2010). Gratier (2010) find that the masses of molecular clouds decrease with increasing radius in M33, and this should result in lower star cluster masses, which in turn may lead to a comparatively rapid disruption timescale for clusters in the outer regions of M33. Lamers et al. (2005a) estimates that a $10^4$ M$_{\\odot}$ cluster in the main body of M33 typically disrupts after $\\sim 1$ Gyr. Assuming no radial changes in the sources of dynamical heating, the ambient mass mixture that dominates the gravitational field, and the mean SFE within M33, then if the cluster disruption timescale $\\propto$ mass$^\\gamma$, where $\\gamma =$ 0.62 (Baumgardt \\& Makino 2003; Lamers et al. 2005b), then the majority of stars in the outer disk of M33 formed in clusters with masses $\\leq 50 - 250$ M$_{\\odot}$ if they are disrupted on timescales of $\\sim 100$ Myr. This characteristic cluster mass is roughly two orders of magnitude lower than the peak of the solar neighborhood initial cluster mass function predicted by Parmentier et al. (2008) and Kroupa \\& Boily (2002). In fact, this is an upper limit to the initial cluster mass, in the sense that the pace with which clusters dissolve depends on factors such as the local mass density and the initial cluster mass, and a 10$^4$ M$_{\\odot}$ cluster in the peripheral regions of the M33 disk would be even longer lived than predicted by Lamers et al. (2005a). Thus, if star clusters are disrupted over $\\sim 0.1$ Gyr timescales in the outer regions of M33 then we predict that the star clusters found there will have (1) young ages, and (2) lower masses than those at smaller radii. We close by noting that the orderly large scale motions induced by secular processes are probably not significant among stars of the age considered here, given that the rotation period of the M33 disk is 200 - 300 Myr (Corbelli \\& Salucci 2000). Rather, stars with ages $\\sim 100$ Myr in this part of M33 appear to have obtained random stellar motions that allow them to populate regions up to $\\sim 2$ kpc from where they formed. This effectively pushes out the observational boundary of the young disk. \\parindent=0.0cm" }, "1101/1101.3483_arXiv.txt": { "abstract": "We report on a $\\sim$63\\,ks \\CXO\\, observation of the X-ray transient Swift\\,J195509.6+261406\\, discovered as the afterglow of what was first believed to be a long duration Gamma-Ray Burst (GRB\\,070610). The outburst of this source was characterized by unique optical flares on timescales of second or less, morphologically similar to the short X-ray bursts usually observed from magnetars. Our \\CXO\\, observation was performed $\\sim$2 years after the discovery of the optical and X-ray flaring activity of this source, catching it in its quiescent state. We derive stringent upper limits on the quiescent emission of Swift\\,J195509.6+261406 which argues against the possibility of this object being a typical magnetar. Our limits show that the most viable interpretation on the nature of this peculiar bursting source, is a binary system hosting a black hole or a neutron star with a low mass companion star ($< 0.12 M_{\\odot}$), and with an orbital period smaller than a few hours. ", "introduction": "\\label{intro} On 2007 June 10 the \\Swift\\, Burst Alert Telescope (BAT) triggered on GRB\\,070610, a typical long-duration GRB (see Gehrels et al.~2007 for a recent review), with a $\\sim$4.6\\,s high-energy prompt emission (Pagani et al.~2007; Tueller et al.~2007). Follow-up soft X-ray observations with the \\Swift\\, X-ray Telescope (XRT) started soon after the event, discovering only one variable X-ray source within the BAT error circle: namely Swift\\,J195509.6+261406\\ (Kasliwal et al.~2008; \\swt\\, hereafter). This transient X-ray source was very different from what expected for the X--ray afterglow of a long GRB: it was decreasing in flux rather slowly, and it showed a strong X-ray flaring activity. The source became undetectable by \\Swift-XRT on 2007 June 29, ranging from a 0.5--10\\,keV flux of $\\sim10^{-9}$ to $< 10^{-12}$\\ergscm2 \\, in 19 days. While in outburst, \\swt\\, had an X-ray spectrum that could be described by a rather hard power-law corrected for the photoelectric absorption ($N_{\\rm H} =7\\times10^{21}$\\cm2 \\, and $\\Gamma$=1.7). Due to spatial and temporal coincidence (it was the only transient source in the BAT error circle), GRB\\,070610 and the X--ray transient \\swt\\, have been associated with high probability (Kasliwal et al.~2008). The most interesting and peculiar features of this transient source came from optical and infrared observations. Many telescopes, triggered by the GRB-like event, promptly observed the position of \\swt\\, during the outburst. A highly variable optical and infrared counterpart was observed, showing large flares for about 11 days after the GRB-like event, when it went back to quiescence. These large flares were characterized by a very short timescale: during the largest flare the source increased its optical flux by more than a factor of 200 in less than 4\\,s. Furthermore, a broad quasi periodic oscillation was observed in the optical band at $\\sim$0.16\\,Hz (Stefanescu et al. 2008). The source distance was constrained by several different methods to be within 3.7--10\\,kpc (mainly red clump study, and detailed measurements of the absorption column in the {\\em mm} waveband; Castro-Tirado et al. 2008). Furthermore, the stringent optical and IR limits derived in the quiescent level (H$>$23; R$>$26.0 and $i^{\\prime}>$24.5; Kasliwal et al. 2008; Castro-Tirado et al. 2008) constrain the type of any companion star to either a main-sequence star with spectral type later than M5V (which means a mass $<$0.12~M$_{\\odot}$), or to a semi-degenerate hydrogen poor star (Castro-Tirado et al. 2008). The large variations of its optical and infrared counterpart during the decay to quiescence, its distance and Galactic nature, set this transient apart from the typical optical afterglows of long-duration GRBs (see Liang et al.~2007 for a recent review). The resemblance of the optical bursts of \\swt\\, with the short X-ray bursts from magnetars (see Mereghetti 2008 for a recent review) led to the idea of a new kind of X-ray and optical transient event in a Galactic magnetar (Castro-Tirado et al. 2008; Stefanescu et al. 2008). On the other hand, its X-ray flaring activity was also proposed to resemble the emission of the fast X-ray nova V4641 Sgr (Markwardt et al. 2008; Kasliwal et al. 2008), an unusual 9~M$_{\\odot}$ black hole in orbit with a 5--8~M$_{\\odot}$ B9 III companion star (in't Zand et al. 2000; Orosz et al. 2001) In this Letter we present the results of a $\\sim$63\\,ks \\CXO\\, observation of \\swt\\, (see \\S\\,\\ref{obs}) aimed at unveiling its X-ray properties during quiescence (see \\S\\,\\ref{results}), and compare them with the current quiescent levels of the magnetar and X-ray binary populations (see \\S\\,\\ref{discussion}). ", "conclusions": "We derived deep upper limits with \\CXO\\, on the X-ray quiescent emission of the optical bursting transient \\swt . We showed that a magnetar scenario is very unlikely: the source is too faint in quiescence for any realistic scenario of magnetar cooling. We suggest that \\swt\\, is most likely an X-ray binary, hosting a black hole or a neutron star with an orbital period faster than a few hours, possibly in an ultra-compact system. High-time resolution optical observations of X-ray binaries during outburst might reveal energetic optical flares, a peculiarity that \\swt\\, does not share yet with any other source." }, "1101/1101.1954_arXiv.txt": { "abstract": "We present the stellar kinematics in the central 2\\arcsec\\ of the luminous elliptical galaxy M87 (NGC~4486), using laser adaptive optics to feed the Gemini telescope integral-field spectrograph, NIFS. The velocity dispersion rises to 480~\\kms\\ at 0.2\\arcsec. We combine these data with extensive stellar kinematics out to large radii to derive a black-hole mass equal to $(6.6\\pm0.4)\\times10^9~\\Msun,$ using orbit-based axisymmetric models and including only the NIFS data in the central region. Including previously-reported ground-based data in the central region drops the uncertainty to $0.25\\times10^9~\\Msun$ with no change in the best-fit mass; however, we rely on the values derived from the NIFS-only data in the central region in order to limit systematic differences. The best-fit model shows a significant increase in the tangential velocity anisotropy of stars orbiting in the central region with decreasing radius; similar to that seen in the centers of other core galaxies. The black-hole mass is insensitive to the inclusion of a dark halo in the models --- the high angular-resolution provided by the adaptive optics breaks the degeneracy between black-hole mass and stellar mass-to-light ratio. The present black-hole mass is in excellent agreement with the Gebhardt \\& Thomas value, implying that the dark halo must be included when the kinematic influence of the black hole is poorly resolved. This degeneracy implies that the black-hole masses of luminous core galaxies, where this effect is important, may need to be re-evaluated. The present value exceeds the prediction of the black hole-dispersion and black hole-luminosity relations, both of which predict about $1\\times10^9~\\Msun$ for M87, by close to twice the intrinsic scatter in the relations. The high-end of the black hole correlations may be poorly determined at present. ", "introduction": "The masses of central black holes in galaxies appear to be closely related to the luminosity (Dressler 1989; Kormendy 1993; Kormendy \\& Richstone 1995; Magorrian et al. 1998) and stellar velocity dispersion (Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000) of their host galaxies. These relationships, which are determined from local samples of galaxies, provide the means to assay the cosmological mass distribution function of massive black holes, and provide the empirical foundation for establishing the role of black holes in galaxy formation and evolution (e.g. Hopkins et al. 2008). At present the black-hole galaxy-property relationships are derived from several dozen black-hole mass determinations made over the last few decades (see G\\\"ultekin et al. 2009). The relationships remain poorly observed at both their high-mass and low-mass ends. Lauer et al. (2007) show, for example, that the $M_{BH}-\\sigma$ and $M_{BH}-L$ relationships must be in conflict at high black-hole mass due to curvature in the Faber \\& Jackson (1976) relationship between galaxy velocity dispersion and luminosity. Small uncertainties in the high-mass end of the relations can lead to uncertainties of up to two orders of magnitude in the implied volume density of black holes with $M_{BH}>10^9~\\Msun,$ due to the high-end exponential cutoff of the galaxy luminosity and velocity-dispersion distribution functions. Such estimates also depend critically on knowledge of the intrinsic scatter in the the relationships (G\\\"ultekin et al. 2009). Thus, there remains a need to measure accurate black-hole masses in a sample of the most massive galaxies. Apart from the need to enlarge the sample of galaxies used to define the black-hole galaxy-property relationships, it appears that we may also need to test and potentially revise some black-hole mass measurements already made, especially in the massive ``core galaxies.'' Recent work shows that black-hole masses are subject to several systematic errors that have not been generally incorporated in the models used for analyzing the data so far. Some of these are discussed in G\\\"ultekin et al. (2009) and include the radial variations in the mass-to-light ratio due to changes in stellar populations or the presence of a dark halo, uncertainties in the deprojection of the surface brightness, and triaxiality, among others. The most important of these systematic effects are: {\\it Dark halo:} Gebhardt \\& Thomas (2009) show that the measured black-hole mass for M87 increases by more than a factor of two when a dark halo is included in the models; the reason for the change is that the black-hole's kinematic influence is poorly resolved in the data that they use, so that there is substantial covariance between the black-hole mass and stellar mass-to-light ratio. In-turn the best-fit stellar mass-to-light ratio, assumed independent of radius, is affected by whether or not a dark halo is included in the models. It is well understood that the mass-to-light profile for ellipticals changes with radius and not including that trend biases the black hole determination. An obvious, but challenging, solution to this degeneracy is to obtain data at radii where the kinematics are strongly dominated by the black hole rather than the stars. {\\it Incomplete orbit library:} Shen \\& Gebhardt (2010) find an increase of two in the black-hole mass for NGC~4649 when using a more complete orbital sampling compared to models using a less coverage (Gebhardt et al. 2003). They argue that the orbital structure near the black hole is dominated by tangential orbits and that the older models did not have adequate coverage of these tangential orbits (as discussed in Thomas et al. 2004). Having too few tangential orbits (i.e., too many radial orbits) can be compensated by having a smaller black-hole mass. {\\it Triaxiality:} Van den Bosch \\& de Zeeuw (2009) find an increase of two in the measured black-hole mass for NGC~3379 by using triaxial models compared to triaxial models (although they find the same black-hole mass for M32). All three of these systematic effects tend to increase the black-hole mass. The increases are generally larger than the statistical uncertainties and suggest that systematic effects still dominate. By observing stars close to the black hole, many model assumptions are no longer needed. For example, if the gravitational potential is dominated by the black hole, then the stellar contribution to the enclosed mass is not important; hence, uncertainties in the stellar mass-to-light ratio, which may arise from uncertainties in the dark-halo properties, can be mitigated by probing well inside the influence region of the black hole. Since it is among the most luminous galaxies nearby, has the largest black hole known (from spatially resolved kinematics) and has one of the nearest and best-studied AGNs, M87 is a natural and important target. An accurate black-hole mass determination for M87 helps to pin down the sparsely sampled upper end of the black-hole mass distribution and provides insights into formation and evolution of the most luminous galaxies. The previous analysis of M87 from Gebhardt \\& Thomas is based on ground-based kinematic data taken in natural seeing under moderately good conditions (FWHM=1\\arcsec). In this paper, we present kinematics based on the integral field spectrograph, NIFS, on the Gemini Telescope, taken with adaptive optics correction. The spatial FWHM of the kinematics is 0.1\\arcsec\\ on average, with the best seeing image at 0.08\\arcsec. At larger radii we incorporate new kinematic data out to 245\\arcsec\\ or 2.5 effective radii that will appear in a companion paper. The extreme improvement in the data quality of M87 allows us to model black-hole mass with smaller systematic uncertainty. This paper focuses on the determination of the mass of the central black hole; the analysis of the stellar mass-to-light ratio and the dark halo properties will be given in Murphy et al. (2011). Obtaining the kinematics at spatial resolution down to 0.1\\arcsec, at the same signal-to-noise obtained here, would have required about 100 orbits (90 hours) of Hubble Space Telescope, due to the faint stellar surface brightness. This adaptive optics study using Gemini/NIFS took about 10 hours in total, highlighting one of the great advantages for ground-based adaptive optics. We assume a distance to M87 of 17.9 Mpc. The value of the black-hole mass scales linearly with assumed distance. ", "conclusions": "\\subsection{M87 Specific Results} Our best-fit black-hole mass is $(6.6\\pm0.4)\\times10^9\\Msun$. Sargent et al. (1979) report a black-hole mass of $6\\times10^9\\Msun$ (after scaling to our assumed distance), which is within 1-sigma of our reported value. Their model is based on lower spatial resolution data (about 1.5\\arcsec), assumes that the velocity distribution is isotropic, and does not include a dark halo. It is impressive that after three decades of improvement in data quality, modeling, and understanding, there is essentially no change in the measured black-hole mass. Part of the reason for the robustness of the Sargent et al. result is that the radial influence on the projected kinematics from the black hole extends to nearly 10\\arcsec\\ (see Fig.~5), so the influence of the black hole was clearly visible in their kinematic data. They also use isotropic models, whereas we run axisymmetric models with no restrictions on the anisotropy. To study the effect of the assumption of isotropy, we fit isotropic models to the kinematic data presented in this paper. The comparison between the projected dispersions of the isotropic models and the data is poor, with an increase in $\\chi^2$ by over a factor of two. The poor fit makes it difficult to assign a best-fit mass and the range of equally poor fitting models have black-hole masses that range from $6-8\\times10^9~\\Msun$, consistent with the models of \\S3, which show significant tangential anisotropy (Fig. 7). Thus, in M87, the assumption of isotropy does not have a significant effect on the measurement of the black-hole mass, although isotropic models provide a poor fit to the data. Sargent et al. also do not include a dark halo, which has been shown to cause the black hole to be underestimated. Their velocity dispersions at large radii are lower than ours (245 compared to 300~\\kms), which is most likely because their template library was incomplete and their spectra had lower S/N. The lower dispersion causes the assumed mass-to-light ratio of the stars to be lower, an error of the opposite sign to the error caused by neglect of the dark halo. Thus, the impressive agreement between our value and that of Sargent et al. (1978) appears to be due in part to the competing effects of observational errors (dispersions too small, which makes the stellar mass-to-light ratio too low and the black-hole mass too large) and oversimplified models (no dark halo or velocity anisotropy, both of which make the black-hole mass too small). Another often-quoted black-hole mass determination from stellar kinematics comes from Magorrian et al. (1998) who report a value of $4.2\\times10^9\\Msun$ (for our distance). The likely reason for the difference is that they do not include a dark halo and thus overestimate the stellar mass-to-light ratio. The black-hole mass reported here is nearly the same as that reported in Gebhardt \\& Thomas (2009), within 4\\%. There is very little kinematic data in common between the two studies. The kinematic data in Gebhardt \\& Thomas come from older long-slit data at spatial resolution of 1.0\\arcsec\\ (van der Marel 1994), while in this paper we use two-dimensional coverage at spatial resolution of 0.08\\arcsec. We further use ground-based data from Murphy et al. (2011) that have excellent S/N and radial extent. There is some data from SAURON (Emsellem et al. 2004) in common between the two studies, but this provides only 10\\% of the LOSVDs used in the models. Thus, the dynamical models from the two studies use nearly independent kinematic datasets, and give approximately the same answer. The uncertainties on the black-hole mass from these two studies are similar even though the data presented here are superior in many ways; the previous uncertainty is $0.5\\times10^9$ whereas the uncertainty with the NIFS data is $0.4\\times10^9$. In order to keep the black-hole mass measures independent, the models presented in this paper do not include the SAURON data inside of 2.5\\arcsec. The similarity in the black-hole mass uncertainty is then due primarily to the fact that the two sets of data have similar accuracy on the kinematics in the central 2.5\\arcsec. Combining all NIFS data, the accuracy on the velocity dispersion is 0.2\\% (1~\\kms). Combining all SAURON data within 2.5\\arcsec\\ provides the same accuracy. Thus, as long as one has a reliable PSF and no systematic differences in the kinematic extractions, then it is expected that the uncertainty on the black-hole mass is similar using either dataset. We have run a subset of models including both the NIFS data and all SAURON data; in this case, the uncertainty on the black-hole mass decreases to $0.25\\times10^9$ (with no change in the best-fit mass). We report and utilize the result using only the NIFS data within 2.5\\arcsec\\ in order to 1) provide as independent result as realistically possible and 2) control potential systematic differences in the kinematic extractions. Murphy et al. (2011) find a difference in the velocity dispersion of the SAURON data at large radii compared to their measurements, which they attribute to template issues. While we do not find an offset in the dispersion values in the central region, we desire to maintain the independence. The major difference, however, is that there is no degeneracy with the stellar mass-to-light ratio using the NIFS data, whereas the degeneracy is very strong otherwise. Thus, the systematic uncertainty from the mass-to-light ratio profile is effectively removed with the adaptive optics data, making the result on the black-hole mass and orbital structure much more robust. For M87, the AO data has removed the systematics due to the mass-to-light ratio profile but the systematics due to the extraction of the kinematics remain important. These systematics include continuum placement, template mismatch, and removal of the AGN contribution. The first two are general and the latter is specific to M87. Getting any of these controlled to better than 1\\% of the velocity dispersion will be very difficult. Corrected to our distance, the black-hole masses reported from gas kinematics are ($2.9\\pm0.8)\\times10^9~\\Msun$ in Harms et al. (1994) and ($3.8\\pm1.1)\\times10^9~\\Msun$ in Macchetto et al. (1997). As discussed in Gebhardt \\& Thomas (2009) the mass reported here is in conflict with these by about 2-sigma. Possible reasons for the differences are discussed in Gebhardt \\& Thomas, with the most likely reason being uncertainty in the inclination of the gas disk. Macchetto et al. assume a value of 51 degrees based on the gas kinematics. Harms et al. assume a value of 42 degrees based on the imaging of the gas emission. The reported difference provides a measure of the systematic uncertainty in the inclination (i.e., whether the gas kinematics or the gas distribution are more affected by non-gravitational forces). Applying this 9 degree difference in the inclination changes the Macchetto et al. black-hole mass from $(3.8\\pm1.1$) to $(5.4\\pm1.3)\\times10^9~\\Msun$, which would lead to an insignificant difference of 0.6-sigma between our result and theirs. Of course, the analysis is more complicated than this simple application since one would need to re-model the gas kinematics with a different inclination. A proper treatment would be to include the gas kinematics with the stellar dynamical models. Our focus in this paper is on the stellar kinematics, and we do not attempt to merge the gas kinematic analysis. \\subsection{General Implications for Black-Hole Mass Measurements} While the kinematics obtained from the adaptive optics study produce effectively the same black-hole mass and its uncertainty from kinematics taken in native image quality, the robustness of the measures is greatly strengthened. For example, the black-hole mass is not dependent on the assumption of constant mass-to-light ratio. Trying to generalize this result to other galaxies with black-hole mass determinations is difficult since the measure of the black-hole mass depends on many aspects. There are two observational extremes that we highlight as examples. The first is having a measure of black-hole mass that comes from observations that resolve well the kinematic influence of the black hole. In the most extreme case, high S/N spectra could potentially see the high velocity wings in the LOSVD due to the black hole (as discussed in van der Marel 1994). The second example would be to allow poorer resolution of the black hole but provide a very accurate measure of the mass-to-light profile. In this paper, we rely on the first strategy; Gebhardt \\& Thomas rely on the second. That the two strategies give consistent results, at least for M87, suggests that both may be reliable. Other studies have reported robust measures of the black-hole mass from ground-based studies that only poorly resolve the black hole's kinematic influence. Shapiro et al. (2006) measure a black-hole for NGC~3379 from SAURON data that is consistent with that measured from {\\it HST} data using both stars (Gebhardt et al. 2000c) and gas kinematics. Kormendy (2004) summarizes the history of black-hole mass measures for many galaxies and finds that, in general, the differences are within the reported uncertainties. If one has sufficient signal-to-noise and two-dimensional coverage (e.g., SAURON or VIRUS-P), then it should be possible to measure a black-hole mass robustly. Thus, it is not necessarily required to resolve the region influenced by the black hole. Being able to use data that does not well resolve the black hole's influence on the kinematics allows us to study black holes that are either distant or low mass. Both of these regimes are important for understanding the physical nature of the black hole correlations with the host galaxy. For example, McConnell et al. (2011) measure a black hole mass in NGC~6086, which is 133 Mpc distant. The kinematic influence of the black hole is barely resolved, and the degeneracy between the black hole mass and M/L profile is strong. However, as demonstrated for M87, as long as one properly characterizes the mass profile at large radii, then high signal-to-noise data can measure the black hole mass accurately. It is possible that systematic uncertainties bias the current crop of black-hole correlations. One obvious consequence could be that without accounting for the effect of systematic uncertainties, the measured intrinsic scatter would increase. G\\\"ultekin et al. (2009) measure scatter of 0.44 dex for the full sample of galaxies with measured black-hole masses and 0.31 dex for ellipticals. Once systematic effects are understood and included, the intrinsic scatter may decrease. Other consequences include increasing the mass density of black holes, if black-hole masses are all underestimated, and changing the slope or curvature of any correlation. Schulze \\& Gebhardt (2011) re-analyse the set of 12 galaxies from Gebhardt et al. (2003) including a dark halo. They find an increase of 50\\% in the black-hole mass, due primarily to improved dynamical modeling (more complete orbit sampling) and partly to including a dark halo. The increase correlates with black-hole mass. It is important to re-evaluate all black-hole mass estimates. The key to understanding all of these effects comes from high spatial resolution data. Data from Hubble Space Telescope (mainly from STIS) is generally regarded as providing the most significant results for black-hole mass studies. The small and stable PSF is a central aspect for the robustness of the data from HST. Future uses of STIS will play an important role for quantifying black-hole masses. The main obstacle for HST though is that it is a relatively small mirror and requires substantial observing time. For example, in order to measure the black-hole mass in M87 at the same accuracy presented here would require nearly 100 orbits. While this amount of time could be justified for a small number of objects, going to a much larger sample using HST is difficult. Fortunately, adaptive optics observations are in a mature stage where they can provide much larger samples." }, "1101/1101.5020_arXiv.txt": { "abstract": "We report detailed timing and spectral analysis of {\\em RXTE}-PCA data obtained from observations during the outburst of a transient X-ray pulsar 1A 1118--61 in January 2009. The pulse profile showed significant evolution during the outburst and also significant energy dependence $-$ a double peaked profile upto 10 keV and a single peak at higher energy. We have also detected quasi-periodic oscillations (QPO) at 0.07--0.09 Hz. The rms value of the QPO is 5.2 $\\%$ and it shows a significant energy dependence with highest rms of 7$\\%$ at 9 keV. The QPO frequency changed from 0.09 Hz to 0.07 Hz within 10 days. The magnetic field strength calculated using the QPO frequency and the X-ray luminosity is in agreement with the magnetic field strength measured from the energy of the cyclotron absorption feature detected in this source. The 3-30 keV energy spectrum over the 2009 outburst of 1A 1118--61 can be well fitted with a partial covering power-law model with a high energy cutoff and an iron fluorescence line emission. The pulse phase resolved spectral analysis shows that the partial covering and high energy cutoff model parameters have significant changes with the pulse phase. ", "introduction": "The hard X-ray transient pulsar 1A 1118--61 was discovered with the Rotation Modulation Collimator ({\\em RMC}) experiment on {\\em Ariel V} in 1974 \\citep{Eyles1975}. Pulsations were detected in this source with a period of 405.3 s \\citep{Ives1975} and the optical counterpart of this source was identified as He 3--640 = WRA 793 \\citep{Chevalier1975} which is a highly reddened Be star. The star has a visual magnitude of V = 12.1 and is classified as a O9.5IV-Ve \\citep{Janot-Pacheco1981, Motch1988} with strong Balmer emission lines indicating the presence of an extended envelope. The UV spectrum shows many absorption features; especially the CIV line indicating a stellar outflow. The P-Cygni profile gives a wind velocity in the range of 1600 $\\pm$ 300 km s$^{-1}$ and the general spectral profile is similar to that of the optical counterparts of other transient systems \\citep{Coe1985}. The extinction value of $A_v$ = 2.8 $\\pm$ 0.3 mag. suggested the distance to be 4 kpc. From the Corbet diagram \\citep{Corbet1984} for high magnetic field accreting pulsars, the orbital period is expected to be around 350 days. The known correlation between the orbital period of Be star binaries and $H_{\\alpha}$ EW of the optical companion also indicates a similar large orbital period \\citep{Reig1997} for 1A 1118--61. However, recently, detection of a 24-day binary period was reported by \\cite{Staubert2010} using {\\em RXTE}/PCA data. {\\em Einstein} and {\\em EXOSAT} observations of this source were carried out in 1979 and 1985 respectively. During these observations the source was in a quiescent state and the luminosity calculated from the {\\em EXOSAT}/ME observations was 0.5-3.0 $\\times$10$^{34}$ ergs s$^{-1}$ at 3-7 kpc. This confirms that in the quiescent state, centrifugal inhibition of accretion was not complete. Three outbursts have so far been detected in this source. First outburst was in 1974 \\citep{Maraschi1976}. The source had a second outburst that was first detected with Burst and Transient Source Experiment (BATSE) on the {\\em Compton Gamma Ray Observatory} ({\\em CGRO}) in 1991/1992. During this period the source had a peak flux of about 145 mCrab and a spin down rate of 0.016 s/day \\citep{coe1994}. The source was also observed by the WATCH all sky monitor on {\\em Granat} \\citep{Lund1992}. 1A 1118--616 was in quiescence for $\\sim$20 years and became highly active in 2009 January. The main outburst lasted only for about $\\sim$20 days. This third outburst, observed by {\\em Swift}/XRT in January 2009 revealed a pulsation period of 407.68 s \\citep{Mangano2009a} indicating a spin-down in between the outbursts. About three weeks after the start of the outburst, the source was also observed with the {\\em International Gamma-Ray Astrophysics Laboratory} ({\\em INTEGRAL}/JEM-X/ISGRI) which detected a flaring activity after the main outburst \\citep{Leyder2009}. Many observations of this source were carried out with {\\em Rossi X-ray Timing Explorer}/Proportional Counter Array ({\\em RXTE}/PCA) during this period, and the combined analysis of the {\\em RXTE}-PCA and High Energy X-ray Timing Explorer (HEXTE) spectra revealed a cyclotron line absorption feature at 55 keV which gives a magnetic field strength of 4.8 $\\times$ 10$^{12}$ G for the neutron star \\citep{Doroshenko2010}. Most of the transient High Mass X-ray Binary (HMXB) pulsars are known to have a Be star companion. In Be/X-ray binaries, X-ray emission is thought to be due to the accretion of matter by the neutron star from the slow, dense, radial outflow of the Be star \\citep{Negueruela1998}. These systems are observed to exhibit two different types of X-ray outbursts. One is short X-ray outbursts (Type-I outbursts) lasting for a few days ($L_x \\leq 10^{36} - 10^{37}$ ergs s$^{-1}$) occuring in several successive binary orbits at orbital phase close to the time of periastron passage and other is giant X-ray outbursts (Type-II outbursts) lasting for several weeks ($L_x \\geq 10^{37}$ ergs s$^{-1}$) and may start at any orbital phase. The current outburst in 1A 1118-61 appears to be a Type-II outburst but of shorter duration and smaller peak luminosity than the Type-II outbursts in most other Be/X-ray binary pulsars. We have carried out a detailed timing and spectral analysis of the {\\em RXTE}-PCA observation of this source during the 2009 January outburst, to detect any intensity or energy dependence of the pulse profile, aperiodic variabilities and also to find a suitable spectral model in the 3-30 keV band. One component of the aperiodic variabilities seen in X-ray binaries is the Quasi Periodic Oscillations (QPO), generally thought to be related to the innermost regions of the accretion disk. Any inhomogeneous matter distribution or blobs of material in the inner disk may result in QPOs in the power spectrum. This can give useful information about the interaction between accretion disk and the central object at different intensity levels. HMXB pulsars show QPOs only at low frequency, i.e. in the range of 10 mHz upto about 1 Hz. Black hole X-ray binaries and low magnetic field neutron stars show QPOs over a wide range of frequency from a few Hz to a few hundred Hz. Studying QPOs and their variations with energy and luminosity gives important clues about the mechanism of QPO production. In section 2 we describe the observations and the data used in the present work. In section 3 we present the pulsation analysis, the power density spectra, the pulse phase averaged and pulse phase resolved spectroscopy using the {\\em RXTE}-PCA archival data followed by a discussion of the results in section 4. ", "conclusions": "\\subsection{Quasi Periodic Oscillations} QPOs have so far been detected in 19 accretion powered high magnetic field pulsars which include mostly HMXBs and a few LMXBs, in both transient and persistent sources. In Table 2 we have listed the sources, the spin frequency $\\nu_{s}$, the QPO frequency $\\nu_{QPO}$ and its range, and the ratio of the two. \\begin {table*} \\caption{List of QPO sources} ~\\\\ \\begin {tabular}{|c|c|c|c|c|c|} \\hline Source&Type&$\\nu_{s}$&$\\nu_{QPO}$&$\\nu_{QPO}$/$\\nu_{s}$&Reference\\footnotemark{}\\\\ & &(mHz)&(mHz)&&\\\\ \\hline Transient pulsars& & & &\\\\ \\hline KS 1947+300&HMXB/Be&53&20&0.38&1\\\\ SAX J2103.5+4545&HMXB/?&2.79&44&15.77&2\\\\ A0535+26&HMXB/Be&9.7&50&5.15&3\\\\ V0332+53&HMXB/Be&229&51&0.223&4\\\\ 4U 0115+63&HMXB/Be&277&62&0.224&5\\\\ 1A 1118--61&HMXB/Be&2.5&92&36.8&This work\\\\ XTE J1858+034&HMXB/Be&4.53&110&24.3&6\\\\ 4U 1901+03&HMXB/?&361.9&130&0.359&7\\\\ EXO 2030+375&HMXB/Be&24&200&8.33&8\\\\ SWIFT J1626.6-5156 &HMXB/Be&65&1000&15.38&9\\\\ XTE J0111.2-7317&HMXB/B0.5-B1Ve&32&1270&39.68&10\\\\ GRO J1744-28&LMXB&2100&20000&9.52&11\\\\ \\hline Persistent pulsars& & & &\\\\ \\hline SMC X-1&HMXB/B0&1410&10&0.0071&12\\\\ Her X-1&LMXB&806&13&0.016&13\\\\ LMC X-4&HMXB/O-type&74&0.65-20&0.0087-0.27&14\\\\ Cen X-3&HMXB/O-type&207&35&0.17&15,16\\\\ 4U 1626-67&LMXB&130&48&0.37&17,18\\\\ X Per&HMXB/Be&1.2&54&45&19\\\\ 4U 1907+09&HMXB/OB&2.27&69&30.4&20,21\\\\ \\hline \\end{tabular} \\vskip 1cm \\footnotetext{}{References: (1) \\cite{James2010}; (2)\\cite{Inam2004}; (3)\\cite{Finger1996}; (4) \\cite{Takeshima1994}; (5) \\cite{Soong&Swank1989}; (6) \\cite{Paul&Rao1998}; (7) \\cite{James2011}; (8) \\cite{Angelini1989}; (9) \\cite{Reig2008}; (10) \\cite{Kaur2007}; (11) \\cite{Zhang1996}; (12) \\cite{Angelini1991}; (13) \\cite{Moon2001b}; (14) \\cite{Moon2001a}; (15) \\cite{Takeshima1991}; (16) \\cite{Raichur2008}; (17) \\cite{Shinoda1990}; (18) \\cite{Kaur2008}; (19) \\cite{Takeshima1997}; (20) \\cite{Zand1998}; (21) \\cite{Mukerjee2001} } \\end{table*} The most commonly used models for explaining the QPO mechanism are Keplerian frequency model (KFM), beat frequency model (BFM) and accretion flow instabilities. In the KFM, the QPOs arise due to inhomogeneities at the inner edge of the accretion disk modulating the light curve at the Keplerian frequency. In the BFM, the accretion flow onto the neutron star is modulated at the beat frequency between the Keplerian frequency of the inner edge of the accretion disk and the spin frequency $\\nu_{QPO}= \\nu_{k} -\\nu_{s}$ \\citep{Shabazaki&Lamb1987}. The third model applies only to the sources that have luminosities close to the Eddington limit \\citep{Fortner1989}. The occurrence of QPOs in accretion powered pulsars is quite complex. In some sources like 4U 1626--67, QPOs are detected most of the time while in some others like Cen X-3, QPOs are rare. Recently, in two transient pulsars (KS 1947+300 and 4U 1901+03) QPOs were observed only at the end of the outbursts when the source intensity had fallen to a few percent of the peak of the outbursts \\citep{James2010, James2011}. Some transient pulsars like XTE J1858+034, showed QPOs in all observations while some other transient pulsars like EXO 2030+375, showed QPOs in only some of the outbursts. In the 405 s recurrent transient pulsar 1A 1118--61, QPOs are observed in the range of 0.07 - 0.09 Hz in most of the observations near the peak of the outburst. The frequency evolution of the QPO in 1A 1118--61 during the outburst and presence of QPOs during most of the outburst makes this source similar to the transient XTE J1858+034 \\citep{Mukherjee2006}. In 1A 1118--61, the QPO frequency is higher than the spin frequency of the pulsar and can be explained in either KFM or BFM and since the QPO frequency is a few hundred times larger than the spin frequency, the radius at which the QPOs are generated is quite similar in the two models. We have, \\begin{equation} R_{qpo} = \\left({GM}\\over{4 \\pi^{2} \\nu^{2}_{qpo}}\\right)^{1/3} \\end{equation} For an assumed neutron star mass of $1.4 M_{\\odot}$, $R_{qpo}$ = 9$\\times$ 10$^{3}$ km. The average 3-30 keV X-ray flux during the period of QPO detection is 1.95$\\times$10$^{-9}$ erg cm$^{-2}$ sec$^{-1}$, which for a distance of 4 kpc corresponds to an X-ray luminosity of $0.37\\times10^{37}$ erg sec$^{-1}$. The radius of the inner accretion disk can be expressed in terms of the luminosity and magnetic moment as \\cite{Frank1992} \\begin{equation} R_{m}= 3\\times10^{8}L^{-2/7}_{37}\\mu^{4/7}_{30} \\end{equation} where $L_{37}$ is the X-ray luminosity in units of 10$^{37}$ erg s$^{-1}$ and $\\mu_{30}$ is the magnetic moment in units of 10$^{30}$ cm$^{3}$ Gauss. Equating $R_{qpo}$ with $R_{m}$, we determine a magnetic moment of $\\mu_{30}$ = 3.38 $\\times$10$^{30}$ cm$^{3}$ G, which for a neutron star canonical radius of 10 km corresponds to a surface magnetic field strength of 3.38 $\\times$10$^{12}$ G. The magnetic field strength of the neutron star derived using the QPO frequency and the X-ray luminosity is in excellent agreement with the strength of the magnetic field obtained from the cyclotron line absorption feature (4.8 $\\times$10$^{12}$ G, Doroshenko et al. 2010). \\subsection{Pulse profile evolution and energy dependence} Accretion powered X-ray pulsars are known to show interesting intensity dependence of the pulse profile, both in transient and persistent sources \\citep{Nagase1989, Raichur2010}. A most remarkable example of pulse profile evolution was investigated in EXO 2030+375 with EXOSAT observations \\citep{Parmar1989} . Evolution of the pulse profile during the outbursts indicate changes in the accretion flow from the inner accretion disc to the neutron star, especially changes in the mass accretion rate. The X-ray beaming pattern may also change during the outburst depending on the density, dimensions and structure of the accretion column. The 2-60 keV pulse profiles of 1A 1118--61, presented here also show a complex profile, and profile evolution, as the relatively short outburst reached its peak and then decayed. The pulse profiles shown in Figure 2 are clearly intensity dependent, but it is not a simple function of intensity. For example, for similar count rates per PCU, the pulse profile is quite different during the rise of the outburst (7-10 January 2009, Figure 2) and during the decay (22-23 January 2009, Figure 2). Most accretion powered pulsars also show strong energy dependence of the pulse profile. In 1A 1118--61, the position of the main peak is at the same phase in all energies in the 2-30 keV band, but several other features including a second peak at low energy, a narrow dip in the low energy and a leading edge of the main peak at high energy are energy dependent features clearly seen in energy resolved pulse profiles (Figure 3). An intriguing aspect is that, the pulse shape change appears at around 10 keV. Strong energy dependence of pulse profiles - in other words strong pulse phase dependence of the energy spectrum - are known to be present also in magnetars \\citep {Enoto2010} and multi component emission models are invoked there to describe the same. However, we find that multi component spectral model for accretion powered pulsars, including a low temperature black body cannot explain the change in pulse profile at around 10 keV. For pulse phase averaged and pulse phase resolved spectroscopy, we have therefore used a model that can fit features in the medium energy band. \\subsection{Pulse phase averaged and pulse phase resolved spectra} The continuum spectra of accreting X-ray pulsars are often best described by a power law with a high-energy cutoff and interstellar absorption. And in most cases, the presence of a Gaussian component, for iron line fluorescent emission is also evident. Cyclotron resonance absorption features have also been detected in about 20 bright pulsars, usually at energies above 10 keV. In 1A 1118-61, Doroshenko et al. (2010) detected a cyclotron absorption feature at $\\sim$55 keV using the combined PCA and HEXTE data in the 4-120 keV band. Considering the complex energy dependence of the pulse profile and from pulse phase resolved spectroscopy we find that a partial covering power-law model describes the data very well. With a partial covering model we don't see a feature at 8 keV. The 8 keV feature mentioned in Doroshenko et al. (2010) and several other contemporary papers probably has same reason that a very simple, single continuum model is used while the sources have complex absorption. If the 8 keV feature is an instrument artifact, it should have been seen in all kinds of sources. However, such a feature has not been reported in the X-ray spectrum of black hole binaries and is also not seen in the spectrum of the Crab Nebula \\citep{Kirsch2005,Weisskopf2010}. We would also like to note here that at higher energy, the partial covering absorption model is the same as a simple power-law model, and thus the cyclotron absorption feature reported by Doroshenko et al. (2010) is not in question here. Lin et al. (2010) fitted the 0.5-10.0 keV energy spectrum of this source obtained with the {\\em Swift}-XRT and found that in this limited energy band, the spectra are fitted best with two black body components. However, since the source shows strong hard X-ray emission (as detected with {\\em Swift}-BAT, Figure 1) and the two blackbody model cannot produce the hard X-ray photons, this model is inappropriate for this pulsar as well as for any hard X-ray pulsar. In the partial covering absorption model, a part of the continuum source is obscured, resulting in a harder spectrum \\citep{Wang2010}. If the absorbing component is in the form of an accretion stream or is a part of the accretion column, it can be phase locked with the neutron star, resulting into a phase dependent column density and covering fraction. The pulse phase resolved spectroscopy reported here shows a strong modulation of the partial covering fraction as well as the column density supporting such a scenario for this pulsar. We also notice a systematic variation of the cutoff energy value and the e-folding energy at the pulse phases with highest column density of the partial covering material. A similar energy dependence of the pulse profile and pulse phase dependence of the spectral parameters have recently been found in another accretion powered transient pulsar GRO J1008-57 \\citep{Naik2010} in the broad band data obtained with the {\\em Suzaku} X-ray observatory." }, "1101/1101.0605_arXiv.txt": { "abstract": "We report on the performance of our cold-dark matter cosmological $N$-body simulation which was carried out concurrently using supercomputers across the globe. We ran simulations on 60 to 750 cores distributed over a variety of supercomputers in Amsterdam (the Netherlands, Europe), in Tokyo (Japan, Asia), Edinburgh (UK, Europe) and Espoo (Finland, Europe). Regardless the network latency of 0.32 seconds and the communication over 30.000 km of optical network cable we are able to achieve $\\sim 87$\\% of the performance compared to an equal number of cores on a single supercomputer. We argue that using widely distributed supercomputers in order to acquire more compute power is technically feasible, and that the largest obstacle is introduced by local scheduling and reservation policies. ", "introduction": "Some applications for large scale simulations require a large amount of compute power. This is often hard to acquire on a single machine. Combining multiple supercomputers to do one large calculation can lift this limitation, but such wide area computing is only suitable for certain algorithms. And even then the political issues, like arranging the network, acquiring the compute time, making reservations, scheduling runtime and synchronizing the run start, and technical limitations are profound. Earlier attempts based on interconnecting PC clusters were quite successful \\cite{inteugrid,Gualandris2007,Manos,Bal20083,QCGEscience2009}, but lacked the raw supercomputer performance required for our application. Running simulations across multiple supercomputers has been done a few times before \\cite{IWay,paragon,arthropod,CosmoGrid}, though the performance of simulations across three or more supercomputers has not yet been measured in detail. Here we report on the performance of our parallel astronomical simulations which use up to 4 supercomputers and predict the performance for simulations which use 5 or more supercomputers. In our experiments we use an international infrastructure of supercomputers. These machines include an IBM Power6 supercomputer located at SARA in Amsterdam (the Netherlands) and three Cray XT-4 supercomputers located at the Edinburgh Parallel Computing Centre in Edinburgh (United Kingdom), the IT Center for Science in Espoo (Finland) and the Center For Computational Astrophysics in Tokyo (Japan). The Edinburgh site is equipped with a 1 Gbps interface while the other three sites are equipped with a 10 Gbps interface. We achieved a peak performance of 0.610 TFLOP/s and a sustained performance of 0.375 TFLOP/s using 120 cores distributed over 4 sites. To provide a comparison with the international tests we also run the code over up to 5 sites on a Dutch grid of PC clusters. Our wide area simulations are realized with the development of a software environment for Simulating the Universe Structure formation on Heterogeneous Infrastructures, or SUSHI for short. ", "conclusions": "We have run a few dozen cosmological $N$-body simulations and analyzed the scalability of our SUSHI integrator on a national distributed computer and across a global network of supercomputers. Our results confirm that SUSHI is able to efficiently perform simulations across supercomputers. We were able to run a simulation using $1024^3$ particles across three supercomputers with $\\sim 10\\%$ communication overhead. The communication performance can be further improved by tuning the optical networks. Based on our model predictions we conclude that a long-term cosmological simulation using $2048^3$ particles and $256^3$ mesh cells scales well over up to $\\sim 16$ sites, given that sufficient bandwidth is available and the number of cores used per site is limited to $\\sim 256$. We also predict that tree codes with a shared time step scheme run efficiently across multiple supercomputers, while tree codes with a block time step scheme do not. Considerable effort is still required to obtain acceptable message passing performance through a long distance optical network. This is due to three reasons. First, it may take up to several months to arrange an intercontinental light path. Second, optical networks are generally used for high volume data streaming such as distributed visualization or bulk data transfer, and are therefore not yet tuned to achieve optimal message passing performance. Third, intercontinental networks traverse a large number of different institutes, making it politically difficult for users to diagnose and adjust settings on individual sections of the path. For our experiments we therefore chose to optimize the wide area communications by tuning our application, rather than requesting system-level modifications to the light path configuration. The main challenges in running simulations across supercomputers are now political, rather than technical. During the GBBP project, we were able to overcome many of the political challenges in part due to good will of all organizations involved and in part through sheer patience and perseverance. However, orchestrating a reservation spanning across multiple supercomputers is a major political undertaking. The use of a meta-scheduler and reservation system for supercomputers and optical networks greatly reduces this overhead, and also improves the workload distribution between supercomputer centers. Once the political barriers are overcome, we will be able to run long lasting and large scale production simulations over a grid of supercomputers." }, "1101/1101.1581.txt": { "abstract": "We predict the performance of the \\planck\\ satellite in determining the bulk flow through kinetic Sunyaev-Zeldovich (kSZ) measurements. As velocity tracers, we use ROSAT All-Sky Survey (RASS) clusters as well as expected cluster catalogs from the upcoming missions \\planck\\ and eRosita (All-Sky Survey: EASS). We implement a semi-analytical approach to simulate realistic \\planck\\ maps as well as \\planck\\ and eRosita cluster catalogs. We adopt an unbiased kinetic SZ filter (UF) and matched filter (MF) to maximize the cluster kSZ signal to noise ratio. We find that the use of \\planck\\ CMB maps in conjunction with the currently existing ROSAT cluster sample improves current upper limits on the bulk flow determination by a factor $\\sim5$ ($\\sim10$) when using the MF (UF). The accuracy of bulk flow measurement increases with the depth and abundance of the cluster sample: for an input bulk velocity of 500 km/s, the UF recovered velocity errors decrease from 94 km/s for RASS, to 73 km/s for \\planck\\ and to 24 km/s for EASS; while the systematic bias decreases from 44\\% for RASS, 5\\% for \\planck, to 0\\% for EASS. The $95\\%$ upper limit for the recovered bulk flow direction $\\Delta\\alpha$ ranges between $4^{\\circ} ~\\rm{and} ~ 60 ^{\\circ}$ depending on cluster sample and adopted filter. The kSZ dipole determination is mainly limited by the effects of thermal SZ (tSZ) emission in all cases but the one of EASS clusters analyzed with the unbiased filter. This fact makes the UF preferable to the MF when analyzing \\planck\\ maps. ", "introduction": "Peculiar velocities, along with inhomogeneities, can be used to constrain cosmology. A coherent, large scale peculiar velocity, also called bulk flow, may originate from spatial inhomogeneities in the mass distribution of large scale structures around us. The standard inflationary model predicts that the rms bulk velocity within a sphere of radius R decreases linearly with comoving distance in the $\\Lambda$CDM universe with $V_{\\rm rms}(r>50-100\\mpch)\\approx250(\\frac{100 \\mpch}{r})$ km/s~\\citep{Kashlinsky1991}. Galaxy cluster peculiar velocity surveys at scales $R\\le60 \\mpch$ generally agree with theoretical predictions of the cluster bulk velocities. However, recent measurements at larger scales ($R\\ge100 \\mpch $) indicate that the bulk flow velocity is significantly larger than the $\\Lambda$CDM prediction with statistical significance up to $3\\sigma$~\\citep{Kashlinsky2008,Feldman2010}. At scales $R\\le60 \\mpch $, an enhancement of the bulk flow with respect to the predicted $\\Lambda$CDM value in the local universe is attributed to a large scale void or overdensity at these depth. This is thought to be the cause of the Local Group (LG) motion with respect to the CMB rest frame, with velocity $v=627\\pm22$ km/s towards $l=276^{\\circ}$, $b=+30^{\\circ}$, in alignment with the CMB dipole~\\citep{Kogut1993}. The measured value is within the cosmic variance limit.~\\citet{Dressler1987} and~\\citet{Lynden1988} identified the Great Attractor (GA), a mass concentration of $\\sim10^{16}M_{\\odot}$, as the origin of the flow at $<60 \\mpch $. On large scales where $\\Lambda$CDM predicts bulk flows with negligible amplitude, observations show the contrary; ~\\citet{Lauer1994} found a strong bulk flow signature with $v=561\\pm284$ km/s towards $l=220^{\\circ}$, $b=-28^{\\circ}$ ($\\pm27^{\\circ}$) at a depth of 110 $\\mpch$. Using a sample of 119 Abell clusters within $150 \\mpch $ they found that the flow originate from a mass concentration beyond $100 \\mpch $. Similarly,~\\citet{Feldman2010} found a bulk flow on scales of $\\sim 100 \\mpch $ with $v=416\\pm78$ km/s towards $l=282^{\\circ}\\pm11^{\\circ}$, $b=+6^{\\circ}\\pm6^{\\circ}$, in disagreement with the WMAP5 cosmological parameters at 99.5\\% confidence. The direction and scale of these bulk flows are shown in Figure~\\ref{f:dipole}. Several theories have been suggested to explain the high bulk flow velocities at large scales. One explanation is that pre-inflationary fluctuations in scalar fields on superhorizon scales gives a titled universe (\\citealt{Turner1991},~\\citealt{Kashlinsky1994},~\\citealt{Mersini2009}). In this picture, matter slides from one side of our Hubble volume to the other, producing an intrinsic CMB dipole anisotropy as seen in the matter rest frame. This inhomogeneity generates a bulk flow with correlation length of order the horizon size. Alternatively,~\\citet{Wyman2010} have showed that a strengthened gravitational attraction at late times can speed up structure formation and increase peculiar velocities. Using N-body simulations, they found an enhancement in large scale, $R>100 \\mpch $, bulk flow velocities of up to $\\sim40\\%$ relative to the $\\Lambda$CDM cosmology. A similar approach using modified gravity is discussed in~\\citet{Afshordi2009} and~\\citet{Khoury2009}. \\citet{Kashlinsky2000} have proposed a method aimed at determining the largest scale bulk flows from galaxy cluster peculiar velocities measured using the kinetic Sunyaev-Zeldovich (kSZ) effect. If many galaxy clusters are moving with a coherent motion with respect to the CMB rest frame, the kinematic part of the SZ signal acquires a dipole moment. Since the kSZ signal is proportional to line of sight velocity, such a measurement directly probes the bulk flow, free of distance measurement errors. Several authors attempted to measure the bulk flow using the measured kSZ effect in \\wmap\\ data. \\cite{Kashlinsky2008} (hereafter KAKE, and later~\\citealt{Kashlinsky2010}) first utilized this method, claiming a large-scale flow with $v>600$ km/s out to $\\sim575 \\mpch $, without sign of convergence to the $\\Lambda$CDM predicted value. However, by repeating the same method~\\citet{Keisler2009} did not detect a statistically significant bulk flow.~\\citet{Osborne2010} (hereafter \\osb), used filters constructed to enhance the signal to noise of the kinetic signal and found no significant velocity dipole in the WMAP 7 year data. More specifically, they found a 95\\% bulk flow upper limits of the order of 4600 km/s in the direction of the KAKE They also showed that the matched filter outperforms the unbiased one when \\wmap\\ data are used, and demonstrated that CMB and instrument noise dominate the uncertainties. In this work, we apply the scheme of \\osb\\ to study the capability of \\planck\\ data and future cluster surveys to measure bulk flows. The use of Planck maps is expected to produce improved results with respect to the WMAP case because of reduced instrument noise, wider frequency coverage (ensuring better foregrounds' subtraction) and increased spatial resolution of this mission. In addition, ~\\planck\\ will also produce the first all-sky SZ survey with a median redshift of $z=0.3$ (\\planck\\ Blue Book) and will ensure a better performance in separating the tSZ from the kSZ signal for any cluster sample considered. We investigate to what extent the use of Planck maps, in combination with data for existing ROSAT clusters, improves on bulk flow determination from WMAP. We also assess the expected performances of bulk flow measurements for upcoming all--sky cluster catalogs, such as the ones derived from \\planck\\ and eRosita satellites. Such samples are more abundant and extend to higher redshifts than the one in hand. Our goals are: (i) to determine the sensitivity of the cluster velocity dipole measurement with the \\planck\\ specifications and assess the nature of the uncertainty; (ii) to study which cluster survey can best constrain the bulk flow; (iii) to study the performance of the filters used in \\osb\\ with the \\planck\\ setup. This paper is organized as follows. The bulk flow velocity expected from the $\\Lambda$CDM model is calculated in section~\\ref{s:theory}. In section~\\ref{s:outline}, we briefly describe the procedure we use to measure the bulk flow velocity. In sections~\\ref{s:cat} and~\\ref{s:szmap} we give details of the SZ and X-ray cluster catalogs we use and describe the procedure we adopt to generate simulated SZ maps. In section~\\ref{s:filter}, we present the two filters we use to reconstruct the kSZ signal from the CMB maps. In section~\\ref{s:dipole}, we describe the analysis pipeline we use to measure and calibrate the cluster dipole. % and calibrate the measurement. In section~\\ref{s:error}, we describe the systematic effects that may contaminate our results. The results are presented in section~\\ref{s:results}, followed by our conclusions in section~\\ref{s:con}. Throughout this paper, we assume a $\\Lambda$CDM cosmological model with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $h=0.72$, $w=-1$, $\\sigma_8=0.8$. \\begin{figure*} \\begin{center} \\leavevmode \\includegraphics[width=90mm, angle=90]{figures/dipole.eps} \\caption{The dipole direction in Galactic coordinates of current bulk flow measurements. (a) CMB $b=48.26\\pm0.03^{\\circ}$, $l=263.99\\pm0.14^{\\circ}$~\\citep{Jarosik2010}; (b) Local Group $b=30\\pm3^{\\circ}$, $l=276\\pm3^{\\circ}$~\\citep{Kogut1993}; (c) KAKE $b=34^{\\circ}$, $l=267^{\\circ}$; (d)~\\citet{Watkins2009} $b=8\\pm6^{\\circ}$, $l=287\\pm9^{\\circ}$; (e)~\\citet{Lauer1994} $b=-28\\pm27^{\\circ}$, $l=220\\pm27^{\\circ}$; (f)~\\citet{Hudson2004} $b=0\\pm11^{\\circ}$, $l=263\\pm13^{\\circ}$; (f)~\\citet{Feldman2010} $b=6\\pm6^{\\circ}$, $l=282\\pm11^{\\circ}$. The size of the colored region is proportional to the amplitude of the measured bulk flow. The color code represents the convergence depth of the bulk flow, in units of $\\mpch $. } \\label{f:dipole} \\end{center} \\end{figure*} ", "conclusions": "\\label{s:con} We investigated \\planck\\ performances in determining bulk flow velocities using the kinetic Sunyaev-Zeldovich effect. We characterize the sensitivity of the bulk flow velocity measurement using simulated \\planck\\ data combined with three representative all sky galaxy cluster surveys: the archived ROSAT All-Sky Survey, the \\planck\\ cluster catalog and the eRosita All-Sky Survey. We employ two different types of filters, a matched filter (MF) and an unbiased kinetic SZ filter (UF), to maximize the cluster signal to noise ratio. The main results are (see also Table~\\ref{t:summary}): \\begin{enumerate} \\item The use of simulated \\planck\\ sky maps instead of \\wmap\\ ones in combination with the RASS catalog reduces the velocity that can be detected at $95\\%$ CL, from $\\sim5000$ km/s ($\\sim10,000$ km/s) to $\\sim1000$ km/s ($\\approx1000$ km/s) when filtered by the MF (UF). \\item Using all clusters with $z\\le0.5$ we find that a bulk flow of 500 km/s would be measured in a strongly biased way with the RASS sample with both filters. The same velocity is measured with 0--5\\% bias with EASS and Planck clusters, with uncertainties of 25--75 km/s. These numbers should be compared with the rms velocity out to $z=0.5$ in $\\Lambda$CDM: 30 km/s (95\\% upper limit: 48 km/s). \\item If the bulk flow is consistent with $\\Lambda$CDM prediction of $v=30$ km/s at $z=0.5$, our analysis pipeline would obtain a 95\\% upper limit to the recovered velocity (when the UF is used) of $v=470$ km/s for RASS cluster, $v=160$ km/s for \\planck\\ cluster sample and $v=60$ km/s for EASS cluster sample. This allows us to measure the departure from $\\Lambda$CDM if the measured bulk flow is in excess of these estimates. \\item \\planck\\ can also constrain a recovered bulk flow direction. For $V_{\\rm input}=500$ km/s, the $95\\%$ upper limit to the angle errors (when the UF is used) are $\\Delta\\alpha_{95}\\approx30^{\\circ}$ for the RASS clusters, $\\Delta\\alpha_{95}\\approx15^{\\circ}$ for the \\planck\\ clusters, and $\\Delta\\alpha_{95}\\approx5^{\\circ}$ for the EASS clusters. These uncertainties are lower than the discrepancies observed in bulk flow directions at different optical depth up to now (Figure~\\ref{f:dipole}). The errors in the directions contained within the galactic plane are larger than in the perpendicular one due to the galactic cut at low latitudes. \\item The error in the kinetic SZ dipole is dominated by the unfiltered CMB and instrument noise and intrinsic scatter of the Y-parameter, with a systematic thermal SZ bias. The contamination from extragalactic radio sources is negligible and can be safely ignored. The noise level is lowest for the EASS sample since it contains the most clusters. \\item The UF is more sensitive and effective than the MF in suppressing the bias induced by the thermal SZ in the velocity reconstruction. This is in contrast to the performance of the two filters on the \\wmap\\ maps, where the MF was much more sensitive than the UF. Differences in performances are to be expected, as the two experiments have different characteristics. As Planck has much lower noise, the tSZ contamination now plays a more major role in setting the bias and errors.~\\planck\\ increased frequency coverage and resolution, however, enables the filters (UF in particular) to mitigate the effects of the thermal SZ signal, especially when a large cluster sample is considered. \\end{enumerate} We have not considered here the effect of infrared point sources, correlations in the signal between CMB and clusters, large--scale (non bulk flow) peculiar velocities, and error determination in extracting optical depth from the survey in hand. Some of these could present further improvements to this initial study, which demonstrate the superior potentials \\planck\\ has in determining the bulk flow. %\\begin{landscape} \\begin{deluxetable}{cccccccc} \\tablecaption{Summary of the results} \\tablewidth{0pc} \\tablecolumns{8} \\tabletypesize{\\scriptsize} \\tablehead{ \\colhead{} & \\colhead{noise components} & \\multicolumn{2}{c}{RASS} & \\multicolumn{2}{c}{\\planck} & \\multicolumn{2}{c}{EASS} \\\\ \\colhead{} & \\colhead{} & \\colhead{MF} & \\colhead{UF} & \\colhead{MF} & \\colhead{UF} & \\colhead{MF} & \\colhead{UF} \\\\ } \\startdata $V^{500}_{\\rm rec}$ & noise + CMB + thermal SZ bias &$1085\\pm82$ &$720\\pm94$ & $530\\pm126$& $516\\pm73$ & $507\\pm33$& $500\\pm24$ \\\\ $\\left \\langle v_x \\right \\rangle$ (75 km/s)& & $129\\pm94$ & $131\\pm138$ & $108\\pm133$ & $88\\pm68$ &$80\\pm40$ & $75\\pm29$ \\\\ $\\left \\langle v_y \\right \\rangle$ (-426 km/s)& & $-75\\pm105$ & $-434\\pm133$ & $-432\\pm123$ & $-431\\pm73$ & $-427\\pm37$& $-428\\pm26$\\\\ $\\left \\langle v_z \\right \\rangle$ (250 km/s)& & $1065\\pm82$ & $526\\pm66$ & $ 240\\pm91$ & $255\\pm49$ & $256\\pm28$& $243\\pm19$\\\\ $\\Delta\\alpha^{500}_{95}$ & &$62^{\\circ}$ & $34^{\\circ}$ & $34^{\\circ}$ & $14^{\\circ}$ & $9^{\\circ}$& $4^{\\circ}$ \\\\ $\\sigma_{\\log V,\\rm bin1}$ & & 0.71& 0.40& 0.12& 0.09 & 0.09& 0.09 \\\\ $\\sigma_{\\log V,\\rm bin2}$ & & 0.28& 0.14& 0.09& 0.04 & 0.02& 0.01 \\\\ $\\sigma_{\\log V,\\rm bin3}$ & & 0.08& 0.05& 0.04& 0.02 & 0.01& 0.01 \\\\ $\\sigma_{\\log V,\\rm bin4}$ & & 0.02& 0.03& 0.03& 0.02 & 0.01& 0.01 \\\\ $\\sigma_{\\log V, \\rm whole}$ & & 0.42& 0.24& 0.08& 0.06 & 0.05& 0.05 \\\\ \\hline % $V^{500}_{\\rm rec}$ & only noise + CMB & $515\\pm47$ & $521\\pm83$& $514\\pm30$& $497\\pm42$ & 479\\pm17$& $484\\pm26$ \\\\ % $\\left \\langle v_x \\right \\rangle$ (75 km/s)& & $69\\pm55$ & $67\\pm87$ & $79\\pm32$ & $75\\pm44$ & $73\\pm18$ & $77\\pm22$ \\\\ % $\\left \\langle v_y \\right \\rangle$ (-426 km/s)& & $-437pm53$ & $-436\\pm101$ & $-441\\pm33$ & $-419\\pm43$ & $-408\\pm17$& $-414\\pm27$ \\\\ % $\\left \\langle v_z \\right \\rangle$ (250 km/s)& & $255\\pm38$ & $251\\pm49$ & $ 248\\pm24$ & $249\\pm33$ & $239\\pm14$& $235\\pm16$ \\\\ % $\\Delta\\alpha^{500}_{95}$ & &$1^{\\circ}$ &$1^{\\circ}$ & $<1^{\\circ}$ & $<1^{\\circ}$ & $<1^{\\circ}$& $1^{\\circ}$ \\\\ % $\\sigma_{\\log V,\\rm bin1}$ & & 0.14& 0.18& 0.09& 0.07 & 0.04& 0.06 \\\\ % $\\sigma_{\\log V,\\rm bin2}$ & & 0.03& 0.04& 0.03& 0.04 & 0.02& 0.02 \\\\ % $\\sigma_{\\log V,\\rm bin3}$ & & 0.02& 0.02& 0.03& 0.02 & 0.01& 0.02 \\\\ % $\\sigma_{\\log V,\\rm bin4}$ & & 0.03& 0.01& 0.03& 0.02 & 0.05& 0.01 \\\\ % $\\sigma_{\\log V, \\rm whole}$ & & 0.084& 0.11& 0.06& 0.04 & 0.03& 0.03 \\\\ % \\hline $V^{500}_{\\rm rec} (\\left \\langle V_{\\rm rec} \\right \\rangle)$ & only thermal SZ bias & $1100\\pm60$ & $724\\pm82$ & $530\\pm115$ & $515\\pm63$ & $511\\pm30$& $504\\pm14$ \\\\ $\\left \\langle v_x \\right \\rangle$ (75 km/s)& & $125\\pm73$ & $142\\pm98$ & $82\\pm109$ & $82\\pm60$ & $104\\pm30$ & $101\\pm15$ \\\\ $\\left \\langle v_y \\right \\rangle$ (-426 km/s)& & $-71\\pm75$ & $-445\\pm110$ & $-437\\pm112$ & $-436\\pm63$ & $-427\\pm34$& $-428\\pm14$ \\\\ $\\left \\langle v_z \\right \\rangle$ (250 km/s)& & $1082\\pm63$ & $534\\pm46$ & $ 253\\pm86$ & $250\\pm46$ & $255\\pm29$& $244\\pm11$ \\\\ $\\Delta\\alpha^{500}_{95}$ & &$62^{\\circ}$ & $33^{\\circ}$& $32^{\\circ}$ & $14^{\\circ}$ & $9^{\\circ}$& $3^{\\circ}$ \\\\ $\\sigma_{\\log V,\\rm bin1}$ & & 0.68& 0.40& 0.12& 0.10 & 0.07& 0.04 \\\\ $\\sigma_{\\log V,\\rm bin2}$ & & 0.27& 0.15& 0.06& 0.03 & 0.03& 0.02 \\\\ $\\sigma_{\\log V,\\rm bin3}$ & & 0.09& 0.06& 0.04& 0.02 & 0.02& 0.01 \\\\ $\\sigma_{\\log V,\\rm bin4}$ & & 0.03& 0.03& 0.03& 0.02 & 0.01& 0.01 \\\\ $\\sigma_{\\log V, \\rm whole}$ & & 0.41& 0.24& 0.08& 0.06 & 0.05& 0.03 \\enddata \\tablecomments{This table lists the main results of the bulk flow measurements for clusters within $z\\le0.5$ in the three cluster surveys considered and for maps containing different components of uncertainties. The first five quantities are based on 100 realizations with input velocity of 500 km/s at $l=280^{\\circ}$ and $b=30^{\\circ}$. $V^{500}_{\\rm rec}$ is the average value of the recovered velocity, $\\left \\langle v_x \\right \\rangle$ , $\\left \\langle v_x \\right \\rangle$ , $\\left \\langle v_x \\right \\rangle$ are the average of the individual directions. For ideal recovery, they should have values (75, -426, 250) km/s. $\\Delta\\alpha^{500}_{95}$ is the $95\\%$ upper limit to the angle error. The last five quantities $\\sigma_{\\log V, \\rm bin i}$ are the deviation parameters as defined in equation~\\ref{eq:dispersionvv}, where bin i refers to the binned velocity range considered in the calculation: bin 1= 100--500 km/s, bin 2= 500--1000 km/s, bin 3= 1000--5000 km/s, bin 4= 5000--10000 km/s. The whole range is for velocities 100--10000 km/s. } \\label{t:summary} \\end{deluxetable} %\\end{landscape}" }, "1101/1101.2093_arXiv.txt": { "abstract": "{Submillimeter galaxies are a population of dusty star-forming galaxies at high redshift. Measuring their properties will help relate them to other types of galaxies, both at high and low redshift. This is needed in order to understand the formation and evolution of galaxies. }{The aim is use gravitational lensing by galaxy clusters to probe the faint and abundant submillimeter galaxy population down to a lower flux density level than what can be achieved in blank-field observations.}{We use the LABOCA bolometer camera on the APEX telescope to observe five cluster of galaxies at a wavelength of 870 $\\mu$m. The final maps have an angular resolution of 27.5\\arcsec and a point source noise level of 1.2--2.2 mJy. We model the mass distribution in the clusters as superpositions of spherical NFW halos and derive magnification maps that we use to calculate intrinsic flux densities as well as area-weighted number counts. We also use the positions of Spitzer MIPS 24~$\\mu \\mathrm{m}$~sources in four of the fields for a stacking analysis.}{We detected 37 submm sources, out of which 14 have not been previously reported. One source has a sub-mJy intrinsic flux density. The derived number counts are consistent with previous results, after correction for gravitational magnification and completeness levels. The stacking analysis reveals an intrinsic 870~$\\mu \\mathrm{m}$~signal of $390\\pm 27$~$\\mu$Jy at $14.5\\sigma$ significance. We study the $S_{\\mathrm{24 \\, {\\mu}m}}$ -- {$S_{\\mathrm{870 \\, {\\mu}m}}$}~relation by stacking on subsamples of the 24~$\\mu \\mathrm{m}$~sources and find a linear relation at $S_{\\mathrm{24 \\, {\\mu}m}}$$<300 \\, \\mu$Jy, followed by a flattening at higher 24~$\\mu \\mathrm{m}$~flux densities. The signal from the significantly detected sources in the maps accounts for 13\\% of the Extragalactic Background Light discovered by COBE, and the stacked signal accounts for 11\\%.} {} \\date{\\today} ", "introduction": "\\label{sec:introduction} Submillimeter galaxies form a population of high-redshift, dusty star-forming galaxies that are highly obscured in the visible and in the near-infrared, and have a spectral energy distribution (SED) that peaks in the submillimeter (submm) waveband (see e.g. \\cite{BlainSmail:2002aa} for a review). Most recent searches for submm galaxies have been based on surveys of blank sky with no known large-scale structure along the line-of-sight. These surveys have exploited large-format sensitive bolometer arrays (e.g. SCUBA on the James Clerk Maxwell Telescope (JCMT), \\citealt{CoppinChapin:2006aa}; LABOCA~on the Atacama Pathfinder EXperiment (APEX), \\citealt{WeisKovacs:2009ab}; AzTEC on ASTE: \\citealt{AustermannDunlop:2010aa,ScottYun:2010ab}; the South Pole Telescope (SPT): \\citealt{VieiraCrawford:2010aa}; MAMBO on the IRAM 30~m telescope: \\citealt{GreveIvison:2004aa}, % \\citealt{BertoldiCarilli:2007aa}). Those maps cover large areas at a nearly uniform noise level, leading to a simple selection function with a constant completeness across the field. The observations showed that source number counts increase steeply with decreasing flux density at mm and submm wavelengths \\cite[e.g.][]{WeisKovacs:2009ab,PatanchonAde:2009ab}. In order to probe the faint (below a few mJy) population of submm galaxies, several authors have taken advantage of the gravitational magnification induced by massive clusters of galaxies (e.g. \\citealt{SmailIvison:1997aa,SmailIvison:2002aa,ChapmanSmail:2002aa,Knudsenvan-der-Werf:2005aa,KnudsenBarnard:2006aa,Knudsenvan-der-Werf:2008aa,JohanssonHorellou:2010ab,WardlowSmail:2010aa, RexAde:2009aa,EgamiRex:2010aa}). A large magnification, produced for example when a source lies close to a critical line of the lens, may make it possible to detect a source with an intrinsic flux density much lower than the formal root mean square of the noise of the observation. This is the only method of detecting such dim sources directly. Cluster field observations have sensitivities that vary across the map, as magnified sources are ``lifted'' above the detection limit, and the selection functions are therefore more complicated. The most comprehensive study to date of submm galaxies behind lensing clusters is that of \\cite{Knudsenvan-der-Werf:2008aa}, who analyzed SCUBA data from 12 galaxy clusters and one blank field, resulting in an effective surveyed area of 71.5 arcmin$^2$ on the sky, but an area in the source plane almost twice as small. Seven sources with sub-mJy fluxes were detected. The sources revealed by gravitational lensing are prime targets for observations across the electromagnetic spectrum. \\cite{SwinbankSmail:2010aa} discovered a very bright submm source, situated at $z=2.33$, with flux density {$S_{\\mathrm{870 \\, {\\mu}m}}$}$\\sim$106~mJy, and molecular line observations showed that the amount of molecular gas is similar to that in local ultra-luminous infra-red galaxies (ULIRGs, \\citealt{DanielsonSwinbank:2010aa}). The $z\\sim 3.4$ submm source studied by \\cite{IkarashiKohno:2010aa} and discovered through use of AzTEC at 1.1~mm \\citep{WilsonAustermann:2008aa} has a 880~$\\mu \\mathrm{m}$~flux density measured by the Submillimeter Array (SMA) of $\\sim 73$~mJy and seems to be a ULIRG as well. On the other hand, the $z=2.79$ galaxy behind the Bullet Cluster \\citep{GonzalezPapovich:2010aa}, with a flux density of about 48~mJy at 870~$\\mu \\mathrm{m}$, is more representative of the normal galaxy population with an intrinsic far-infrared luminosity of a few times 10$^{11} L_\\odot$ \\citep{2008MNRAS.390.1061W,GonzalezClowe:2009aa,JohanssonHorellou:2010ab}. Large surveys in the mm (SPT) and the far-infrared (Herschel) are discovering bright lensed submm galaxies \\citep{VieiraCrawford:2010aa,NegrelloHopwood:2010aa}. Another way to probe the faint part of the submm galaxy population is to perform a stacking analysis using known positions obtained from complementary observations at another wavelength. \\cite{2006A&A...451..417D} % used the positions of sources detected with Spitzer Space Telescope at 24~$\\mu$m to measure the contribution of those sources to the 70 and 160~$\\mu$m far-infrared background, gaining up to one order of magnitude in depth. Greve et al. (2010) carried out a stacking analysis of the LABOCA submm map of the Extended Chandra Deep Field (ECDF) using a large sample of near-infrared detected galaxies. \\begin{table*}[t!] \\centering \\caption{Observed cluster fields. } \\begin{tabular}[t]{l l l c c c c } \\hline \\hline Target & $\\alpha$\\tablefootmark{a} [J2000] & $\\delta$\\tablefootmark{a} [J2000] & $z$ & rms\\tablefootmark{b} & $\\Omega$\\tablefootmark{c} \\\\ & [h~~~~m~~~~s] & [\\degr~~~~\\arcmin~~~~\\arcsec] & & [$\\mathrm{mJy\\, beam}^{-1}$] & [$\\mathrm{arcmin}^2$]\\\\ \\hline {Abell~2163} & 16 15 45.1 & $-$06 08 31 & 0.203 & 2.2 & 150 \\\\ Bullet Cluster\\tablefootmark{1} & 06 58 29.2 & $-$55 56 45 & 0.296 & 1.2 & 220 \\\\ Abell~2744\\tablefootmark{2} & 00 14 15.0 & $-$30 22 60 & 0.308 & 1.5 & 220 \\\\ AC~114\\tablefootmark{3} & 22 58 52.3 & $-$34 46 55 & 0.312 & 1.2 & 130 \\\\ MS~1054-03\\tablefootmark{4} & 10 57 00.2 & $-$03 37 27 & 0.823 & 1.6 & 200 \\\\ \\hline \\end{tabular} \\tablefoot{\\tablefoottext{a} Central coordinates of the $2\\times 2$ raster square scanning pattern; these positions differ slightly from the central X-ray positions. \\tablefoottext{b} The noise level measured in the central 10 arcminutes of each map, as described in Sect.~\\ref{sec:survey-completeness}. \\tablefoottext{c} Extent of the LABOCA~maps. \\\\ \\tablefoottext{1} Alternative name 1E~0657-56. Project's observing identification (Obs. ID): O-079.F-9304A-2007, E-380.A-3036A-2007. \\\\ \\tablefoottext{2} Alternative name AC 118. Obs. ID O-081.F-9319A-2008. \\\\ \\tablefoottext{3} Alternative name Abell~S~1077, Obs. ID E-081.A-0451A-2008, E-078.F-9032A-2007. \\\\ \\tablefoottext{4} Obs. ID O-083.F-9300A-2009} \\label{tab:targ} \\end{table*} In this paper, we extend the analysis of submm sources behind the Bullet Cluster~recently presented in \\cite{JohanssonHorellou:2010ab} by four additional galaxy cluster fields observed with the LABOCA~receiver on the APEX telescope. The deep observations allow us to detect submm galaxies with observed flux densities above $\\sim 4.5$~mJy, while the gravitational magnification reveals galaxies with intrinsic sub-mJy flux densities. We derive the magnification of the foreground clusters by using the lens equation for clusters modeled as a superposition of Navarro, Frenk and White (NFW, 1997) mass density profiles whose parameters are inferred from published papers on the selected clusters. From the magnification maps, we calculate intrinsic flux densities and derive submm number counts for the entire survey. We carry out a stacking analysis on 24~$\\mu \\mathrm{m}$~detected sources in the fields to probe the correlation between submm and mid-infrared emission and detect stacked 870~$\\mu \\mathrm{m}$~observed flux densities of {$S_{\\mathrm{870 \\, {\\mu}m}}$}~$<800$~$\\mu$Jy for sources that are undetected individually in the maps. This paper is organized as follows: in Sect.~\\ref{sec:observations} we describe the submm observations and data reduction and the Spitzer MIPS 24~$\\mu \\mathrm{m}$~archival data; in Sect.~\\ref{sec:results} we present the resulting maps. In Sect.~\\ref{sec:cluster-lens-models} we discuss the lensing models and the number counts and in Sect.~\\ref{sec:stacking-analysis} we present a stacking analysis. Section~\\ref{sec:contribution-ebl} discusses the contribution of our submm signals to the Extragalactic Background Light discovered by COBE. The results are summarized in Sect.~\\ref{sec:conclusions}. Throughout the paper, we adopt the following cosmological parameters: a Hubble constant $H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$, a matter density parameter $\\Omega_0 = 0.27$, and a dark energy density parameter $\\Omega_{\\Lambda0} = 0.73$. The redshift $z=0.3$ where three of our clusters reside corresponds to an angular-diameter distance of 911~Mpc and a scale of 4.42~kpc/arcsec. $z=2.2$, the median redshift of known submm galaxies, corresponds to an angular-diameter distance of 1728~Mpc and a scale of 8.38~kpc/arcsec\\footnote{We used Ned Wright's cosmology calculator \\citep{Wright:2006aa} available at \\texttt{http://www.astro.ucla.edu/{\\textasciitilde}wright/cosmocalc.html}.}. ", "conclusions": "\\label{sec:conclusions} We used the LABOCA~receiver on APEX to carry out a submm survey of five clusters of galaxies. The clusters act as gravitational lenses and magnify background sources. The main results of the survey are summarized below. \\begin{enumerate} \\item We discovered 37 submm sources, out of which 14 are new submm detections. \\item We modeled the galaxy clusters as superpositions of spherical NFW halos and generated magnification maps for the five clusters. \\item The magnification maps were used to correct for the gravitational lensing and to obtain the intrinsic flux densities of the detected sources. \\item We constructed number counts taking into account both the gravitational lensing and the varying completeness level. The number counts probe the sub-mJy level and are consistent with previous work within the uncertainties. \\item We performed a stacking analysis in the LABOCA~maps on positions of detected 24~$\\mu \\mathrm{m}$~sources in the fields. The stacking yields $>4.9 \\sigma$ detections in all fields with MIPS coverage, reaching noise levels below $100$~$\\mu$Jy, more than an order of magnitude deeper than the individual maps. \\item By dividing the 24~$\\mu \\mathrm{m}$~catalog in MS~1054-03, AC~114~and Abell~2744~into six equal halves, we find a linear relation between $S_{\\mathrm{24 \\, {\\mu}m}}$~and {$S_{\\mathrm{870 \\, {\\mu}m}}$}~at low 24~$\\mu \\mathrm{m}$~fluxes, followed by a flattening of the relation at $S_{\\mathrm{24 \\, {\\mu}m}} \\sim 300$~$\\mu$Jy. This behavior can be explained if the low MIPS fluxes trace star formation while the higher values are dominated by AGN heating. \\item The observations reveal a total of $\\sim$24\\% of the infrared extragalactic background light, where $\\sim$13\\% comes from the significant submm sources and $\\sim$11\\% comes from the stacked signal. \\end{enumerate}" }, "1101/1101.1745_arXiv.txt": { "abstract": "% {The formulation of the radio interferometer measurement equation (RIME) for a generic radio telescope by Hamaker et al. has provided us with an elegant mathematical apparatus for better understanding, simulation and calibration of existing and future instruments. The calibration of the new radio telescopes (LOFAR, SKA) would be unthinkable without the RIME formalism, and new software to exploit it.}% {The MeqTrees software system is designed to implement numerical models, and to solve for arbitrary subsets of their parameters. It may be applied to many problems, but was originally geared towards implementing Measurement Equations in radio astronomy for the purposes of simulation and calibration. The technical goal of MeqTrees is to provide a tool for rapid implementation of such models, while offering performance comparable to hand-written code. We are also pursuing the wider goal of increasing the rate of evolution of radio astronomical software, by offering a tool that facilitates rapid experimentation, and exchange of ideas (and scripts).}% {MeqTrees is implemented as a Python-based front-end called the meqbrowser, and an efficient (C++-based) computational back-end called the meqserver. Numerical models are defined on the front-end via a Python-based Tree Definition Language (TDL), then rapidly executed on the back-end. The use of TDL facilitates an extremely short turn-around time (hours rather than weeks or months) for experimentation with new ideas. This is also helped by unprecedented visualization capabilities for all final and intermediate results. A flexible data model and a number of important optimizations in the back-end ensures that the numerical performance is comparable to that of hand-written code.}% {MeqTrees is already widely used as the simulation tool for new instruments (LOFAR, SKA) and technologies (focal plane arrays). It has demonstrated that it can achieve a noise-limited dynamic range in excess of a million, on WSRT data. It is the only package that is specifically designed to handle what we propose to call {{\\em third-generation}} calibration (3GC), which is needed for the new generation of giant radio telescopes, but can also improve the calibration of existing instruments.}% {} ", "introduction": "} The MeqTrees software system has been designed to implement an arbitrary Measurement Equation (i.e. a numerical model of an instrument and/or process), and to solve for arbitrary subsets of its parameters. In this paper we will concentrate on the simulation and calibration of data taken with radio telescopes. After all, that is the subject for which MeqTrees was developed originally, and for which it is most urgent. It is also an excellent subject for demonstrating the special capabilities of MeqTrees. Until recently, radio interferometers like WSRT, VLA, ATCA, GMRT were designed so that they could be approximated by a relatively simple instrumental model. Their {\\em stations}\\footnote{Throughout this paper, we will use the generic term {\\em station} for an element of an interferometer array. A station can be a parabolic dish or an aperture array, or something more exotic like a parabolic cylinder. Each station has two output signals, one for each polarization.} were carefully designed so that the shape of their spatial response beams could be assumed to be `identical' at all times. In addition, the instrumental error associated with a station can be represented in this model by only two complex gain factors (one for each polarization), which may vary in time and frequency. Because their stations are parabolic dishes that are pointed towards the observed field, instrumental polarization effects can be treated (or not) as small `leakage' terms. Before 1980, first-generation calibration (1GC) was based on `open-loop' methods, making separate calibrator observations before and after, and relying on instrumental stability in between. The result was a dynamic range\\footnote{Dynamic range is the ratio between the flux of the brightest source in the field, and either the thermal noise or the calibration artifacts, whichever is higher.} of about 100:1. However modest, this was enough for a plethora of important discoveries. Around 1980, the invention of {\\em self-calibration} \\citep[see also summary by \\citealt{Ekers:Selfcal}]{Cornwell:selfcal} ushered in the era of second-generation calibration (2GC). Selfcal is a ``closed-loop'' method which continuously estimates the complex station gain factors with the help of one or more bright sources in the field of view. In this process, the utilized Sky Model is improved also. Selfcal has been spectacularly successful, and has led to a blossoming of techniques, software packages, and beautiful results. It allows astronomers to achieve dynamic ranges in excess of $10^4-10^5$ as a matter of routine, depending on how well the instrument approximates its simplified model. Record dynamic ranges of well over $10^6$ have been achieved at the WSRT by \\citet{deBruyn:million,deBruyn:3C147} (see also Sect. \\ref{sec:Results}). Radio astronomy is currently going through a remarkable worldwide burst of building new telescopes and upgrading existing ones\\footnote{An important incentive is the preparation for the building of the multi-billion Euro Square Kilometer Array (SKA) later in this decade.}. These instruments present a new, two-pronged calibration challenge. On the one hand, they are much more sensitive, so more subtle instrumental effects will have to be taken into account to reach the thermal noise. On the other hand, the use of new technology like phased arrays complicates the instrumental model. Therefore, what we propose to call {\\em third-generation calibration} (3GC) will require a more complex, able, and a general form of selfcal. In 1996, \\citeauthor{ME1} developed a formulation of an explicit radio interferometer measurement equation (RIME) for a generic radio telescope. Further work by \\citet{ME4} led to a fully $2\\times2$ matrix formulation of the RIME, which provides the mathematical underpinnings for 3GC. Without this full-polarization formalism, calibration of the new telescopes would be difficult. However, although the RIME is widely recognized as being correct, complete and universal, its actual adoption has been slow. This is caused to a large extent by the sustained success of the existing 2GC data reduction packages (AIPS, NEWSTAR, MIRIAD, DIFMAP). The ensuing low rate of evolution of calibration techniques could be a risk factor for the new telescopes. One of the most important challenges of 3GC is dealing with Direction-Dependent Effects (DDEs), i.e. instrumental effects that can no longer be assumed to be constant over the field of view. The most important DDEs are typically caused by the ionosphere (mostly phase and Farady rotation), and by station beamshapes that differ substantially from each other, and/or vary individually in frequency and time. Tackling DDEs implies that one has to solve for a much larger number of RIME parameters than before. Besides the practical problem of extra processing (which may well turn out to be a major bottleneck, but will not be discussed here), this raises some more fundamental issues about whether there is enough information available for 3GC. And, last but not least, methods are needed to correct for DDEs once they are known, which is non-trivial. At this moment, it is not yet clear how some of the new generation of telescopes will be calibrated to the necessary precision. The MeqTrees software system is a tool that can play a role in building that understanding on several levels. First of all, it is firmly based on the explicit RIME. Secondly, it has many built-in features for {\\em generalized selfcal}, such as allowing for arbitrary RIMEs with arbitrary parameterizations, and solving for arbitrary subsets of RIME parameters, including source parameters. These parameterizations may be experimented with rather rapidly, since the art of modelling (in Python) is separated from the complex and efficient numerical machinery (in C++) ``under the hood''. Rapid progress is also greately helped by the many possibilities for visualization of intermediate results, enabling one to easily see what is actually going on. The heart of this paper, Sects.~\\ref{sec:forest}--\\ref{sec:performance}, is a detailed description of how MeqTrees works. Sect.~\\ref{sec:RIME} gives a brief description of the RIME, explains why it is such a powerful formalism, and shows how it pertains to MeqTrees. In Sect.~\\ref{sec:Results}, we present some recent results that give a taste of what MeqTrees can do. ", "conclusions": "Conclusions} MeqTrees raises the art of instrumental modelling to the level where user-developers can concentrate on the physics of the problem, while the complex numerical machinery, e.g. for solving for arbitrary subsets of parameters, is hidden ``under the hood''. In the special case of radio astronomy, the correct treatment of various instrumental effects is greatly facilitated by the elegant matrix formalism of the Measurement Equation of a generic radio telescope (RIME). The latter is well on its way to becoming the new Common Language of radio astronomy. The present collection of node classes offers basic functionality, with bias towards radio astronomy (see Appendix \\ref{sec:Available-MeqNode-classes}). There are various TDL (Python) frameworks to help the user in building complex trees; these in fact evolve much more rapidly than the binary release cycle. We also intend to offer a MeqWizard tool to help both novices and experts to find their way in the multiverse of possibilities. The MeqTrees kernel is robust and efficient, and has been tested thoroughly on real data (see Sect. \\ref{sec:Results}.) MeqTrees is (slowly) beginning to find its place, propelled by the increasingly urgent need for 3GC simulation and calibration software for radio astronomy. It plays its designated role as pathfinder for LOFAR calibration, as illustrated by impressive all-sky LOFAR images \\citep{Yatawatta:lofar}. It has been used as the main education tool in several international workshops, to train the the new generation of radio astronomers in the use of the RIME and 3GC. MeqTrees is freely distributed under the terms of the GNU General Public License. A stable binary release (version 1.1.1 at time of writing) is available. This is shipped as binary packages for the major Linux distros, so installation is relatively painless (while users of unsupported platforms always have the option of building from source.) A Mac OSX version has been tested, but is not (yet) part of the binary release. MeqTrees is natively multi-threaded to take full advantage of multi-core machines common today. An experimental MPI-based cluster version was developed jointly with Oxford Astrophysics and OeRC, and is currently being tested by Tony Willis at DRAO (Canada.) For further information on downloading and installing the software, please refer to the MeqTrees Wiki: {\\tt http://www.astron.nl/meqwiki}. You can also join the MeqTrees forum hosted at UCL: {\\tt https://great08.projects.phys.ucl.ac.uk/meqtrees/}." }, "1101/1101.1868_arXiv.txt": { "abstract": "We report the clearest detection to date of dusty torus signatures in a Weak-Line Radio Galaxy (WLRG). The deep Spitzer InfraRed Spectrograph (IRS) rest-frame mid-infrared (MIR) spectrum of the WLRG PKS 0043-42 (z=0.116) shows a clear spectral turnover at $\\lambda\\ga$ 20 \\micron~suggestive of warm dust, as well as a 9.7 \\micron~silicate absorption feature. In addition, the hard X-ray results, based on Chandra data, strongly support a picture in which PKS 0043$-$42 has a torus and accretion disc more typical of Strong-Line Radio Galaxies (SLRGs). The MIR and X-ray spectra are markedly different from those of other WLRGs at similar redshifts, and here we show that the former can be successfully fitted with clumpy torus models with parameters characteristic of Type-2 AGN tori: close to edge-on ($i$=74\\degr) and relatively broad ($\\sigma$=60\\degr), with an outer radius of 2 pc, N$_H$=1.6$\\pm^{0.2}_{0.1}\\times$10$^{23}~cm^{-2}$, and AGN bolometric luminosity L$_{bol}^{AGN}$ = 1.6$\\pm^{0.2}_{0.1}\\times$10$^{44}~erg~s^{-1}$. The presence of a compact torus in PKS 0043-42 provides evidence that this WLRG is fuelled by cold, rather than hot, gas accretion. We suggest that WLRGs are a diverse population, and PKS 0043-42 may represent a type of radio galaxy in which the AGN activity has been recently re-triggered as a consequence of intermittent gas supply, or in which the covering factor of the Narrow-Line Region (NLR) clouds is relatively low. ", "introduction": "\\label{intro} Most powerful radio galaxies (PRGs) can be classified as Narrow- and Broad-Line Radio Galaxies/quasars (NLRGs, BLRGs/QSOs) on the basis of their spectral features. The unified model for active galactic nuclei (AGN; \\citealt{Antonucci93,Urry95}), proposes the existence of a pc-scale obscuring toroidal structure to account for the observed differences between their spectra. However, there is also a third class, the WLRGs (also known as Low-Excitation Galaxies; LEGs), whose optical spectra are dominated by the stellar continua of the host galaxies \\citep{Laing94,Tadhunter98} without the prominent emission lines characteristic of NLRGs and BLRGs (which can be grouped as SLRGs, or alternatively, High-Excitation Galaxies; HEGs). The nature of the physical parameters which lead to the division between SLRGs and WLRGs still remains unclear. Previous studies failed to explain the differences between the two types as due to different or time-varying accretion rates \\citep{Ghisellini01}. An alternative interpretation is that SLRGs are powered by cold gas accretion, while WLRGs are fuelled by accretion of hot gas provided by the reservoir of their X-ray gaseous coronae \\citep{Allen06,Best06,Hardcastle07,Balmaverde08,Buttiglione10}. The high temperatures of the hot gas would prevent the formation of the ``cold'' structures (e.g. the Broad-Line Region and the torus). This hypothesis explains the fact that WLRGs do not normally show the high level of X-ray absorption expected from a molecular torus \\citep{Hardcastle09}. SLRGs are almost invariably associated with Fanaroff-Riley II sources (FRII). On the other hand, while most WLRGs show Fanaroff-Riley I (FRI) radio morphologies, some are classified as FRIIs \\citep{Laing94}. Evidence for nuclear warm dust emission has been found in the MIR for FRII/SLRGs \\citep{Ogle06,Wolk10} and, to a lesser extent, for a few FRI and FRII WLRGs \\citep{Leipski09,Wolk10}. However, none of the WLRGs so far studied have provided definitive evidence for a warm compact obscuring region e.g. in the form of a spectral turnover at $\\sim$20 \\micron~and a clear silicate absorption feature. This lack of conclusive detection of torus signatures in previous studies may be due to contamination of the MIR data by stellar photospheric and starburst-heated dust components, which generally produces a broad minimum in the 8-10 \\micron~range and can be confused with a silicate absorption feature. Strong PAH (in particular the 8.2 and 11.3 \\micron~features) and non-thermal emission from the synchrotron-emitting core sources can also mask the torus signatures. Based on extrapolation of the radio core data, the latter emission appears particularly important in WLRGs at MIR wavelengths \\citep{Dicken08,Wolk10}. All of these diluting components need to be accurately subtracted before the presence of warm dust can be deduced. At optical wavelengths the galaxy PKS 0043-42 is classified as a WLRG at redshift z=0.116 \\citep{Tadhunter93}, whereas at radio wavelengths it shows a clear FRII morphology \\citep{Morganti93,Morganti99}. According to \\citet{Tadhunter98}, the WLRGs are defined as a group of galaxies with [O III]$\\lambda$5007 emission line equivalent widths below 10\\AA. PKS 0043-42 not only fullfills this criterion, but also has an [O III]$\\lambda$5007 luminosity a factor of 10 lower than those of the general population of SLRGs at similar redshifts and radio powers (L$_{[O III]}$=5$\\times$10$^{40}$ erg~s$^{-1}$; \\citealt{Tadhunter98,Dicken09}), as well as line ratios indicating a low ionization state (e.g., [O III]$\\lambda$5007/[O II]$\\lambda$3727 $<$ 1; \\citealt{Tadhunter98,Lewis03}). The galaxy is at the center of a group or cluster, and from NIR imaging, \\citet{Inskip10} reported a central isophotal twist and an excess of emission along the NS direction. They claimed that both are likely to be associated either with the presence of a dust lane, or with the apparent interaction with the companion object $\\sim$15~kpc to the North. However, the deep optical image presented in \\citet{Ramos10} does not reveal any dust features in PKS 0043-42, but appears instead to confirm the presence of the bridge with the companion galaxy \\citep{Ramos10}. In this study we report the detection of torus signatures in the MIR and X-ray spectra of the WLRG PKS 0043-42. In the following sections we describe the observations, the SED modelling, the main results and their implications. Throughhout this paper we assume a cosmology with H$_0$=73 km~s$^{-1}$~Mpc$^{-1}$; $\\Omega_m$=0.27, and $\\Omega_{\\Lambda}$=0.73, which provides a luminosity distance of 516 Mpc for PKS 0043-42 (from the NASA/IPAC Extragalactic Database; NED). ", "conclusions": "" }, "1101/1101.4026_arXiv.txt": { "abstract": "Cool core clusters of galaxies require strong feedback from their central AGN to offset cooling. We present a study of strong cool core, highly-luminous (most with $\\Lx\\ge~10^{45}~\\ergps$), clusters of galaxies in which the mean central AGN jet power must be very high yet no central point X-ray source is detected. Using the unique spatial resolution of $Chandra$, a sample of 13 clusters is analysed, including A1835, A2204, and one of the most massive cool core clusters, RXCJ1504.1-0248. All of the central galaxies host a radio source, indicating an active nucleus, and no obvious X-ray point source. For all clusters in the sample, the nucleus has an X-ray bolometric luminosity below 2 per cent of that of the entire cluster. We investigate how these clusters can have such strong X-ray luminosities, short radiative cooling-times of the inner intracluster gas requiring strong energy feedback to counterbalance that cooling, and yet have such radiatively-inefficient cores with, on average, $\\Lkin$/$\\Lnuc$ exceeding 200. Explanations of this puzzle carry significant implications for the origin and operation of jets, as well as on establishing the importance of kinetic feedback for the evolution of galaxies and their surrounding medium. ", "introduction": "Clusters of galaxies with steeply rising X-ray surface brightness profiles are known as cool core clusters, and require strong feedback from their central AGN to offset cooling of the intracluster medium (ICM). For highly-luminous cool core clusters (with $\\Lx\\ge10^{45}\\ergps$), the central AGN must be injecting on average $10^{45}\\ergps$ into the surrounding medium. We do not know the black hole (BH) mass for most of these AGN but we expect that they lie between $10^9$ and $10^{10}\\Msun$, based on the few reliable measurements in the literature ({\\rm e. g.} M87, \\cite[Macchetto et al. 1997]{Macchetto et al. 1997}). Given that these objects are in an exceptional environment and are active continuously, it is not clear that the standard $M_\\mathrm{BH}-\\sigma$ or $M_\\mathrm{BH}-M_\\mathrm{K}$ relations are relevant for them. For $M_\\mathrm{BH}\\sim10^9M_\\odot$, this means that these black holes must be operating at high enough Eddington rates that they should be radiatively efficient (see Fig. \\ref{fig1}). We would therefore expect to see an X-ray point source. We present a sample of strong cool core, highly-luminous clusters, for which there is {\\emph no} evidence of a nuclear X-ray point source in the $Chandra$ images. Using these images, we derive upper limits of the nuclear luminosities with the web interface of {\\sc pimms} (\\cite[Mukai 1993]{Mukai 1993}), which converts a count rate into an expected flux. We also investigate whether there is a hidden power law in the X-ray spectra, but find no such evidence in any of our objects. Finally, we calculate the energy that must be injected by the AGN in order to offset the cooling (\\Loutflow) of the ICM, and compare it with the nuclear luminosity of the AGN. \\begin{figure}[t] \\begin{center} \\includegraphics[width=3.in]{Fig1.eps} \\caption{Sketch of black hole energy release as a function of mass accretion (\\cite[Churazov et al. 2005]{Churazov et al. 2005}). The energy in radiation dominates at high accretion rates. If a black hole is releasing $10^{45}\\ergps$, then for $M_\\mathrm{BH}\\sim10^9M_\\odot$ the power exceeds $10^{-2}\\LEdd$, indicating that $\\Pradiation\\ge\\Poutflow$, i.e. we should see a nuclear point source. For $M_\\mathrm{BH}\\ge10^{10}M_\\odot$, the power exceeds $10^{-3}\\LEdd$. Here, $\\Poutflow$ could dominate over \\Pradiation, if the black hole is in a low accretion state. } \\label{fig1} \\end{center} \\end{figure} ", "conclusions": "We have identified a population of objects which require powerful jets to be present but have no X-ray detectable nucleus. We have also identified a range of possible explanations, each of which carries significant implication for the origin and operation of jets. The black holes may be ultramassive ($M_\\mathrm{BH}\\gg10^9M_\\odot$), or have very high spin, or be highly obscured. They may also be mostly off, yet unobservable when they are switched on, or have highly radiatively inefficient jets." }, "1101/1101.0718_arXiv.txt": { "abstract": "{Chamaeleon I is the most active region in terms of star formation in the Chamaeleon molecular cloud complex. Although it is one of the nearest low-mass star forming regions, its population of prestellar and protostellar cores is not known and a controversy exists concerning its history of star formation.} {Our goal is to search for prestellar and protostellar cores and characterize the earliest stages of star formation in this cloud.} {We used the bolometer array LABOCA at the APEX telescope to map the cloud in dust continuum emission at 870~$\\mu$m with a high sensitivity. This deep, unbiased survey was performed based on an extinction map derived from 2MASS data. The 870~$\\mu$m map is compared with the extinction map and C$^{18}$O observations, and decomposed with a multiresolution algorithm. The extracted sources are analysed by carefully taking into account the spatial filtering inherent in the data reduction process. A search for associations with young stellar objects is performed using \\textit{Spitzer} data and the SIMBAD database.} {Most of the detected 870~$\\mu$m emission is distributed in five filaments. We identify 59 starless cores, one candidate first hydrostatic core, and 21 sources associated with more evolved young stellar objects. The estimated 90$\\%$ completeness limit of our survey is 0.22~M$_\\odot$ for the starless cores. The latter are only found above a visual extinction threshold of 5 mag. They are less dense than those detected in other nearby molecular clouds by a factor of a few on average, maybe because of the better sensitivity of our survey. The core mass distribution is consistent with the IMF at the high-mass end but is overpopulated at the low-mass end. In addition, at most 17$\\%$ of the cores have a mass larger than the critical Bonnor-Ebert mass. Both results suggest that a large fraction of the starless cores may not be prestellar in nature. Based on the census of prestellar cores, Class~0 protostars, and more evolved young stellar objects, we conclude that the star formation rate has decreased with time in this cloud.} {The low fraction of candidate prestellar cores among the population of starless cores, the small number of Class 0 protostars, the high global star formation efficiency, the decrease of the star formation rate with time, and the low mass per unit length of the detected filaments all suggest that we may be witnessing the end of the star formation process in Chamaeleon~I.} ", "introduction": "\\label{s:intro} Large-scale, unbiased surveys in dust continuum emission have greatly improved our knowledge of the process of star formation in nearby, star-forming molecular clouds. During the past decade, (sub)mm dust continuum surveys established a close relationship between the prestellar core mass function (CMF) and the stellar initial mass function (IMF), suggesting that the IMF is already set at the early stage of cloud fragmentation \\citep[e.g.][]{Motte98,Testi98,Johnstone00}. In these submm surveys, prestellar cores are found only above a visual extinction of 5--7~mag, suggesting the existence of an extinction threshold for star formation to occur \\citep[see, e.g.,][, but also \\citeauthor{Onishi98} 1998 based on molecular line observations]{Johnstone04,Enoch06,Kirk06}. In the mid-infrared range, five of these nearby molecular clouds -- Chamaeleon~II, Lupus, Perseus, Serpens, and Ophiuchus -- were mapped with good angular resolution and high sensitivity with the \\textit{Spitzer} satellite in the frame of the c2d legacy project \\citep[\\textit{From molecular cores to planet forming disks}, see][]{Evans09}. This extensive project produced a deep census of young stellar objects. It delivered star formation efficiencies ranging from 3 to 6$\\%$ and \\citet{Evans09} estimate that these efficiencies could reach 15--30$\\%$ if star formation continues at current rates in these clouds. Star formation is found to be slow compared to the free-fall time and highly concentrated to regions of high extinction \\citep{Evans09}, consistent with the threshold mentioned above. Recently, the far-infrared/submm \\textit{Herschel} Space Observatory revealed an impressive network of parsec-scale filamentary structures in nearby molecular clouds in the frame of the Gould Belt Survey \\citep[][]{Andre10a}. This survey greatly increases the statistics of prestellar cores and confirms the close relationship between the CMF and the IMF. In the Aquila molecular cloud, several hundreds of prestellar cores are identified, and \\citet[][]{Andre10a} show that most of these prestellar cores are located in supercritical filaments. They confirm the existence of an extinction threshold for star formation at $A_V \\sim 7$~mag and propose that it may naturally come from the conditions required for a filament to become gravitationally unstable, collapse, and fragment into prestellar cores \\citep[see also][]{Andre10b}. This finding will be further investigated in the frame of this \\textit{Herschel} survey, but it can also be partly tested in other nearby clouds in the submm range, thanks to the advent of very sensitive cameras like LABOCA at the Atacama Pathfinder Experiment telescope (APEX). Chamaeleon I (Cha~I) is one of the nearest, low-mass star forming regions in the southern sky \\citep[150--160~pc,][; see Appendix~\\ref{ss:distance} for more details]{Whittet97,Knude98}. Together with Cha~II and III, it belongs to the Chamaeleon cloud complex whose population of T-Tauri stars has been well studied \\citep[see][ and references therein]{Luhman08c}. Cha~I contains nearly one order of magnitude more young stars than Cha~II, while Cha~III does not have any. Surveys performed in CO and its isotopologues showed that Cha~I contains a larger fraction of dense gas than Cha~II and III although the latter are somewhat more massive \\citep[][]{Boulanger98,Mizuno99,Mizuno01}. Finally, several indications of jets and outflows were found in Cha I \\citep[][]{Mattila89,Gomez04,Wang06,Belloche06}. Only three Herbig-Haro objects are known in Cha~II and none has been found in Cha~III \\citep[][]{Schwartz77}. Cha~I is therefore much more actively forming stars than Cha~II and III. A controversy exists, however, concerning the history of star formation in Cha~I. Based on the H-R diagram of the young stellar population known at that time (47 members with known spectral type), \\citet{Palla00} conclude that star formation in Cha~I began within the last 7 Myr and that its rate steadily increased until recently. They find similar results in several other clouds and draw the general conclusion that star formation in nearby associations and clusters started slowly some 10 Myr ago and \\textit{accelerated} until the present epoch. They speculate that this acceleration arises from contraction of the parent molecular cloud and results from star formation being a critical phenomenon occuring above some threshold density. However, based on a larger sample of known members, \\citet{Luhman07} recently came to the conclusion that star formation began 3--6~Myr ago in Cha~I and has continued to the present time at a \\textit{declining} rate. Since little is known about the earliest stages of star formation in Cha~I, investigating whether there is presently a population of condensations that are \\textit{prestellar}, i.e. bound to form stars, should provide strong constraints on these two competing scenarii for the star formation history in Cha~I. To unveil the present status of the earliest stages of star formation in Cha~I, we carried out a deep, unbiased dust continuum survey of this cloud at 870~$\\mu$m with the bolometer array LABOCA at APEX. The observations and data reduction are described in Sects.~\\ref{s:obs} and \\ref{s:reduction}, respectively. The maps are presented in Sect.~\\ref{s:results}. Section~\\ref{s:sourceid} explains the source extraction and the sources are analysed in Sect.~\\ref{s:analysis}. The implications are discussed in Sect.~\\ref{s:discussion}. Section~\\ref{s:conclusions} gives a summary of our results and conclusions. ", "conclusions": "" }, "1101/1101.2830_arXiv.txt": { "abstract": "In order to produce high dynamic range images in radio interferometry, bright extended sources need to be removed with minimal error. However, this is not a trivial task because the Fourier plane is sampled only at a finite number of points. The ensuing deconvolution problem has been solved in many ways, mainly by algorithms based on CLEAN. However, such algorithms that use image pixels as basis functions have inherent limitations and by using an orthonormal basis that span the whole image, we can overcome them \\cite{SBY10}. The construction of such an orthonormal basis involves fine tuning of many free parameters that define the basis functions. The optimal basis for a given problem (or a given extended source) is not guaranteed. In this paper, we discuss the use of generalized prolate spheroidal wave functions as a basis. Given the geometry (or the region of interest) of an extended source and the sampling points on the visibility plane, we can construct the optimal basis to model the source. Not only does this gives us the minimum number of basis functions required but also the artifacts outside the region of interest are minimized. ", "introduction": "Due to the increase in computing power, high dynamic range imaging in radio interferometry (only limited by calibration errors) is achievable and is essential to produce novel scientific results. One of the main obstacles to reach a high dynamic range is the fact that only a finite number of samples of the Fourier (visibility) plane data is observed. Moreover, in earth rotation synthesis, these sampling points are not on a regular grid. If the observed field of view contains bright extended sources, complete removal (deconvolution) of such sources to reveal the faint background remains a challenge. The commonly used algorithm for such problems is CLEAN \\cite{Hogbom}. However, as shown in \\cite{SBY10} (and references therein), the usage of a set of image pixels (clean components) in CLEAN based algorithms has limitations. An alternative approach to using clean components is to use an orthonormal basis to model bright extended structure in the observation. Once such a model is obtained, it can be subtracted from the visibilities to reveal the residual (or the fainter background). The most obvious method of deriving such a basis is to use the observed data itself \\cite{Levanda}. However, this has the drawback that the derivation of the basis can only be done once the complete observation is available (problematic for real time imaging etc.). On the other hand, we can adopt any arbitrary orthonormal basis to a given observation, completely independent of the data. For instance, in \\cite{SBY10} Gauss-Hermite polynomials (shapelets) were used to produce high dynamic range images. The drawback in such an approach is the selection of free parameters (such as the number of basis functions, and the scale) cannot be determined in an optimal fashion. It requires experience of the user as well as a trial and error approach to fine tune such a basis for a given observation. Due to the noise floor, any extended source has finite support \\cite{Slepian}, and the region of interest (ROI) or the support of a given extended source might not be optimal for a given arbitrary basis. In order to tackle the problem of finding the optimal basis for a given extended source, independent of the observed data, we select prolate spheroidal wave functions (PSWF) \\cite{Slepian61,Landau}. A similar problem has been solved in magnetic resonance imaging \\cite{Lindquist06} and we extend that result to radio interferometry in this paper. Unlike data derived basis functions, PSWF basis can be precomputed and can be reused for observations at different epochs. Furthermore, unlike shapelets, we suffer less from artifacts outside the ROI with minimal number of basis functions used. In fact, PSWF are already being used in radio astronomical imaging to construct a regular grid of sampling points in the Fourier plane \\cite{Brouw75}. We refer the reader to \\cite{Simons} for similar applications of PSWF in geoscience. Notation: We denote vectors in bold lowercase and matrices in bold uppercase. The matrix transpose, Hermitian, pseudoinverse are denoted by $(.)^T$, $(.)^H$ and $(.)^\\dagger$ respectively. The identity matrix is given by ${\\bf I}$. ", "conclusions": "We have presented the use of prolate spheroidal wave functions in radio interferometric image deconvolution and have demonstrated its feasibility by application to a real observation. We get results comparable with existing techniques, but with fewer basis functions and with less artifacts outside the ROI. Future work will focus on widefield imaging and reducing the computational cost." }, "1101/1101.0835_arXiv.txt": { "abstract": "We consider two classes of supersymmetric flipped SU(5) models with gravity mediated supersymmetry breaking such that the thermal neutralino relic abundance provides the observed dark matter density in the universe. We estimate the muon flux induced by neutrinos that arise from neutralino annihilations in the Sun and discuss prospects for detecting this flux in the IceCube/Deep Core experiment. We also provide comparisons with the corresponding fluxes in the constrained minimal supersymmetric standard model and non-universal Higgs models. Regions in the parameter space that can be explored by the IceCube/DeepCore experiment are identified. ", "introduction": "It is now widely accepted that approximately 23 $\\%$ of the Universe's energy density consists of non-baryonic cold dark matter \\cite{Komatsu:2008hk}. A large number of experiments consisting of direct, indirect and accelerator searches are currently underway all hoping to discover the underlying, presumably massive ($\\sim$ GeV -TeV), weakly interacting dark matter particle (WIMP). The lightest neutralino in supersymmetric models with conserved matter parity is a particularly attractive cold dark matter candidate and has attracted a great deal of attention. The direct detection searches have already yielded important constraints on the spin independent neutralino-nucleon cross sections in the constrained minimal supersymmetric standard model (CMSSM) and some related models (see \\cite{Feng:2010gw}, and references therein). Indirect WIMP searches rely on the capture and subsequent annihilation, say in the Sun's center, of relic dark matter particles. The neutralinos, in particular, can annihilate into the known SM particles, for example, $\\chi\\chi \\rightarrow \\tau^+\\tau^-$. The tau particles in turn produce energetic muon neutrinos which interact with the polar ice to produce muons which can be identified by the $\\rm km^3$ IceCube/Deep Core detector \\cite{Collaboration:2010hr}. Neutralinos in the galactic halo passing through a massive body like the Sun can get captured if they scatter off the nuclei with velocities smaller than the escape velocity. In the core of the Sun, where they eventually accumulate, these neutralinos can annihilate into known SM particles, for e.g., $\\chi\\chi \\rightarrow \\tau^+\\tau^-$ . These particles decay (e.g. $\\tau \\rightarrow \\nu_\\tau \\bar{\\nu}_\\mu\\mu$) and produce energetic muon neutrinos which can then be detected at IceCube after they interact with the polar ice and produce muons (for e.g. via processes like $\\nu_\\mu+N \\rightarrow \\mu^- +X$, N being the nucleon and X some hadronic system). We investigate the possibility of detecting these energetic neutrinos by estimating the flux of muons that they induce. In addition to the well studied CMSSM, we explore other well motivated models, namely, flipped SU(5), non-universal Higgs models (NUHM2), and flipped SU(5) with universal Higgs masses at $M_{GUT}$. The prospects for detecting this neutrino induced muon flux by the IceCube/DeepCore experiment is discussed. In this paper we are mainly interested in studying the implications of supersymmetric flipped SU(5) models for indirect dark matter WIMP searches, with the lightest neutralino being the dark matter candidate. Flipped SU(5) has several distinct features which are not easily replicated in other GUTs such as SU(5) and SO(10). For instance, the well-known doublet-triplet splitting problem is easily solved in flipped SU(5) \\cite{ant}. Primordial inflation with predictions for the cosmological parameters in good agreement with the 7 year WMAP data are readily obtained , which in turn, lead to testable predictions for proton decay \\cite{Rehman:2009yj}. Our paper is organized as follows. In section \\ref{framework}, following \\cite{Gogoladze:2009mc}, we briefly describe the two flipped SU(5) models under discussion. Consistent with the underlying gauge group (SU(5) x U(1)), both classes of models work with non-universal gaugino masses. Their difference stems from the non-universal soft scalar Higgs masses employed in one of the models. In section \\ref{flux-sd} we review the calculations of the conversion factors relating the muon flux and spin dependent (SD) cross section in the IceCube/Deep Core experiment. Section \\ref{procedure} contains a description of the experimental constraints and the scanning procedure employed to generate the benchmark points. Our predictions for the muon flux and SD cross section are presented in section \\ref{sec:results}, and the conclusions are summarized in section \\ref{conclusions}. ", "conclusions": "} We have considered indirect neutralino dark matter detection in two sets of supersymmetric flipped SU(5) models. These two sets of models have non-universal soft gaugino masses at $M_{GUT}$ that are related by the underlying $SU(5)\\times U(1)_X$ gauge symmetry. The supersymmetry breaking soft Higgs $\\rm {masses}^2$, associated with $H_u$ and $H_d$, are equal at $M_{GUT}$ in one set of models (FSU(5)-UH) but not in the other (FSU(5)). We have provided estimates of the flux, from annihilating neutralinos in the Sun, of neutrino induced muons, and considered prospects of detecting this flux in the IceCube/DeepCore detector. Some uncertainties arise in converting the muon flux into spin dependent neutralino-nucleon cross sections that we have briefly discussed. We offer comparisons with previously studied CMSSM and NUHM2 models, and also highlight some benchmark models in flipped SU(5) with varying neutralino compositions which can be tested by the IceCube/DeepCore experiment. Our results for NUHM2, FSU(5)-UH and FSU(5) show more points above the projected IceCube limit compared to the CMSSM, and hence a greater prospect of detection. This is to be expected because the models are less constrained than the CMSSM and possess additional free parameters." }, "1101/1101.5883_arXiv.txt": { "abstract": "We study the statistical properties of faint X-ray sources detected in the {\\it Chandra Bulge Field}. The unprecedented sensitivity of the Chandra observations allows us to probe the population of faint Galactic X-ray sources down to luminosities $L_{\\rm 2-10 keV}\\sim10^{30}$ \\ergss~ at the Galactic Center distance. We show that the luminosity function of these CBF sources agrees well with the luminosity function of sources in the Solar vicinity \\citep{sazonov06}. The cumulative luminosity density of sources detected in the CBF in the luminosity range $10^{30}-10^{32}$ \\ergss~ per unit stellar mass is $L_{\\rm 2-10 keV}/M_\\star=(1.7\\pm0.3)\\times10^{27}$ \\ergss~ $M_\\odot^{-1}$. Taking into account sources in the luminosity range $10^{32}-10^{34}$ \\ergss~ from \\cite{sazonov06}, the cumulative luminosity density in the broad luminosity range $10^{30}-10^{34}$ \\ergss~ becomes $L_{\\rm 2-10 keV}/M_\\star=(2.4\\pm0.4)\\times10^{27}$ \\ergss~ per $M_\\odot$. The majority of sources with the faintest luminosities should be active binary stars with hot coronae based on the available luminosity function of X-ray sources in the Solar environment. ", "introduction": "The hot interstellar plasma in galaxies \\cite[see e.g.][]{forman85} provides a valuable tool to probe different galactic phenomena, including energy and mass ejection from supernova explosions \\cite[e.g.][]{loewenstein91,maiz01,strickland07}, AGN activity \\cite[e.g.][]{padovani93,forman05}, and the contribution of non-thermal pressure in hot atmospheres\\cite[][]{churazov10}. However, before one can study the genuine X-ray emission of the hot interstellar plasma in early type galaxies, one must excise or account in some manner for the contribution of other constituents of galaxy emission. Revnivtsev et al. (2008) summarized the contributions from low mass X-ray binaries (see also \\citealt{gilfanov04}) and the fainter emission from the population of unresolved cataclysmic variables (CVs) and coronally active binary stars (ABs) (see Revnivtsev for a discussion of the origin of the faint diffuse emission). The appearance of the latest generation of X-ray observatories with fine angular resolution (e.g., as high as $0.6''$ from the {\\it Chandra} X-ray Observatory) provides the possibility to subtract/mask the contribution of the brightest X-ray binaries in nearby galaxies. The contribution of fainter X-ray binaries is often estimated via extrapolation of the observed luminosity function of bright sources towards lower luminosities \\cite[e.g.][]{sarazin01}, which Gilfanov (2004) has shown to be a good estimator for this component. In the past, is has also been assumed that all emission lines at energies $<$1-1.5 keV, visible in the spectra of galaxies, originate in a hot interstellar plasma \\cite[e.g.][]{sarazin01,irwin02,david06}. However, it is now clear (see \\citep{revnivtsev06,sazonov06,revnivtsev07,revnivtsev08,revnivtsev09,revnivtsev10}) that faint discrete galactic X-ray sources, namely cataclysmic variables and coronally active (usually binary) stars, provide an important contribution to the X-ray emission of galaxies (after subtraction of the contribution from the brightest sources) and this should be taken into account. Moreover, as the X-ray emission from these stars naturally includes a contribution from optically thin plasma of moderate or low temperatures, their spectra contain exactly the same lines that are usually thought to originate in the interstellar plasma. Until now we have had very limited knowledge of the broad band luminosity function of the population of faint Galactic X-ray sources (however, there are studies of different sub-populations, see e.g. \\citealt{pallavicini81,worrall83,hunsch99,schwope02}). Essentially, the best available broad band ($10^{27}$ \\ergss$~ 32$ from \\cite{sazonov06}, the cumulative luminosity density of sources with $\\log L_{\\rm 2-10 keV}>30$ is $(2.4\\pm0.4)\\times10^{27}$ \\ergss~ per $M_\\odot$. \\item The majority of the faintest sources should be active binaries with hot coronae, which makes their cumulative emission virtually indistinguishable from that of the hot interstellar plasma typically associated with old stellar populations. This indicates the importance of carefully accounting for the emission of such sources, based on detailed knowledge of their statistical properties and not only on the shape of the energy spectrum. \\end{itemize}" }, "1101/1101.1580.txt": { "abstract": "We present the results of a high resolution study with the Submillimeter Array towards the massive star forming complex G9.62+0.19. Three sub-mm cores are detected in this region. The masses are 13, 30 and 165 M$_{\\sun}$ for the northern, middle and southern dust cores, respectively. Infall motions are found with HCN (4-3) and CS (7-6) lines at the middle core (G9.62+0.19 E). The infall rate is $4.3\\times10^{-3}~M_{\\odot}\\cdot$yr$^{-1}$. In the southern core, a bipolar-outflow with a total mass about 26 M$_{\\sun}$ and a mass-loss rate of $3.6\\times10^{-5}~M_{\\odot}\\cdot$yr$^{-1}$ is revealed in SO ($8_{7}-7_{7}$) line wing emission. CS (7-6) and HCN (4-3) lines trace higher velocity gas than SO ($8_{7}-7_{7}$). G9.62+0.19 F is confirmed to be the driving source of the outflow. We also analyze the abundances of CS, SO and HCN along the redshifted outflow lobes. The mass-velocity diagrams of the outflow lobes can be well fitted by a single power law. The evolutionary sequence of the cm/mm cores in this region are also analyzed. The results support that UC~H{\\sc ii} regions have a higher blue excess than their precursors. ", "introduction": "%The mechanism of how high-mass stars are formed has long been %debated. Several theoretical modes have been proposed. High mass %stars would be formed via disk accretion with jets/outflow, merging %of smaller stars, turbulent accretion, competitive accretion, and %ionized accretion flows. High-mass stars play a major role in the evolution of the Galaxy. They are the principal sources of heavy elements and UV radiation \\citep{zin07}. However, the formation and evolution of high-mass stars are still unclear. A possible evolution sequence of high-mass stars from infrared dark clouds to classic H{\\sc ii} regions has been suggested \\citep{van05}. But one of the major topics whether high-mass stars form through accretion-disk-outflow, like low-mass ones \\citep{shu87}, or form via collision-coalescence \\citep{wol87,bon98} is still far from solved. Yet more and more observations at various resolutions seem to support the accretion-disk-outflow models rather than collision-coalescence models. Disks are detected in several high-mass star forming regions \\citep{pat05,jia05,sri05}. Outflows are found with a high detection rate as in low-mass cores in single-dish surveys \\citep{wu04,zha05,qin08a}. High resolution studies have also confirmed that molecular outflows are common in high-mass star forming regions \\citep{su04,qiu07,qin08b,qin08c,qiu09}. Searching for inflow motions also has made large progress in recent years \\citep{wu03,ful05,wyr06,kla07,wu07,wu09,fur11}. Both infall and outflow motions in the massive core JCMT 18354-0649S are detected \\citep{wu05}, and further confirmed by higher resolution observations \\citep{liu11}. Although accretion-disk-outflow systems are found in high-mass star forming regions, there may be differences between low- and high-mass formation. The infall motion can be detected via \"blue profile\", a double-peaked profile with the blueshifted peak being stronger for optically thick lines and a single peak at the absorption part of optically thick lines for optically thin lines, which is caused by self absorption of the cooler outer infalling gas towards the warmer central region \\citep{zho93}. In contrast, the \"red profile\" where the redshifted peak of a double-peaked profile being stronger for optically thick lines is suggested as indicators for outflow motions. \\cite{mar97} defined the \"blue excess\" in a survey, E, as E~=~(N$_{B}$-N$_{R}$)/N$_{T}$ (Mardones et al. 1997), where N$_{T}$ is number of sources, N$_{B}$ and N$_{R}$ mark the number of sources with blue and red profiles, respectively. The blue excess seems to be no significant differences among the low-mass cores in different evolutionary phases. However, using the IRAM 30 m telescope, \\cite{wu07} found that UC~H{\\sc ii} regions show a higher blue excess than their precursors, indicating fundamental differences between low- and high-mass-star forming conditions. The searches need to be expanded. Located at a distance of 5.7 kpc \\citep{hof94}, G9.62+0.19 is a well studied high-mass star forming region containing a cluster of H{\\sc ii} regions, which are probably at different evolutionary stages. Multiwavelength VLA observations have identified nine radio continuum sources (denoted from A-I) \\citep{gar93,tes00}, and components C-I are very compact ($<5\\arcsec$ in diameter) \\citep{gar93,tes00}. As revealed in NH$_{3}$ (4,4), (5,5) and CH$_{3}$CN (J=6-5), component F is a hot molecular core (HMC) and hence likely the youngest source in the region \\citep{ces94,hof94,hof96b}. G9.62+0.19 E is a young massive star surrounded by a very small UC~H{\\sc ii} region and a dusty envelope \\citep{hof96b}, while G9.62+0.19 D a small cometary UC~H{\\sc ii} region excited by a B0.5 ZAMS star \\citep{hof96b,tes00}. Both G9.62+0.19 E and G9.62+0.19 D seem to be at a more evolved stage than G9.62+0.19 F. Thus G9.62 complex is an ideal sample to examine massive star forming activities including outflow and infall motions. Maser emissions of NH$_{3}$, H$_{2}$O, OH, and CH$_{3}$OH, as well as the strong thermal NH$_{3}$ emissions were detected along a narrow region with projected length $20\\arcsec$ and width$\\leq2\\arcsec$ \\citep{hof94}. A possible explanation for this alignment is compression of the molecular gas by shock front originating from an even more evolved H{\\sc ii} region to the west of the star-forming front \\citep{hof94}. High-velocity molecular outflows also have been detected in this region, and G9.62+0.19 F is believed to be the driving source \\citep{gib04,hof01,su05}. However, most of previous work was carried out at low frequencies, probing low excitation conditions. To exam the hot dust/gas environment and dynamical processes in this region, higher resolution studies at high frequencies are needed. In this paper we report the results of the Submillimeter Array (SMA\\footnote {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.}) observations toward G9.62+0.19 region at 860 $\\micron$. ", "conclusions": "\\subsection{Rotational temperature of H$_{2}$CS transitions} Six transitions of H$_{2}$CS have been detected in the middle and southern cores, enabling us to estimate the rotational temperature. Under the assumptions that the gas is optically thin under local thermodynamic equilibrium and the gas emission fills the beam, the rotation temperature and beam-averaged column density can be estimated using the Rotational Temperature Diagram (RTD) by \\citep{cum86,tur91,liu02} \\begin{equation} \\textrm{ln}(\\frac{N_{u}}{g_{u}}) = \\textrm{ln}(\\frac{N_{T}}{Q_{rot}})-\\frac{E_{u}}{T_{rot}} = \\textrm{ln}[2.04\\times10^{20}\\frac{\\int~I(Jy~beam^{-1})dv(km~s^{-1})}{\\theta_{a}\\theta_{b}(arcsec^{2})g_{I}g_{K}\\nu^{3}(GHz^{3})S\\mu^{2}(debye^{2})}] \\end{equation} where N$_{u}$ is the observed column density of the upper energy level, g$_{u}$ is the degeneracy factor in the upper energy level, N$_{T}$ is the total beam-averaged column density, Q$_{rot}$ is the rotational partition function, E$_{u}$ is the upper level energy in K, T$_{rot}$ is the rotation temperature, $\\int$~I~dv is the integrated intensity of the specific transition, $\\theta_{a}$ and $\\theta_{b}$ are the FWHM beam size, g$_{K}$ is the K-ladder degeneracy, g$_{I}$ is the degeneracy due to nuclear spin, $\\nu$ is the rest frequency, and S is line strength and $\\mu$ the permanent dipole moment. For H$_{2}$CS, the interchangeable nuclei are spin $\\frac{1}{2}$, leading to ortho- and para-forms with g$_{I}$ equaling $\\frac{3}{4}$ and $\\frac{1}{4}$, respectively \\citep{bl87,tur91}. The partition function Q$_{rot}$ of H$_{2}$CS is \\citep{bl87} \\begin{equation} Q_{rot} = 2[\\frac{\\pi(kT_{rot})^{3}}{h^{3}ABC}]^{\\frac{1}{2}} \\end{equation} where k and h are the Boltzmann and Planck constants, respectively, and A, B, and C are the rotation constants. Thus the rotation temperature T$_{rot}$ and total column density N$_{T}$ can be estimated by least-squares fitting to the multiple transitions. We applied the RTD method towards D, E, F, G (see Figure 14), and the fitting results are listed in the second and third columns of Table 3. The rotational temperature of the middle core (E) is 83$\\pm$21 K. In the southern core, the rotational temperature estimated decreases from G (91 K) to F (83 K) and D (43 K), suggesting the temperature gradient in the southern core. The total column density of H$_{2}$CS ranges from 1.3$\\times$10$^{15}$ (G) to 3.8$\\times$10$^{15}$~cm$^{-2}$ (D). However, the filling factor and the optical depth correction were not taken account of in the RTD method. To investigate their effect we applied the Population Diagram (PD) analysis \\citep{gol99,wang10}. In the PD analysis, we have \\begin{equation} \\textrm{ln}(\\frac{\\hat{N_{u}}}{g_{u}}) = \\textrm{ln}(\\frac{N_{T}}{Q_{rot}})-\\frac{E_{u}}{T_{rot}}+\\textrm{ln}(f)-\\textrm{ln}(\\frac{\\tau}{1-e^{-\\tau}}) \\end{equation} where $\\hat{N_{u}}$ is the inferred column density of the upper energy level from the PD analysis, f is the source filling factor and $\\tau$ is the optical depth. The optical depth $\\tau$ can be expressed by \\citep{re04} \\begin{equation} \\tau = \\frac{8\\pi^{3}S\\mu^{2}\\nu}{3k\\Delta\\textrm{v}T_{rot}}\\frac{N_{T}}{Q_{rot}}e^{-\\frac{E_{u}}{T_{rot}}} \\end{equation} where $\\Delta$v is the FWHM line width. Under LTE, the upper-level populations, $\\hat{N_{u}}$, can be predicted according to the right-hand side of Equation (3) for a given set of total column density, N$_{T}$, rotational temperature, T$_{rot}$, and source filling factor, f. The expected $\\hat{N_{u}}$ were evaluated for the parameter space of T$_{rot}$ = 10-500 K, N$_{T}$ = 10$^{14}$-10$^{17}$ cm$^{-2}$, and f between 0.01 and 1.0. To compare the observed $N_{u}$ and the inferred $\\hat{N_{u}}$, we calculate the $\\chi^{2}$ as: \\begin{equation} \\chi^{2} = \\sum(\\frac{N_{u}-\\hat{N_{u}}}{\\delta~N_{u}})^{2} \\end{equation} where $\\delta~N_{u}$ is the 1 $\\sigma$ error of observed upper-state column density. Although the $\\chi^{2}$ is a good representation of the goodness of fit, the parameter set with the lowest $\\chi^{2}$ may not actually represent physical parameters very well due to the uncertainties of the observed data. In order to find a representative parameter set, we compute a weighted mean and standard deviation for all the parameters, with the weights being the inverse of the $\\chi^{2}$. All the parameter sets where the inferred upper-level population $\\hat{N_{u}}$ corresponds with the observed upper-level population $N_{u}$ within 3 $\\sigma$ are used to compute the weighted means and standard deviations. The derived rotational temperature, total column density and filling factor of each component are list in the [3-5] columns of Table 3. The inferred optical depths of each line transition are listed in the last six columns of Table 3. The rotational temperatures of D, E, F, G are estimated to be 42$\\pm$34, 92$\\pm$74, 51$\\pm$23 and 105$\\pm$37 K, respectively. A temperature gradient in the southern core is also revealed as in the RTD method. The four components D, E, F, G has similar total column densities as high as 4$\\times10^{16}$ cm$^{-2}$, about an order of magnitude higher than those obtained from RTD method, which are mainly due to the small source filling factor ($<$0.5). The optical depths of H$_{2}$CS (10$_{0,10}$-9$_{0,9}$) at the four components are all much larger than one, while the other transitions are always optically thin except H$_{2}$CS (10$_{2,9}$-9$_{2,8}$) line at G. \\subsection{Core properties} In the optically thin case, the total dust and gas masses of the three sub-mm cores can be obtained with the formula $M=S_{\\nu}D^{2}/\\kappa_{\\nu}RB_{\\nu}(T_{d})$ \\citep{hil83}, where $S_{\\nu}$ is the flux at 860 $\\micron$, D is the distance, R=0.01 is the mass ratio of dust to gas, and $\\kappa_{\\nu}$ is dust opacity per unit dust mass. $B_{\\nu}(T_d)$ is the Planck function at a dust temperature of T$_{d}$. We assume that T$_{d}$ equals the rotational temperature of H$_{2}$CS. For the northern core, since only CS (7-6) (upper energy E$_{u}$ = 65.8 K) and HCN (4-3) (E$_{u}$ = 42.5 K) exhibit strong emission lines, we assume T$_{d}$ to be 50 K. Together with the measurements at centimeter and millimeter wavelengths, \\cite{su05} extrapolated the ionized gas emission at mm/submm wavelengths, and found that the 0.85 mm continuum associated with components D, E, and F are dominated by thermal dust emission. They have derived opacity index $\\beta$ of components E and F to be 1.2, and 0.8, respectively. For the northern sub-mm core, ${\\beta}=1.5$ is assumed. Using the above dust opacity indexes, we adopt $\\kappa_{\\nu}$=2.0, 1.8, and 1.5~cm$^{2}$g$^{-1}$ for the northern, middle and southern cores, respectively \\citep{ossen94}. At the distance of 5.7 kpc , we get the total dust and gas masses for these three cores, and list all the parameters in Table 1. The deduced masses for the northern, middle and southern cores are 13, 30, 165 M$_{\\sun}$, respectively. The column density of H$_{2}$ are 1.2$\\times10^{24}$ and 2.1$\\times10^{24}$ cm$^{-2}$ for the middle and southern sub-mm cores, respectively. \\subsection{Infall properties in the middle core} In the middle core, both CS(7-6) and HCN(4-3) emission exhibits \"blue profile\" feature, indicating infall motions of the gas envelope toward the central star \\citep{ket88,zho93,zha98,wu03,wu05,wu07,ful05,wyr07,sun08}. The velocity difference (0.9 km~s$^{-1}$) between the absorption dip in CS (7-6) spectrum (3 km~s$^{-1}$) and the systemic velocity (2.1 km~s$^{-1}$) is taken as the infall velocity $V_{in}$. Since both HCN (4-3) and CS (7-6) emissions are not resolved towards the middle core, we simply take the dust core size as the radius of the infall region, which may underestimate the infall rate derived below. The kinematic mass infall rate can be calculated using dM/dt=$4{\\pi}R_{in}^{2}nmV_{in}$. n=1.5$\\times10^{7}$cm$^{-3}$ is the number density of this dust core. Taking Helium into account, the mean molecular mass m is 1.36 times of H$_{2}$ molecule mass. %Since the core diameter of 50\\% peak intensity is $5''$ which is %much larger than the $3.81''\\times2.83''$ beam size, we note that %the adopted radius is an lower limit. So is the mass infall rate. The infall rate calculated is $4.3\\times10^{-3}~M_{\\odot}\\cdot$yr$^{-1}$. For comparison, the $V_{in}$ from pure free-infall assumption is also derived with the formula $V_{in}^{2}=2GM/R_{in}$. The pure free-infall velocity is $V_{in}=3.6~$km~s$^{-1}$ and thus the \"gravitational\" mass infall rate is $1.7\\times10^{-2}~M_{\\sun}\\cdot$yr$^{-1}$, which is larger than the kinematic infall rate. \\subsection{Outflow properties in the southern core} \\subsubsection{Shock chemistry in the outflow region of the southern core} Observations have suggested that there are important differences in molecular abundances in different outflow regions \\citep{bac97,cho04,jor04,cod05}. Significant abundance enhancements are found in the shocked region for sulfur-bearing molecules \\citep{bac97,jor04}, and the abundance of HCN in outflow regions is related to atomic carbon abundance \\citep{cho02}. However, previous studies of the chemical impact of outflows are confined to the well collimated outflows around Class 0 sources, while such studies especially high resolution studies on massive outflows are rare \\citep{bac97,jor04,arc07}. A red and bright IRAC source is found to be associated with the southern core. The magnitudes of the IRAC source at 3.6 $\\micron$, 4.5 $\\micron$ and 5.8 $\\micron$ are $10.102\\pm0.093$, $8.361\\pm0.108$ and $7.778\\pm0.302$ mag, respectively. The [3.6-4.5] color is as large as 1.74, indicating shocked emission in the southern core \\citep{tak10}. Maser emissions of NH$_{3}$, H$_{2}$O, OH, and CH$_{3}$OH, as well as the strong thermal NH$_{3}$ emissions also uncover the existence of the shocked gas \\citep{hof94}. Outflows can be revealed from shocked H$_{2}$ emission probed by the strong and extended emission at the 4.5 $\\micron$ band \\citep{qiu08,tak10}. Thus the massive outflow in the southern core of G9.62 complex provides an ideal sample to study shock chemistry. The fractional abundance of a certain molecule is defined as $\\chi = N_{T}/N_{H_{2}}$, where $N_{T}$ is the total column density of a specific molecule and $N_{H_{2}}$ is the H$_{2}$ column density. Assuming that the gas is optically thin and the emission fills the beam, the beam-averaged total column density of a specific molecule can be obtained from: \\begin{equation} N_{T} = 2.04\\times10^{20}\\frac{\\int~I(Jy~beam^{-1})dv(km~s^{-1})Q_{rot}e^{E_{u}/T_{rot}}}{\\theta_{a}\\theta_{b}(arcsec^{2})g_{I}g_{K}\\nu^{3}(GHz^{3})S\\mu^{2}(debye^{2})} \\end{equation} Assuming that T$_{rot}$ of HC$^{15}$N equals to that of H$_{2}$CS and the gas is optically thin, $N_{T}$ of HC$^{15}$N is calculated to be $3.0\\times10^{13}$~cm$^{-2}$ at the core region. At the galactocentric distance of 3 kpc for G9.62+0.19 (Scoville et al.1987,Hofner et al.1994), the abundance ratio [$^{14}$N]/[$^{15}$N]$~\\approx~350$ (Wilson and Rood. 1994). Thus the total column density of HCN at the core region should be $1.1\\times10^{16}$~cm$^{-2}$. Therefore, the fractional abundance of HCN relative to H$_{2}$ at the core region is $5.2\\times10^{-9}$. HCN appears to be greatly enhanced in the outflow regions of the L1157 \\citep{bac97}, while has similar abundances in the outflow region and the ambient cloud of NGC 1333\u00a8CIRAS 2A \\citep{jor04}. Owing to the lack of a direct estimation of the H$_{2}$ column density towards the outflow region, the fractional abundance of HCN in the outflow region is also assigned to $5.2\\times10^{-9}$ in calculating the outflow parameters. Since the HC$^{15}$N emission traces outflowing gas at much lower velocity than HCN, perhaps HCN could be more enhanced in the high velocity component. With the possibility of higher opacity and the lack of direct H$_{2}$ column density measurement, the derived fractional abundance perhaps is a lower limit anyway. \\cite{su07} estimate an HCN abundance of $\\sim1-2\\times10^{-8}$ in the massive outflow lobes of IRAS 20126+4104, which is comparable to our estimation here. Since the blueshifted outflow gas traced by CS (7-6) and HCN (4-3) suffers self-absorption, the abundance ratios among SO ($8_{7}-7_{7}$), CS (7-6), and HCN (4-3) were inferred from the beam-averaged spectra taken from the redshifted outflow lobe. The abundance ratio as a function of flow velocity (the outflow velocity relative to the systemic velocity) of [CS/SO] is obtained assuming five different excitation temperatures in the left panel of Figure 15. It can be seen that the abundance ratio of [CS/SO] increases with the excited temperature. At each excitation temperature, the abundance ratio of [CS/SO] has lower values at flow velocities less than 6~km~s$^{-1}$, and higher values when V$_{flow}$ larger than 8~km~s$^{-1}$, whereas the abundance ratio seems to be constant at flow velocities between 6~km~s$^{-1}$ and 8~km~s$^{-1}$. There are two reasons for the lower abundance ratio when V$_{flow}~<$~6~km~s$^{-1}$: first, the flux missing of CS (7-6) due to the interferometer is more serious than SO ($8_{7}-7_{7}$); second, CS (7-6) may be more optically thick at lower flow velocities than SO ($8_{7}-7_{7}$). As shown in the P-V diagrams, the emission region of CS (7-6) is much larger than SO ($8_{7}-7_{7}$) at high velocities. The higher abundance ratio when V$_{flow}~>$~8~km~s$^{-1}$ is due to the smaller filling factor of SO ($8_{7}-7_{7}$) emission. We propose the mean observed value between 6~km~s$^{-1}$ and 8~km~s$^{-1}$ can represent the actual abundance ratio of [CS/SO]. Assuming a typical excitation temperature of T$_{ex}$=30 K \\citep{wu04}, the abundance ratio of [CS/SO] at the redshifted lobe is inferred as 0.7. \\cite{nil00} find that the [SO/CS] abundance ratios are strongly enhanced in the Orion A and NGC 2071 outflows where the [SO/CS] ratios are estimated to be about 24 and 2.2, respectively. However, the [SO/CS] abundance ratio in the outflow of G9.62+0.19 is found to be 1.4, much lower than that found in Orion A outflow. As shown in the right panel of Figure 15, the abundance ratio of [CS/HCN] decreases linearly with the flow velocity. To avoid the missing flux difficulty, the abundance ratio is calculated at high flow velocities larger than 7 km~s$^{-1}$. The decreasing of the abundance ratio with velocity is because that the emission region traced by CS (7-6) is always smaller than HCN (4-3), leading to smaller filling factor for CS (7-6), which can be verified easily by comparing the channel maps between CS (7-6) in Figure 11 and HCN (4-3) in Figure 12 at high velocities. We fitted the observed data with a linear function, and adopted the value at flow velocity of 10 km~s$^{-1}$ as the actual abundance ratio of [CS/HCN] in the outflow region, which is [CS/HCN]=1.2. Since HCN fractional abundance is $5.2\\times10^{-9}$, the fractional abundances of CS and SO are deduced to be $6.2\\times10^{-9}$ and $8.9\\times10^{-9}$, respectively. \\subsubsection{Properties of the bipolar-outflow traced by SO ($8_{7}-7_{7}$) emission} The SO ($8_{7}-7_{7}$) emission in the southern core shows line wings, suggesting outflow motions. From the integrated intensity map in Figure 10(c), we find the outflow lobes revealed by SO ($8_{7}-7_{7}$) emission peak at different position with different position angle compared with previously reported H$_{2}$S ($2_{2,0}-2_{1,1}$) \\citep{gib04} and HCO$^{+}$ (1-0) data \\citep{hof01}. But in the same sense, the blue- and red-lobes revealed by SO overlap to a large extent as well as HCO$^{+}$ (1-0) and H$_{2}$S ($2_{2,0}-2_{1,1}$) data, consistent with the argument of the outflow being viewed pole-on \\citep{hof01}. The total mass of each outflow lobe is given by: \\begin{equation} M_{flow} = 1.04~\\times~10^{-4}D^{2}\\frac{Q_{rot}e^{E_{u}/T_{rot}}}{\\chi\\nu^{3}S\\mu^{2}}\\int\\frac{\\tau}{1-e^{-\\tau}}S_{\\nu}dv \\end{equation} where M$_{flow}$, D, S$_{\\nu}$, $\\chi$, and $\\tau$ are the outflow gas mass in M$_{\\sun}$, source distance in kpc, line flux density in Jy, relative abundance to H$_{2}$, and optical depth. The other parameters have the same units as in equation (1). The fractional abundance of SO is taken as $8.9\\times10^{-9}$ (see Sec.4.4.1). Assuming an excitation temperature of 30 K and the outflowing gas is optically thin, the inferred outflow masses are 13 M$_{\\sun}$ for each of red and blueshifted lobes. Thus, the momentum can be calculated by $P=\\sum$M(v)dv, and the energy by $E=\\sum{\\frac{1}{2}}$M(v)v$^{2}$dv, where $v$ is the flow velocity. The derived parameters are listed in Table 4. The momentum and energy of the red lobe are 82 M$_{\\sun}\\cdot$km~s$^{-1}$ and $5.4\\times10^{45}$ erg. For the blue lobe, the momentum and energy are calculated to be 86 M$_{\\sun}\\cdot$km~s$^{-1}$ and $5.8\\times10^{45}$ erg. The dynamical timescale t$_{dyn}$ is estimated as R/V$_{char}$, where R ($\\sim$ 0.06 pc) is adopted as the mean size of the outflow lobes assuming a collimation factor of unity, and V$_{char}$ ($\\sim$ 5.5 km~s$^{-1}$) is assumed as the mass weighted mean velocity. Thus, the dynamic timescale is estimated to be $1\\times10^{4}$ year, which may be underestimated due to the uncertainty of the outflow scale. The mechanical luminosity L, and the mass-loss rate $\\dot{M}$ are calculated as L=E/t, $\\dot{M} = P/(tV_{w})$, where the wind velocity V$_{w}$ is assumed to be 500 km~s$^{-1}$ \\citep{lam95}. The mechanical luminosity L and the total mass-loss rate are estimated to be 9.3 L$_{\\sun}$ and $3.6\\times10^{-5}$~M$_{\\sun}$$\\cdot$yr$^{-1}$, respectively. \\subsubsection{Very high-velocity gas detected in CS (7-6) emission} The CS (7-6) emission at the southern core shows \"red-profile\" with wide wings. We take $6.2\\times10^{-9}$ as the fractional abundance of CS relative to H$_{2}$ along the outflow lobes. Assuming T$_{ex}$ = 30 K, we derive the parameters for the CS outflow (Table 4) with the same method used for SO ($8_{7}-7_{7}$). The outflow masses at very high velocities (v$_{flow}~>~$10~km~s$^{-1}$) are 3.7 M$_{\\sun}$ and 5.5 M$_{\\sun}$ for the blueshifted and redshifted lobes, respectively. The momentum and energy of the blueshifted lobe at very high velocities are calculated to be 47 M$_{\\sun}\\cdot$km~s$^{-1}$ and $6.0\\times10^{45}$ erg. For the redshifted lobe, the momentum and energy at extremely high velocities are calculated to be 68 M$_{\\sun}\\cdot$km~s$^{-1}$ and $8.7\\times10^{45}$ erg, which are similar to the blueshifted lobe. \\subsubsection{Very high-velocity gas detected in HCN (4-3) emission} As discussed before, HCN (4-3) has a velocity extent of at least 60 km~s$^{-1}$, which traces extremely high-velocity (EHV) gas. Adopting an excited temperature of 30 K, and an HCN-to-H$_{2}$ abundance ratio of $5.2\\times10^{-9}$, the parameters of the outflow are calculated and listed in Table 4. The outflow mass at very high velocities (v$_{flow}~>~$10~km~s$^{-1}$) are 5.2 M$_{\\sun}$ and 17.6 M$_{\\sun}$ for the blueshifted and redshifted lobes, respectively. The momentum and energy of the blueshifted lobe at very high velocities are 85 M$_{\\sun}\\cdot$km~s$^{-1}$ and $1.4\\times10^{46}$ erg. For the redshifted lobe, the momentum and energy at very high velocities are 294 M$_{\\sun}\\cdot$km~s$^{-1}$ and $5.5\\times10^{46}$ erg, which are larger than the blueshifted lobe. \\subsubsection{Mass-Velocity diagrams} A broken power law, $dM(v) / dv \\propto v^{-\\gamma}$ usually exhibits in molecular outflows near young stellar objects \\citep{cha96,lad96,rid01,su04,qiu07,qiu09}. The slope, $\\gamma$, typically ranging from 1 to 3 at low outflow velocities, and often steepens at velocities larger than 10 km~s$^{-1}$ --- with $\\gamma$ as large as 10 in some cases \\citep{arc07}. Assuming optically thin, the mass-velocity diagrams of the outflow at the southern core of G9.62+0.19 complex are shown in Figure 16. SO ($8_{7}-7_{7}$), CS (7-6), HCN (4-3) results were all used in the mass spectra. We calculate the outflow mass traced by CS (7-6) and HCN (4-3) from V$_{flow}$ of 10 km~s$^{-1}$ to avoid the absorption of the spectra. Instead of broken power law appearance, the mass-velocity diagram of blueshifted lobe can be well fitted by a single power law with a power indexes of $2.28\\pm0.23$. The mass-velocity diagram of redshifted lobe can be well fitted by a single power law with a power indexes of $1.70\\pm0.17$ even though the mass drops more rapidly after 25 km~s$^{-1}$. As marked by the dashed ellipse in the right panel, the outflow mass revealed by CS (7-6) is much lower than that revealed by HCN (4-3) at very high velocities. Despite the CS data, the mass-velocity diagram of redshifted lobe at velocities smaller than 25 km~s$^{-1}$ can be fitted by a single power law with a much smaller power indexes of $1.08\\pm0.09$. However, no significant slope changes are found in both the red- and blue-shifted lobes of the outflow at the southern core, which are very different from those previous works. \\subsection{Different evolutionary stages of the three dust cores} The northern core has the smallest diameter and mass among the three cores. It seems likely to be a point source after deconvolution. It is located south of the nominal radio UC~H{\\sc ii} region G9.62+0.19 C. In this region, eight near-IR sources are detected in a diffuse near-IR nebulosity at the west of the radio emission peak \\citep{per03}. The reddest one c7 (18$^{\\rm h}$06$^{\\rm m}$14.34$^{\\rm s}$,-20$\\arcdeg$31$\\arcmin$25.0$\\arcsec$) is located within $1\\arcsec$ of the radio peak, while the faintest one c8 (18$^{\\rm h}$06$^{\\rm m}$14.42$^{\\rm s}$,-20$\\arcdeg$31$\\arcmin$27.4$\\arcsec$) seems to be associated with the sub-mm core detected in SMA observation. Source c8 is too faint to be detected even at H band and also shows no emission at 12.5 $\\micron$. In contrast to the bright, rich molecular spectrum forest in the middle and southern sub-mm cores, the northern sub-mm core lacks strong molecular emissions. There is also no other early star forming signature such as masers associated with it. Since it is with near-IR emission and at the edge of the UC~H{\\sc ii} region G9.62+0.19 C, the northern core may be just a remnant core in the envelope of UC~H{\\sc ii} region G9.62+0.19 C, which needs further observations. The middle core is associated with the hyper-compact H{\\sc ii} region G9.62+0.19 E \\citep{gar93,kur02}. OH, H$_{2}$O, and NH$_{3}$ (5,5) masers have been detected near the radio emission peak \\citep{for89,hof94,hof96a}. Periodic class II methanol masers are also found in G9.62+0.19 E \\citep{van09,goe05,nor93}. Methanol masers are believed to be a good tracer of young massive star forming regions at stages earlier than relatively evolved UC~H{\\sc ii} regions \\citep{lon07}. No infrared source coincides with G9.62+0.19 E \\citep{per03}. Hot molecular CH$_{3}$CN lines are detected in this region, and a kinematic temperature of T$_{k}$ = 108 K was obtained from CH$_{3}$CN emission with LVG model \\citep{hof96b}, which is coincident with the rotational temperature (T$_{rot}$ = 92 K) obtained from H$_{2}$CS emission. A spectra forest including hot molecular lines, such as CH$_{3}$OH, is detected towards G9.62+0.20 E, suggesting this core is in a hot phase. Infall motions are traced by CS (7-6) and HCN (4-3) lines, indicating active star forming in this region. All above suggest that G9.62+0.20 E is forming a massive young star. The 860 $\\micron$ dust emission of the southern core peaks at G9.62+0.19 F, and extends from north to south. A hump structure is found to the southeast of the emission peak, indicating another possible sub-mm core. The previously recognized mm/cm cores (G9.62+0.19 D, G) are at the edges of the southern core. G9.62+0.19 G is a weak radio source \\citep{tes00}, while G9.62+0.19 D is consistent with an isothermal UC~H{\\sc ii} region excited by a B0.5 star \\citep{hof96b}. Weaker radio emission was found at core F. H$_{2}$O and OH masers are found across the whole sub-mm core from north to south \\citep{for89,hof96a}. A near-IR source with large NIR excess is found to be associated with G9.62+0.19 F \\citep{tes98,per03}. With higher resolution observations \\cite{linz05} found four near-IR objects (F1-F4) in this core. F4 is with little emission at K band but becomes redder at longer wavelengths, which seems to correspond to the bright IRAC source with large excess at 4.5 $\\micron$. This object is the dominating and closest associated source of core F. Core F is also confirmed to be the driving source of an active outflow. All of above imply that G9.62+0.19 F is a very young massive star forming region. \\subsection{Blue excess in high-mass star forming regions} Wu et al. (2007) found that UC~H{\\sc ii} regions show a higher blue excess than UC~H{\\sc ii} precursors with the IRAM 30 m telescope. \\cite{wyr06} also detected large blue excess in UC~H{\\sc ii} regions. \"Blue profile\" was detected with CS (7-6) and HCN (4-3) lines in UC~H{\\sc ii} region G9.62+0.19 E, while \"red profile\" in hot molecular core G9.62+0.19 F, which coincides with their argument. The detection of infall signature in G9.62+0.19 E also coincides the interpretation that material is still accreted during the UC~H{\\sc ii} phase \\citep{wu07,keto02}. Around younger cores, the outflow is more active and cold than UC~H{\\sc ii} regions, which leads to more \"red profile\". While in UC~H{\\sc ii} regions, the outflows become weak. The surrounding gas of UC~H{\\sc ii} regions is thermalized and the temperature gradient towards the central star is more likely to cause \"blue profile\", which results in the higher blue excess than UC~H{\\sc ii} precursors." }, "1101/1101.3007_arXiv.txt": { "abstract": "The ideal MHD equations are a central model in astrophysics, and their solution relies upon stable numerical schemes. We present an implementation of a new method, which possesses excellent stability properties. Numerical tests demonstrate that the theoretical stability properties are valid in practice with negligible compromises to accuracy. The result is a highly robust scheme with state-of-the-art efficiency. The scheme's robustness is due to entropy stability, positivity and properly discretised Powell terms. The implementation takes the form of a modification of the MHD module in the FLASH code, an adaptive mesh refinement code. We compare the new scheme with the standard FLASH implementation for MHD. Results show comparable accuracy to standard FLASH with the Roe solver, but highly improved efficiency and stability, particularly for high Mach number flows and low plasma $\\beta$. The tests include 1D shock tubes, 2D instabilities and highly supersonic, 3D turbulence. We consider turbulent flows with RMS sonic Mach numbers up to 10, typical of gas flows in the interstellar medium. We investigate both strong initial magnetic fields and magnetic field amplification by the turbulent dynamo from extremely high plasma $\\beta$. The energy spectra show a reasonable decrease in dissipation with grid refinement, and at a resolution of $512^3$ grid cells we identify a narrow inertial range with the expected power-law scaling. The turbulent dynamo exhibits exponential growth of magnetic pressure, with the growth rate twice as high from solenoidal forcing than from compressive forcing. Two versions of the new scheme are presented, using relaxation-based 3-wave and 5-wave approximate Riemann solvers, respectively. The 5-wave solver is more accurate in some cases, and its computational cost is close to the 3-wave solver. ", "introduction": "In order to model complex nonlinear astrophysical phenomena, numerical solutions of the equations of ideal MHD are central. Examples include the dynamics of stellar atmospheres, star formation and accretion discs. Magnetised gas flows in astrophysics are typically highly compressible, nonlinear and often supersonic. Hence, obtaining stable numerical results in such regimes is extremely challenging if the numerical scheme must also be both accurate and efficient. Here we present a new MHD solver that preserves stable, physical solutions to the compressible MHD equations by construction, while at the same time showing improved efficiency and comparable accuracy to standard schemes based on the very accurate, but less stable Roe approximate Riemann solver \\cite{Roe1981}. We consider numerical solutions of the ideal MHD system (here written in conservation form, and in three dimensions, letting $I_3$ denote the $3\\times3$ identity matrix) \\begin{equation}\\begin{array}{l} \\rho_t + \\nabla\\cdot(\\rho \\bfu)=0,\\\\ (\\rho \\bfu)_t + \\nabla\\cdot(\\rho \\bfu\\otimes \\bfu + (p+ \\frac{1}{2} |\\bfB|^2)\\, {I_3} - \\bfB\\otimes \\bfB)=0,\\\\ E_t + \\nabla\\cdot[(E+p+\\frac{1}{2} |\\bfB|^2)\\bfu - (\\bfB\\cdot \\bfu)\\bfB]=0,\\\\ \\bfB_t + \\nabla\\cdot(\\bfB\\otimes \\bfu - \\bfu\\otimes \\bfB)=0,\\\\ \\nabla\\cdot\\bfB=0, \\label{eq:MHDP} \\end{array} \\end{equation} with the specific internal energy $e$, such that the total energy density is given by $E=\\rho e + \\half\\rho\\bfu^2 + \\half|\\bfB|^2$, and the pressure given by the equation of state $p=p(\\rho,e)$. The Cartesian components of the velocity are denoted by $\\bfu=(u,v,w)$. The system fits the generic form of a conservation law $U_t + \\nabla\\cdot \\mathbf{F}(U)=0$, except for the restriction on $\\nabla\\cdot\\bfB$. However, if this restriction is satisfied at the initial time $t=0$, it automatically holds at later times $t>0$ for the exact solution. Shock conditions are common in astrophysics, hence \\eqref{eq:MHDP} should be understood in a weak sense. The lack of regularity and large scale ranges in astrophysical flows make it particularly challenging to devise numerically stable schemes. Recent developments in nonlinear stability analysis have made such schemes possible, however. The central stability notions are entropy stability and positivity of mass density $\\rho$ and internal energy $\\rho e$. This paper presents an implementation of the robust positive second order scheme from \\cite{W1}, which uses the entropy stable approximate Riemann solvers of \\cite{BKW1, BKW2}. Several codes are available for high-performance compressible MHD simulations. We mention Athena \\cite{Athena2008}, AstroBEAR \\cite{Astrobear2009}, Chombo \\cite{Chombo2005}, Enzo \\cite{Enzo2008}, FLASH \\cite{FLASH}, Nirvana \\cite{Nirvana2005}, Pluto \\cite{Pluto2007}, RAMSES \\cite{RAMSES2006} and VAC \\cite{VAC1996}. The positive scheme of \\cite{W1} has so far been unavailable in such codes. We have implemented it as an alternative MHD module in the FLASH code. We focus here on a comparison with the original FLASH code. Benchmark tests in one and two spatial dimensions are presented. Transition to unstable turbulence-like flow is particularly emphasised. Additional tests of the algorithms may be found in \\cite{BKW2, KW2010} and \\cite{W1}. As an example application we consider turbulence in molecular clouds, which is characterised by high Mach numbers and the presence of magnetic fields. Before presenting the results, we give a review of the numerical schemes and their theoretical properties. ", "conclusions": "We presented an implementation of an accurate, efficient and highly stable numerical method for MHD problems. The method is reviewed here, and presented in detail in \\cite{W1}. It is implemented as a modification of the FLASH code \\cite{FLASH}, which enables large-scale, multi-processor simulations and adaptive mesh refinement. In \\cite{W1} it was found that our method could handle significantly larger ranges of the sonic Mach number and plasma $\\beta$ than a standard MHD scheme. This was confirmed in this paper by comparisons with the standard FLASH code. The algorithmic changes underlying the increased stability can be broken down into three parts: \\begin{enumerate} \\item An entropy stable approximate Riemann solver that preserves positivity of density and internal energy (\\cite{BKW1,BKW2}). \\item For second order accuracy, a reconstruction method that ensures positivity (\\cite{W1}). \\item In multidimensions, a stable discretisation of the Powell system (\\cite{W1}). \\end{enumerate} All these ingredients were essential in obtaining the desired stability and efficiency of the overall scheme. The different elements of the new scheme have been studied separately in previous papers. While \\cite{W1} focused on the positive second-order algorithm and multidimensionality, only a single approximate Riemann solver, HLL3R, was considered. The present study contrasts a standard scheme to the combination of these three new ingredients. The new scheme was implemented in two versions featuring the 3-wave (HLL3R) and 5-wave (HLL5R) approximate Riemann solvers of \\cite{BKW2} respectively, while there were two standard implementations of the FLASH code (version 2.5), using the Roe and the HLLE approximate Riemann solvers. We observed some increase of numerical dissipation compared to the Roe solver of FLASH, but it was minor, and due to the replacement of the Roe solver with the robust and efficient HLL-type solvers. The HLLE solver was found to be the most dissipative, while HLL5R showed almost identical dissipation properties and accuracy to the Roe solver. HLL3R was ranked between HLLE and Roe in terms of accuracy. As a physical application, we have considered forced MHD turbulence at high sonic Mach number. We were able to compare the new and old schemes at RMS sonic Mach number 2 with an initial plasma $\\beta=0.25$. The schemes were all found to give similar and reasonable results, but the new schemes HLL3R and HLL5R were altogether about eight times more efficient in this test. The Roe solver-based scheme in the FLASH code was slightly less dissipative, but had to be run at a four times lower CFL number to be stable. At RMS sonic Mach number 10 only the new schemes yielded physical results. We found reasonable dependence of dissipation on numerical resolution at this Mach number, and were able to infer a small inertial range from the velocity power spectra at $512^3$ resolution. Finally, we studied the turbulent dynamo action at RMS Mach number 5 with the new scheme. So far, there have been very few studies of turbulent dynamos in the supersonic regime. We found dynamo-generated exponential growth rates of the magnetic pressure that differed according to the type of forcing mechanism, i.e., solenoidal versus compressive forcing \\cite{FederrathKlessenSchmidt2008,FederrathKlessenSchmidt2009,FederrathDuvalKlessenSchmidtMacLow2010}. Turbulence simulations with the new scheme across a wide range of modest to large sonic and Alfv{\\'e}nic Mach numbers have been presented in \\cite{BruntFederrathPrice2010a,BruntFederrathPrice2010b}. The two relaxation-based approximate Riemann solvers HLL3R and HLL5R have previously not been compared at high order and in higher spatial dimensions than 1D. In many cases we found them to give similar results, with HLL3R being slightly more efficient. However, we presented one case, a 2D Kelvin-Helmholtz instability, where the more detailed HLL5R was significantly less viscous. This was because of velocity shears parallel to both the grid and the magnetic field lines. In a three-dimensional turbulence run at Mach 2, we also found HLL5R to be somewhat less dissipative than HLL3R, giving results that were very close to the Roe solver-based FLASH version. The new FLASH MHD module is freely available upon contact with the corresponding author." }, "1101/1101.4627_arXiv.txt": { "abstract": "{ Radio mini-halos are diffuse, steep-spectrum synchrotron sources associated with a fraction of relaxed clusters of galaxies. Observations of some mini-halo sources indicate a correlation between the radio emission and the X-ray signature of gas sloshing, ``cold fronts.'' Some authors have suggested turbulence associated with the sloshing motions may reaccelerate relativistic electrons, resulting in emission associated with the fronts. We present MHD simulations of core gas sloshing in a galaxy cluster, where we measure the turbulence created by these motions and employ passive tracer particles to act as relativistic electrons that may be reaccelerated by such turbulence. Our preliminary results support such a link between sloshing motions and particle reacceleration. } ", "introduction": "Many galaxy clusters are sources of radio emission. One such class of sources found in galaxy clusters are that of radio mini-halos. Mini-halos are diffuse, steep-spectrum synchrotron sources found in the cores of so-called ``cool-core'' clusters. These sources typically are associated with a central AGN and extend out to approximately the cooling radius of the cluster gas \\citep[for a review see][]{fer08}. A number of mini-halos have emission that is correlated on the sky with spiral-shaped ``cold fronts'' seen in the X-ray emission, believed to be the signature of sloshing of the cluster's cool core gas \\citep{MVF03,asc06}. \\citet{maz08} discovered this correlation in two clusters, and suggested that the correlation resulted from a population of relativistic electrons that was reaccelerated by turbulence generated by the sloshing motions. In order to determine whether or not the reacceleration efficiency resulting from this turbulence is sufficient enough to reaccelerate electrons and produce the corresponding radio emission, we have performed MHD simulations of gas sloshing in a galaxy cluster with tracer particles acting as the relativistic electrons. \\begin{figure*}[t!] \\resizebox{\\hsize}{!}{\\includegraphics[clip=true]{temp.eps}\\includegraphics[clip=true]{bmag.eps}} \\caption{\\footnotesize Slices through the center of the simulation domain of gas sloshing in the center of a galaxy cluster. Left: Temperature (keV), showing prominent spiral-shaped ``cold fronts''. Right: Magnetic field strength ($\\mu$G), showing the field amplification near the front surfaces. Vectors show the magnetic field direction. Each panel is 400~kpc on a side. } \\label{fig:sloshing} \\end{figure*} ", "conclusions": "We have performed MHD simulations of gas sloshing in clusters of galaxies, with the aim of determining whether or not the correlation between radio mini-halo emission and sloshing cold fronts in some clusters can be explained by the reacceleration of relativistic electrons by turbulence associated with the sloshing motions. Our initial results are very promising, indicating that the combined effects of the amplified magnetic field and the turbulence associated with the sloshing motions reaccelerates a population of relativistic electrons within the envelope of the cold fronts which emit synchrotron radiation from these regions. In our procedure we assume a population of seed electrons diffused in the sloshing region. Although we do not model the injection process of these particles, several mechanisms may provide an efficient source of fresh electrons to reaccelerate in cluster cool-cores \\citep[see discussion in][]{cas08}. Further work will also detail differences in prediction with other models for mini-halos." }, "1101/1101.4970_arXiv.txt": { "abstract": "Close pairs of white dwarfs are potential progenitors of Type~Ia supernovae and they are common, with of order 100 -- 300 million in the Galaxy. As such they will be significant, probably dominant, sources of the gravitational waves detectable by \\emph{LISA}. In the context of \\emph{LISA}'s goals for fundamental physics, double white dwarfs are a source of noise, but from an astrophysical perspective, they are of considerable interest in their own right. In this paper I discuss our current knowledge of double white dwarfs and their close relatives (and possible descendants) the AM~CVn stars. \\emph{LISA} will add to our knowledge of these systems by providing the following unique constraints: (i) an almost direct measurement of the Galactic merger rate of DWDs from the detection of short period systems and their period evolution, (ii) an accurate and precise normalisation of binary evolution models at the shortest periods, (iii) a determination of the evolutionary pathways to the formation of AM~CVn stars, (iv) measurements of the influence of tidal coupling in white dwarfs and its significance for stabilising mass transfer, and (v) discovery of numerous examples of eclipsing white dwarfs with the potential for optical follow-up to test models of white dwarfs. ", "introduction": "In the early 1980s it was suggested that Type~Ia supernovae might come from close pairs of white dwarfs merging under the action gravitational radiation losses \\cite{Webbink1984,IbenTutukov1984}. It was later realised that the large number of systems needed to sustain the Type~Ia rate within the Galaxy under these models meant that double white dwarfs (henceforth DWDs) are likely to be a dominant source of gravitational waves for space-based interferometry, to the extent that over some frequency intervals of interest in the context of \\emph{LISA}, DWDs may define \\emph{LISA}'s noise floor \\cite{Evans1987,Hils1990}. Early searches for DWDs produced meagre returns, and predictions that 10\\% of all ``single'' white dwarfs might in fact be double \\cite{Paczynski1985} seemed wide of the mark. \\citeasnoun{Robinson1987} found no DWDs amongst 44 targets, \\citeasnoun{Foss1991} none amongst 25, and \\citeasnoun{Bragaglia1990} found one certain DWD together with a few candidates from 54 targets. Together with the system L870-2, discovered by \\citeasnoun{Saffer1988}, by the early 1990s only two DWDs had measured periods. This changed when advances in our understanding of white dwarf atmospheres led to the identification of white dwarfs of too low a mass for single star evolution \\cite{Bergeron1992}. Optical spectroscopy showed that a large fraction of these objects are DWDs \\cite{Marsh:WD1101+364,Marsh:friends,Holberg:Feige55,Moran:WD0957-666,Maxted:WD1704+481}. Since the turn of the millennium further discoveries have followed from the SPY survey \\cite{Napiwotzki:SPY2} and from the SDSS as detailed later. These discoveries have established the presence of large numbers of DWDs and their importance as gravitational wave sources. There have been numerous studies of the likely impact of DWDs on \\emph{LISA}. These find that below a cutoff frequency of about $2$ to $6\\,$mHz, there are so many systems that their signals are unresolved, while above this frequency individual systems are resolved, with the odd nearby system rising above the noise at somewhat lower frequencies \\cite{Hils1990,Hils2000,Nelemans2001,Ruiter2010,Liu2010,Yu2010}. Most of these studies have been concerned with predicting the gravitational wave (GW) signal from DWDs in \\emph{LISA}. My interest here is more what we can learn about DWDs from \\emph{LISA} that is hard to deduce from electromagnetic (EM) observations. The potential is great, with direct measurement of tidal coupling between white dwarfs and the first detections of DWDs in globular clusters where their numbers are expected to be dynamically enhanced \\cite{Shara2002}, likely to come from \\emph{LISA} data. ", "conclusions": "Double white dwarfs are predicted to be the dominant source population at \\emph{LISA} frequencies. \\emph{LISA}'s sensitivity to short orbital periods will allow the best estimates of the merger rates of these stars, tidal coupling of the two stars and answer questions about their evolution that are hard to solve at the longer periods favoured by electromagnetic observations. The dual combination of the GW and EM observations will be a powerful tool for probing white dwarf astrophysics." }, "1101/1101.1849_arXiv.txt": { "abstract": "{} {The abundances of Fe in the intracluster medium of nearby ($z<0.08$) clusters were measured up to 0.3$\\sim$ 0.5r$_{180}$. }{We analyzed 28 clusters of galaxies observed with XMM-Newton. We derived Fe abundances from the flux ratios of Fe lines to the continuum within an energy range of 3.5--6 keV to minimize and evaluate systematic uncertainties. }{ The radial profiles of the Fe abundances of relaxed clusters with a cD galaxy at their X-ray peak have similar slopes. These clusters show similar enhancements in the Fe abundance within 0.1$r_{180}$, and at 0.1--0.3$r_{180}$, they have flatter Fe abundance profiles at 0.4$\\sim$0.5 solar, with a small scatter. Most other clusters, including merging clusters, also have similar Fe abundance profiles beyond 0.1$r_{180}$. { These clusters may have universal metal enrichment histories}, \\rm and a significant amount of Fe was synthesized at a very early stage in cluster formation. Mergers of clusters can destroy the central Fe peak. } {} ", "introduction": "Clusters of galaxies are the largest gravitationally bound objects in the universe. The intracluster medium (ICM) contains a large amount of metals, which are synthesized mainly by supernovae (SNe) in early-type galaxies (e.g., Arnaud et al. 1992; Renzini et al. 1993). Thus, metal abundances in the ICM provide important clues for understanding the metal-enrichment history and evolution of galaxies in clusters. Because both SN II and SN Ia synthesize Fe, the distribution of Fe in the ICM contains information about the star-formation histories of massive stars and the history of chemical-enrichment attributed to SN Ia. The ASCA satellite (Tanaka et al.\\ 1994) first enabled us to measure the distribution of Fe in the ICM\\@ (e.g., Fukazawa et al.\\ 2000; Finoguenov et al.\\ 2000; 2001). The Fe abundances of these clusters are 0.2--0.3 solar, adopting the solar abundance from the ``photospheric'' values given by Anders and Grevesse (1989). The dependence of the Fe abundance on the temperature of the ICM is weak (Fukazawa et al. 1998). ASCA revealed large-scale abundance gradients from AWM 7 and the Perseus cluster (Ezawa et al. 1997; 2001). De Grandi et al. (2004) derived the Fe abundances of nearby hot clusters observed with Beppo-SAX and found that the abundance profiles are strongly peaked for cool core clusters, whereas they remain constant for other systems. XMM-Newton and Chandra observations show a spatial distribution of the Fe abundance of up to $0.3\\sim 0.4 r_{180}$ (e.g., Tamura et al. 2004; Vikhlinin et al. 2005; Baldi et al. 2007; Maughan et al. 2008; Leccardi \\& Molendi 2008). At the center of most relaxed clusters, the Fe abundance decreases steeply outward. Outside the central region, the Fe abundance % decreases more gradually toward the outer regions. The Suzaku satellite first measured the Fe abundance of the ICM beyond $0.5r_{180}$. (Fujita et al. 2008; Tawa 2008). Within 0.3--0.5$r_{180}$, Suzaku can derive the Fe abundances more accurately than XMM\\@. The Fe abundance gradually decreases from the center to $\\sim $0.7$r_{180}$. In this paper, we describe our study of the Fe abundance in the ICM up to $0.3\\sim 0.5$$r_{180}$ of 28 brightest clusters of galaxies observed with XMM-Newton. Some clusters have cool cores at their center (e.g. Makishima et al. 2001; B\\\"ohringer et al. 2002). We must be careful when deriving elemental abundances with a multi-temperature plasma, because emission lines and a continuum spectrum are different functions of temperature. Therefore, we directly derived the strength of the K$\\alpha$ lines of Fe, considering the temperature dependence of the flux ratios of the Fe lines and the continuum, and derived Fe abundances. We adopted an energy range in which the ICM dominates the background. In Section 2 we summarize the observations and data preparation. Section 3 describes our analysis of the data, and in Section 4 the Fe abundances are determined. We discuss our results in Secion 5. We adopt the new solar abundances given by Lodders (2003), according to which the solar Fe abundance with regard to H is 2.95$\\times 10^{-5}$by number. This value differs from the photospheric value (4.68$\\times 10^{-5}$) given by Anders and Grevesse (1989). \\color{black} Considering the difference in He abundance between the two solar abundance tables, the Fe abundance yielded by the former is 1.5 times higher than that of the latter. \\color{black} We use H$_0$ = 70 km/s/Mpc. Unless otherwise specified, errors are quoted with 68\\% confidence. ", "conclusions": "We derived radial profiles of the Fe abundance of the ICM in nearby ($z<0.08$) clusters observed with XMM-Newton. The Fe abundances of the ICM were derived from the ratio of the flux of the K$\\alpha$ lines of He-like or H-like Fe to those of the continuum at 3.5--6.0 keV, because the systematic uncertainty in the continuum flux in this energy band owing to the background is smaller. The temperature dependence of these ratios constrains the Fe abundances of multi-temperature plasmas. \\color{black} In cluster core regions ($<0.1 r_{180}$), the observed Fe abundances of cD clusters show similar radial profiles. The less-peaked abundance profile \\color{black} compared with the light \\color{black} in the central region indicates the ejection of metals from cD galaxies. \\color{black} In the outer regions, 0.1--0.2$r_{180}$, Fe abundances of 0.4--0.5 solar appear to be universal with no temperature dependence. The observed flatter Fe abundance profiles of the cD clusters beyond 0.1$r_{180}$ indicate early metal enrichment. Chandra and XMM cannot reach beyond $0.5 r_{180}$ owing to a high, unstable particle background. Information is available for only about 10\\% of the cluster volume, and the majority has not yet been revealed. The total amount of Fe synthesized in galaxies can be derived only by the precise abundance measurements beyond $0.5 r_{180}$. Because of its low background, Suzaku is the only satellite available for the next several years to study clusters of galaxies up to the virial radius. With Suzaku, Fe abundances will be determined out to a radius of 0.7$\\sim$ 0.8$r_{180}$." }, "1101/1101.4188_arXiv.txt": { "abstract": "The phenomenon of turbulent thermal diffusion in temperature-stratified turbulence causing a non-diffusive turbulent flux of inertial and non-inertial particles in the direction of the turbulent heat flux is found using direct numerical simulations (DNS). In simulations with and without gravity, this phenomenon is found to cause a peak in the particle number density around the minimum of the mean fluid temperature for Stokes numbers less than 1, where the Stokes number is the ratio of particle Stokes time to turbulent Kolmogorov time at the viscous scale. Turbulent thermal diffusion causes the formation of large-scale inhomogeneities in the spatial distribution of inertial particles. The strength of this effect is maximum for Stokes numbers around unity, and decreases again for larger values. The dynamics of inertial particles is studied using Lagrangian modelling in forced temperature-stratified turbulence, whereas non-inertial particles and the fluid are described using DNS in an Eulerian framework. ", "introduction": "Transport and mixing of small particles (aerosols and droplets) in turbulent fluid flow is of fundamental importance in a large variety of applications (environmental sciences, physics of the atmosphere and meteorology, industrial turbulent flows and turbulent combustion; see, e.g., \\cite{CSA80,BLA97,CST11,WA00,BE10,S03,BH03,KPE07}). There are also astrophysical applications, in particular in the context of protoplanetary accretion discs \\cite{SA72,TA09,HB98,JAB04}. Numerous laboratory \\cite{WA00,AC02,WH05,SA08,XB08,SSA08} and numerical \\cite{BL04,CC05,CG06,YG07,AC08,BB07,BB10} experiments as well as observations in atmospheric \\cite{S03,KPE07,VY00,SGG10} and astrophysical \\cite{SA72,TA09,HB98,JAB04,JY07,Johansen07} turbulent flows have shown different kinds of large-scale and small-scale long-living inhomogeneities (clusters) in the spatial distribution of particles. It is well known that turbulent diffusion causes destruction of large-scale inhomogeneities in the spatial distributions of particles. But how can we explain the opposite process resulting in a formation of large-scale clusters of particles? One of the mechanisms of formation of particle inhomogeneities in temperature-stratified turbulence is the phenomenon of turbulent {\\it thermal} diffusion \\cite{elperin_etal96}. This effect consists of a turbulent non-diffusive flux of inertial particles in the direction of the turbulent heat flux, so that particles are accumulated in the vicinity of the mean temperature minimum. The particular form of the flow field does not play any role in this effect. It is a purely collective phenomenon caused by temperature stratified turbulence resulting in a pumping effect, i.e., appearance of the non-zero mean effective velocity of particles in the direction opposite to the mean temperature gradient. A competition between two different phenomena, namely the turbulent thermal diffusion and the turbulent diffusion determines the conditions for the formation of large-scale particle clusters in the vicinity of the mean temperature minimum. The characteristic scale of the particle inhomogeneity formed due to the turbulent thermal diffusion is much larger than the integral scale of the turbulence. Furthermore, the characteristic time scale of the formation of the particle inhomogeneity is much longer than the characteristic turbulent time, i.e., this is a mean-field effect. The phenomenon of turbulent thermal diffusion has been predicted theoretically in \\cite{elperin_etal96} and detected in different laboratory experiments in stably and unstably temperature-stratified turbulence \\cite{EEKR04,BEE04,EEKR06a,eidelman_etal06,EKR10}. This phenomenon is shown to be important for atmospheric turbulence with temperature inversions \\cite{sofiev09} and for small-scale particle clustering in temperature-stratified turbulence \\cite{EKR10}, but it is also expected to be significant for different kinds of heat exchangers, e.g., industrial boilers where Reynolds numbers and temperature gradients are large. In spite of the fact that turbulent thermal diffusion has already been found in different types of laboratory experiments and atmospheric flows, this effect has never been observed in direct numerical simulations. The main goal of this paper is to find turbulent thermal diffusion of non-inertial and inertial particles in direct numerical simulations (DNS). The paper is organized as follows. In Sect.\\ II we discuss the physics of the phenomenon of turbulent thermal diffusion. The numerical simulations for fluid, inertial and non-inertial particles, and the results of direct numerical simulations are described in Sect.\\ III. Motions of inertial particles are determined using a Lagrangian framework (Sections III-B,C,D), while non-inertial particles are described using an Eulerian framework (Sect.\\ III-E). Conclusions are drawn in Sect.\\ IV. ", "conclusions": "This study is the first numerical demonstration of the existence of the phenomenon of turbulent thermal diffusion of inertial and non-inertial particles in forced, temperature stratified turbulence. The inertial particles are described using a Lagrangian framework, while non-inertial particles and the fluid flow are determined using an Eulerian framework. The phenomenon of turbulent thermal diffusion has been studied for different Stokes and fluid Reynolds numbers, P\\'eclet numbers as well as different forcing scales of the turbulence. Furthermore, the effect of gravity has been included in simulations. In all simulations, with different parameters and different formulations of the dynamic viscosity, we always observe the effect of particle accumulation in the vicinity of the mean temperature minimum due to turbulent thermal diffusion for $\\St<1$. This effect is robust and the results of the numerical simulations are in agreement with theoretical studies \\cite{elperin_etal96,EKR00}, laboratory experiments \\cite{BEE04,eidelman_etal06,EKR10} and atmospheric observations \\cite{sofiev09}. When $\\St>1$, this effect is decreasing with Stokes number." }, "1101/1101.1764_arXiv.txt": { "abstract": "% {Since its formulation by Hamaker et al., the radio interferometer measurement equation (RIME) has provided a rigorous mathematical basis for the development of novel calibration methods and techniques, including various approaches to the problem of direction-dependent effects (DDEs). However, acceptance of the RIME in the radio astronomical community at large has been slow, which is partially due to the limited availability of software to exploit its power, and the sparsity of practical results. This needs to change urgently.} {This series of papers aims to place recent developments in the treatment of DDEs into one RIME-based mathematical framework, and to demonstrate the ease with which the various effects can be described and understood. It also aims to show the benefits of a RIME-based approach to calibration. } {Paper I re-derives the RIME from first principles, extends the formalism to the full-sky case, and incorporates DDEs. Paper II then uses the formalism to describe self-calibration, both with a full RIME, and with the approximate equations of older software packages, and shows how this is affected by DDEs. It also gives an overview of real-life DDEs and proposed methods of dealing with them. Finally, in Paper III some of these methods are exercised to achieve an extremely high-dynamic range calibration of WSRT observations of 3C 147 at 21 cm, with full treatment of DDEs. }% {The RIME formalism is extended to the full-sky case (Paper I), and is shown to be an elegant way of describing calibration and DDEs (Paper II). Applying this to WSRT data (Paper III) results in a noise-limited image of the field around 3C 147 with a very high dynamic range (1.6 million), and none of the off-axis artifacts that plague regular selfcal. The resulting differential gain solutions contain significant information on DDEs and errors in the sky model. }% {The RIME is a powerful formalism for describing radio interferometry, and underpins the development of novel calibration methods, in particular those dealing with DDEs. One of these is the differential gains approach used for the 3C 147 reduction. Differential gains can eliminate DDE-related artifacts, and provide information for iterative improvements of sky models. Perhaps most importantly, sources as faint as 2 mJy have been shown to yield meaningful differential gain solutions, and thus can be used as potential calibration beacons in other DDE-related schemes.} ", "introduction": " ", "conclusions": "Since its original formulation by \\citet{ME1}, the Radio Interferometer Measurement Equation (RIME) has provided the mathematical underpinnings for novel calibration methods and algorithms. Besides its explanatory power, the RIME formalism can be wonderfully simple and intuitive; this fact has become somewhat obscured by the many different directions that it has been taken in. Several authors have developed approaches to the DDE problem based on the RIME, using different (but mathematically equivalent) versions of the formalism. This paper has attempted to reformulate these using one consistent $2\\times2$ formalism, in preparation for follow-up papers (II and III) that will put it to work. Finally, a number of misunderstandings and controversies has inevitably accrued themselves to the RIME over the years. Some of these have been addressed here. It is hoped that this paper has gone some way to making the RIME simple again." }, "1101/1101.4231_arXiv.txt": { "abstract": "Juvenile ultracool dwarfs are late spectral type objects (later than $\\sim$M6) with ages between 10~Myr and several 100 Myr. Their age-related properties lie intermediate between very low mass objects in nearby star-forming regions (ages 1--5 Myr) and field stars and brown dwarfs that are members of the disk population (ages 1--5~Gyr). Kinematic associations of nearby young stars with ages from $\\sim$10--100~Myr provide sources for juvenile ultracool dwarfs. The lowest mass confirmed members of these groups are late-M dwarfs. Several apparently young L dwarfs and a few T dwarfs are known, but they have not been kinematically associated with any groups. Normalizing the field IMF to the high mass population of these groups suggests that more low mass (mainly late-M and possibly L dwarf) members have yet to be found. The lowest mass members of these groups, along with low mass companions to known young stars, provide benchmark objects with which spectroscopic age indicators for juvenile ultracool dwarfs can be calibrated and evaluated. In this proceeding, we summarize currently used methods for identifying juvenile ultracool dwarfs and discuss the appropriateness and reliability of the most commonly used age indicators. ", "introduction": "Juvenile ultracool dwarfs are very low mass stellar and substellar objects. {\\em Juvenile} refers to objects with intermediate ages (10~Myr~$<$~age~$\\lesssim$~600~Myr). They lack substantial ongoing accretion and primordial circumstellar material, but they still exhibit some signatures of youth that are not seen in typical field objects. {\\em Ultracool} dwarfs have a spectral type of $\\sim$M6 or later. Juvenile ultracool dwarfs are typically identified by the combination of a late spectral type with one or more youth indicators: activity signatures, low gravity spectral features, membership in a young cluster or nearby moving group, and/or companionship to a known young star. A small but significant population of these objects are currently known, including such benchmark objects as 2MASS~J1207$-$39, a member of the $\\sim$10~Myr~TW~Hydrae moving group and the host of a planetary-mass companion. The properties of juvenile ultracool dwarfs play a role in many aspects of star formation and stellar evolution. A complete census of the low mass population of young, nearby moving groups is essential for understanding how the initial mass function varies across stellar environments. Characterization of the physical and circumstellar properties of juvenile ultracool dwarfs is crucial for complete understanding of any evolutionary phenomenon with a mass or age dependence, for example: planet formation, disk dissipation, angular momentum evolution, companion frequency, and chromospheric activity. Benchmark juvenile ultracool dwarfs (i.e. objects with well-characterized kinematic and physical properties) will essentially provide calibration data for evolutionary models. Finally, juvenile ultracool dwarfs provide excellent targets for exoplanet searches because they are nearby and young, thus potentially hosting self-luminous giant planets that provide a favorable contrast ratio and angular separation for direct imaging instruments \\citep{Beichman10,Kataria10}. The specific questions about juvenile ultracool dwarfs addressed in the splinter session were: \\begin{enumerate} \\item What is the most efficient and accurate method for identifying juvenile ultracool dwarfs and associating them with young nearby moving groups? \\item What properties/features are reliable age indicators for late-M, L, and T spectral types? \\item How do juvenile ultracool dwarfs fit in with our current understanding of star formation, e.g., mass function, number density, multiplicity, disk fraction, etc.? \\end{enumerate} The first question is addressed in Section~2 and the second in Section~3. Question 3 is not explicitly discussed in this proceeding. As a result of the splinter session it became clear that a more complete answer to 1 and 2 will further our understanding of point 3. Section~4 discusses important caveats for identifying young moving groups, evaluating membership, and using membership as an age indicator. ", "conclusions": "" }, "1101/1101.1622_arXiv.txt": { "abstract": "The degeneracy among the disk, bulge and halo contributions to galaxy rotation curves prevents an understanding of the distribution of baryons and dark matter in disk galaxies. In an attempt to break this degeneracy, we present an analysis of the strong gravitational lens \\lens, discovered as part of the SLACS survey. The lens galaxy is a high inclination, disk dominated system. We present new Hubble Space Telescope multicolor imaging, gas and stellar kinematics data derived from long-slit spectroscopy, and K-band laser guide star adaptive optics imaging, both from the Keck telescopes. We model the galaxy as a sum of concentric axisymmetric bulge, disk and halo components and infer the contribution of each component, using information from gravitational lensing and gas kinematics. This analysis yields a best-fitting total (disk plus bulge) stellar mass of $\\log_{10}(\\Mstar/\\Msun)=10.99^{+0.11}_{-0.25}$. The photometric data combined with stellar population synthesis models yield $\\log_{10}(\\Mstar/\\Msun)=10.97\\pm 0.07$, and $11.21\\pm 0.07$ for the Chabrier and Salpeter IMFs, respectively. Assuming no cold gas, a Salpeter IMF is marginally disfavored, with a Bayes factor of 2.7. Accounting for the expected gas fraction of $\\simeq 20\\%$ reduces the lensing plus kinematics stellar mass by $0.10\\pm0.05$ dex, resulting in a Bayes factor of 11.9 in favor of a Chabrier IMF. The dark matter halo is roughly spherical, with minor to major axis ratio $\\qhalo=0.91^{+0.15}_{-0.13}$. The dark matter halo has a maximum circular velocity of $V_{\\rm max}=276^{+17}_{-18} \\kms$, and a central density parameter of $\\log_{10}\\Delta_{V/2}=5.9^{+0.9}_{-0.5}$. This is higher than predicted for uncontracted dark matter haloes in $\\Lambda$CDM cosmologies, $\\log_{10}\\Delta_{V/2}=5.2$, suggesting that either the halo has contracted in response to galaxy formation, or that the halo has a higher than average concentration. Larger samples of spiral galaxy strong gravitational lenses are needed in order to distinguish between these two possibilities. At 2.2 disk scale lengths the dark matter fraction is $f_{\\rm DM}=0.55^{+0.20}_{-0.15}$, suggesting that \\lens is sub-maximal. ", "introduction": "\\label{sec:intro} \\begin{figure*} \\centering\\includegraphics[width=0.9\\linewidth]{fig1.eps} \\caption{Illustration of the different geometries probed by strong lensing and kinematics. Strong lensing measures mass with a cylinder (or more generally an ellipse), whereas stellar and gas kinematics measure mass within spheres (or more generally ellipsoids).} \\label{fig:cartoon} \\end{figure*} The discovery of extended flat rotation curves in the outer parts of disk galaxies three decades ago (Bosma 1978; Rubin \\etal 1978) was decisive in ushering in the paradigm shift that led to the now standard cosmological model dominated by cold dark matter (CDM). The need for dark matter on cosmological scales is also firmly established from observations of the Cosmic Microwave Background, type Ia Supernovae, weak lensing, and galaxy clustering (see, e.g., Spergel \\etal 2007). Numerical simulations of structure formation within the $\\Lambda$CDM cosmology make firm predictions for the structure and mass function of dark matter haloes in the absence of baryons (e.g., Navarro, Frenk, \\& White 1997; Bullock \\etal 2001; Macci\\`o \\etal 2007; Navarro et al.\\ 2010). It is still unclear, however, whether this standard model can reproduce the observed properties of the Universe at galactic and sub-galactic scales. There are problems related to the inner density profiles of dark matter haloes (e.g., de Blok \\etal 2001; Swaters \\etal 2003; Newman et al. 2009), reproducing the zero point of the Tully-Fisher relation (e.g., Mo \\& Mao 2000; Dutton \\etal 2007), and the amount of small-scale substructure (e.g., Klypin \\etal 1999; Moore et al. 1999; Stewart \\etal 2008). There are three classes of solutions to these problems: those that invoke galaxy formation processes that modify the properties of dark matter haloes; those that change the nature of dark matter itself; and those in which dark matter does not exist. Thus, measuring the density profiles of the dark matter haloes of galaxies of all types is a stringent test for galaxy formation theories. From an observational point of view, little is known about the detailed distribution of dark matter in the inner regions of disk galaxies, despite the great investment of telescope time and high quality measurements of hundreds of rotation curves (e.g., Carignan \\& Freeman 1985; Begeman 1987; Courteau 1997; de Blok \\& McGaugh 1997; Verheijen 1997; Swaters 1999; de Blok \\etal 2001; Swaters \\etal 2003; Blais-Ouellette \\etal 2004; Simon \\etal 2005; Noordermeer \\etal 2005; Simon \\etal 2005; Chemin \\etal 2006; Kuzio de Naray 2006; de Blok \\etal 2008; Dicaire \\etal 2008; Epinat \\etal 2008). The fundamental reason is the so-called {\\it disk-halo degeneracy}: mass models with either maximal or minimal baryonic components fit the rotation curves equally well, leaving the structure of the dark matter halo poorly constrained by the kinematic data alone (e.g., van Albada \\& Sancisi 1986; van den Bosch \\& Swaters 2001; Dutton \\etal 2005). Stellar population models are able to place constraints on stellar mass-to-light ratios, allowing inference about the baryonic contribution to the overall mass profile. However, there are a number of uncertainties which limit the accuracy of this method (e.g., Conroy \\etal 2009, 2010). These include systematic uncertainties such as the unknown stellar initial mass function (IMF), and the treatment of the various stellar evolutionary phases in stellar population synthesis (SPS) models. These result in about a factor of 2 uncertainty in the stellar masses estimated from spectral energy distribution (SED) fitting. Moreover, for a given IMF and SPS model, there are uncertainties in the star formation histories, metallicities and extinction which introduce ($1\\sigma$) random errors in measurements of stellar masses for individual galaxies at the level of 0.15~dex (e.g., Bell \\& de Jong 2001; Auger \\etal 2009, Gallazzi \\& Bell 2009). Nevertheless, galaxy colours and dynamical mass estimates have been used in combination by various authors to place an upper limit on the stellar mass-to-light ratio normalisation, favoring IMFs more bottom light than Salpeter for spiral galaxies and fast rotating low-mass elliptical galaxies (Bell \\& de Jong 2001; Cappellari \\etal 2006; de Jong \\& Bell 2007; see, however, Treu et al. 2010, Auger et al. 2010, and van Dokkum \\& Conroy 2010 for massive ellipticals). However, as theory generally predicts more dark matter in the inner regions of disk galaxies than is consistent with standard IMFs (Dutton \\etal 2007; Dutton \\etal 2010b) a lower limit to the stellar mass would provide a more useful constraint for $\\Lambda$CDM. Several other methods have been used to try and measure disk galaxy stellar masses, independent of the uncertainties in the IMF. These include: 1) vertical velocity dispersions of low inclination disks (Bottema 1993; Verheijen \\etal 2007; Bershady \\etal 2010), 2) bars and spiral structure (Weiner \\etal 2001; Kranz \\etal 2003), and 3) strong gravitational lensing by inclined disks (Maller \\etal 2000; Winn \\etal 2003). None of these methods have thus far yielded conclusive results. \\begin{figure*} \\centering\\includegraphics[width=0.42\\linewidth] {fig2a.eps} \\centering\\includegraphics[width=0.42\\linewidth] {fig2b.eps} \\caption{Differences between projected (cylindrical) mass and enclosed (spherical) mass for a bulge-halo system (left) and a disk-halo system (right). For each system two models are shown (in red and black). The models have baryonic mass profiles (short-dashed lines, upper panels) with the same shape but normalizations that differ by a factor of two. The dark matter profiles (long-dashed lines, upper panels) have been chosen so that the total circular velocity curves are close to identical (solid lines, upper panels). For the bulge-halo system the ratio between projected and enclosed masses (middle panels) is independent of the relative contributions of the bulge and halo, which differ significantly between the two models (lower panels). However, for the disk-halo system there is a significant difference between the projected and enclosed masses, especially at radii smaller than the effective radius. This illustrates the potential of strong lensing plus kinematics to break the disk-halo degeneracy.} \\label{fig:proj} \\end{figure*} An approach combining strong gravitational lensing plus kinematics holds great promise, because it takes advantage of the different geometries of disks and haloes, which results in three effects that enable the disk mass to be measured. 1) An inclined disk will present a much higher projected surface density than a face-on disk, with resulting image positions and shapes that depend on the disk mass fraction. 2) An edge-on disk is highly elliptical in projection, more than expected for any realistic dark matter halo, with resulting total mass ellipticity depending on the disk mass fraction 3) Strong lensing measures mass projected along a cylinder (within the Einstein radius), whereas stellar kinematics (rotation and dispersion) measure mass enclosed within spheres (see \\Fref{fig:cartoon}). For spherical mass distributions of stars and dark matter, the ratio between the projected mass within a cylinder of radius, $r$, and the enclosed mass within a sphere of the same radius, $r$, is {\\it independent} of the relative contribution of the two mass components (left panel \\Fref{fig:proj}). Therefore, in order to break the degeneracy one has to assume a radial profile shape for both components (e.g., Treu \\& Koopmans 2002, 2004; Koopmans \\& Treu 2003; Koopmans et al. 2006; Treu et al.\\ 2010; Auger et al. 2010). Typically this involves assuming the baryonic mass follows the light, and then assuming a functional form for the dark matter halo. However, for a disk plus halo system, this ratio is {\\it dependent} on the relative contribution of the two components (right panel \\Fref{fig:proj}). Thus if the spherical and cylindrical masses can be measured accurately enough, the disk halo degeneracy can be broken without assuming a specific radial profile shape for either component. Furthermore, strong lensing plus kinematics can place constraints on the 3D shape of the dark matter halo (e.g., Koopmans, de Bruyn, \\& Jackson 1998; Maller \\etal 2000) which is of interest because $\\Lambda$CDM haloes are predicted to be non-spherical (e.g. Allgood \\etal 2006; Bett \\etal 2007; Macci\\`o \\etal 2008). The power of the strong gravitational lensing method has not yet been fully realised, primarily due to the scarcity of known spiral galaxy gravitational lenses. Prior to the SLACS Survey (Bolton \\etal 2006, 2008) only a handful of spiral galaxy lenses with suitable inclinations to enable rotation curve measurements were known: Q2237$+$0305 (Huchra \\etal 1985; Trott \\& Webster 2002); B1600$+$434 (Jackson \\etal 1995; Jaunsen \\& Hjorth 1997); PMN\\,J2004$-$1349 (Winn \\etal 2003); CXOCY\\,J220132.8$-$320144 (Castander \\etal 2006). However, most of these systems are doubly-imaged QSOs which provide minimal constraints on the projected mass density. Q2237$+$0305 is a quadruply-imaged QSO, which gives more robust constraints, but since the Einstein radius is small compared to the size of the galaxy, the lensing is mostly sensitive to the bulge mass, not the halo (Trott et al. 2010; van de Ven et al. 2010). The final SLACS lens sample (Auger \\etal 2009) is comprised of 98 strong galaxy-galaxy lenses, among these, 16 have been classified morphologically as type S or S0. Inspired by this, we have extended the Sloan Digital Sky Survey (hereafter SDSS, York \\etal 2000) spectroscopic lens selection technique specifically to spiral galaxy lenses. In the resulting SWELLS survey (Treu \\etal, in prep, referred to hereafter as Paper I) we have assembled a larger sample of \\NSWELLS \\, late-type galaxy-scale gravitational lenses for detailed mass modelling. In this paper, the second of the SWELLS series, we present a detailed and self-consistent mass model of the spiral galaxy lens \\lens (RA=21:41:54.67, DEC=$-$00:01:12.2, J2000), constrained by both kinematic and lensing data. As we will see, this galaxy is disk-dominated, with a disk inclination of $\\simeq 80^\\circ$; this set-up approximately maximises the projected disk mass while allowing an accurate rotation curve to be measured. The original spectroscopic observations of \\lens were obtained on SDSS plate~989, with fiber~35, on MJD~52468. The latest public SDSS-DR7 (Abazajian \\etal 2009) Petrosian magnitudes (uncorrected for extinction) for the lens galaxy are $(u,g,r,i,z)=(20.61,18.62,17.47,16.92,16.48)$ with errors $(0.15,0.01,0.01,0.01,0.02)$. The SDSS measured redshift for the lens galaxy is $\\zd=0.1380 \\pm 0.00015$, and the velocity dispersion is $181\\pm14\\kms$. The spectrum also exhibits nebular emission lines at a background redshift of $\\zs=0.7127$ (Bolton \\etal 2008). With these redshifts the scale in the lens plane is $1\\,{\\rm arcsec}=2.438 \\rm \\,kpc$, while in the source plane it is $1\\,{\\rm arcsec}=7.196 \\rm \\,kpc$. This paper is organised as follows. In \\Sref{sec:imaging} we present the imaging observations of \\lens from Keck and the Hubble Space Telescope, and then infer the structure of the stellar mass distribution of the galaxy in the presence of its dust from these data in \\Sref{sec:mstar}. With this information in hand we then define a three-component mass model for the galaxy in \\Sref{sec:model}, and describe how it is constrained by the imaging data (although we choose not to use the stellar mass inferred from the SED at this stage). In \\Sref{sec:lens}, we describe the preparation and analysis of the strong lensing data. In \\Sref{sec:spec} we present the spectroscopic observations of \\lens from Keck. Then in \\Sref{sec:results} we present fits to the lensing and kinematics data using three combinations: lensing only; kinematics only; and lensing plus kinematics. This joint analysis yields constraints on the stellar mass of the disk and bulge; returning to the stellar masses inferred from the stellar population modelling of the SED we, discuss implications of our results for the stellar IMF in \\Sref{sec:imf}. In \\Sref{sec:halodensity} \\& \\Sref{sec:haloshape} we discuss our results for the density and shape of the dark matter halo. We conclude in \\Sref{sec:concl}. Throughout, we assume a flat $\\Lambda$CDM cosmology with present day matter density, $\\Omega_{\\rm m}=0.3$, and Hubble parameter, $H_0=70 \\rm\\,km\\,s^{-1}\\,Mpc$. All magnitudes are given in the AB system. Unless otherwise stated, all parameter estimates are the median of the marginalised posterior PDF, and their uncertainties are described by the absolute difference between the median and the 84th and 16th percentiles (such that the error bars enclose 68\\% of the posterior probability). \\begin{figure*} \\centering\\includegraphics[width=0.95\\linewidth] {fig3.eps} \\caption{Optical to near-IR high resolution imaging of \\lens. Images are 12~arcsec by 6~arcsec, North up, East is left. The lens is a high inclination, disk dominated, and star forming spiral galaxy. The source ($\\simeq 1.3$ arcsec to the east of the galaxy center) appears to be multiply imaged in the optical, but is a continuous arc in the near IR.} \\label{fig:images} \\end{figure*} ", "conclusions": " \\begin{itemize} \\item The lensing and kinematics constraints yield a stellar mass of $\\log_{10}(\\Mstar/\\Msun) = 10.99^{+0.11}_{-0.25}$ (68\\% confidence interval), independent of the IMF. \\item This value is in excellent agreement with the stellar mass derived from the SED using SPS models and assuming a Chabrier (2003) IMF: $\\log_{10}(\\Mstar/\\Msun) = 10.97^{+0.07}_{-0.07}$. A Salpeter (1955) IMF results in stellar masses 0.24 dex higher: our analysis marginally favors a Chabrier IMF over a Salpeter IMF, by a Bayes factor of 2.7. \\item Accounting for the expected gas mass reduces the lensing and kinematics stellar mass by $0.10\\pm0.05$ dex, and increases the Bayes factor in favor of a Chabrier IMF to 11.9. \\item At 2.2 disk scale lengths the spherical dark matter fraction is $f_{\\rm DM}=0.55^{+0.20}_{-0.15}$, suggesting that the baryons are sub-maximal. \\item The dark matter halo has a maximum circular velocity of $\\Vhalo=276^{+17}_{-18} \\kms$, and a core radius of $\\rhalo=2.4^{+2.4}_{-1.5} \\kpc$. The corresponding central density parameter $\\log_{10}\\Delta_{V/2}=5.9^{+0.9}_{-0.5}$ is higher than expected for uncontracted NFW haloes in the concordance $\\Lambda$CDM cosmology, which have $\\log_{10}\\Delta_{V/2}=5.2$ and an intrinsic scatter of 0.3. \\item This high density could either be evidence for halo contraction in response to galaxy formation (e.g., Blumenthal \\etal 1986), or the result of a selection bias towards high concentration haloes. A larger sample with well-characterised selection function is required to make further progress. \\item The dark matter halo is oblate, $\\qhalo = 0.91^{+0.15}_{-0.13}$, with a probability of 69\\%. This finding provides support for the notion that galaxy assembly turns strongly prolate triaxial dark matter haloes into roughly oblate axisymmetric haloes (e.g., Abadi \\etal 2010). \\end{itemize}" }, "1101/1101.1308_arXiv.txt": { "abstract": "These notes present a detailed introduction to Maldacena's calculation [1] of the cubic terms in the inflationary action. These interactions are important since they produce the most readily observable evidence for a non-Gaussian component in the pattern of primordial fluctuations produced by inflation. In the simplest class of inflationary theories, those with only a single scalar field participating in the inflationary era, these non-Gaussianities are predicted to be extremely small, as will be reviewed here. ", "introduction": " ", "conclusions": "" }, "1101/1101.3633_arXiv.txt": { "abstract": "Recent studies suggest that \\emph{Swift} gamma-ray bursts (GRBs) may not trace an ordinary star formation history. Here we show that the GRB rate turns out to be consistent with the star formation history with an evolving stellar initial mass function (IMF). We first show that the latest \\emph{Swift} sample of GRBs reveals an increasing evolution in the GRB rate relative to the ordinary star formation rate at high redshifts. We then assume only massive stars with masses greater than the critical value to produce GRBs, and use an evolving stellar IMF suggested by Dav\\'{e} (2010) to fit the latest GRB redshift distribution. This evolving IMF would increase the relative number of massive stars, which could lead to more GRB explosions at high redshifts. We find that the evolving IMF can well reproduce the observed redshift distribution of \\emph{Swift} GRBs. ", "introduction": "Gamma-ray bursts (GRBs) are brief flashes of $\\gamma$-rays occurring at an average detection rate of a few events per day at cosmological distances. Because of their very high luminosity, GRBs can be detected out to the edge of the visible Universe (Ciardi \\& Loeb 2000; Lamb \\& Reichart 2000; Bromm \\& Loeb 2002; Gou et al. 2004). Thus, GRBs are ideal tools for probing the star formation rate, the reionization history, and the metal enrichment history of the Universe (Totani 1997; Campana et al. 2007; Bromm \\& Loeb 2007). The advantages of GRBs over quasars for probing the high-redshift Universe had been discussed by Bromm \\& Loeb (2007). In addition, GRBs have been used as standard candles to constrain cosmological parameters and dark energy (Dai, Liang \\& Xu 2004; Friedman \\& Bloom 2005; Wang \\& Dai 2006; Schaefer 2007, and references therein). The association of long GRBs with core-collapse supernovae naturally suggests that the cosmic GRB rate should trace the star formation history. This gave rise to the expectation that GRBs may be a good tracer of cosmic star formation (Totani 1997; Wijers et al. 1998; Lamb \\& Reichart 2000; Blain \\& Natarajan 2000; Porciani \\& Madau 2001). However, it was found that the rate of GRBs increases with cosmic redshift faster than the ordinary star formation rate (SFR) does (Daigne et al. 2006; Le \\& Dermer 2007; Kistler et al. 2008, 2009; Y\\\"{u}ksel \\& Kistler 2007; Cen \\& Fang 2007; Li 2008; Wang \\& Dai 2009; Butler et al. 2010; Wanderman \\& Piran 2010). The reason for this discrepancy has been unknown. By investigating the redshift distribution of \\emph{Swift} GRBs, Guetta \\& Piran (2007) found that the observed high-redshift bursts are more than the expectation from an ordinary star formation history (SFH) and thus the high-redshift GRB rate is inconsistent with the one inferred from the current model for the SFR. Furthermore, Kistler et al. (2008) found that the GRB rate at redshift $z\\simeq 4$ is about four times larger than expected from star formation measurements. Daigne et al. (2006) concluded that GRB properties or progenitors must evolve with cosmic redshift to reconcile the observed GRB redshift distribution with the measured SFH. Li (2008) explained the observed discrepancy between the GRB rate history and the star formation rate history as being due to cosmic metallicity evolution, by assuming that long GRBs tend to occur in galaxies with low metallicities. However, very recently Levesque et al. (2010a,b) found several high-metallicity long GRB host environments, which suggests that a low-metallicity cut-off is unlikely (also see Graham et al. 2009). Xu \\& Wei (2008) used a factitious stellar initial mass function (IMF) evolving with redshift to interpret the GRB redshift distribution. Cheng et al. (2010) suggested that this discrepancy could be eliminated if some high-redshift GRBs are ascribed to electromagnetic bursts of superconducting cosmic strings, although the existence of the superconducting cosmic strings has remained controversial. In this Letter, we first enlarge the GRB sample with 122 long GRBs observed by \\emph{Swift}. Then we interpret the latest \\emph{Swift} GRB redshift distribution using a reasonable evolving stellar initial mass function (IMF) proposed by Dav\\'{e} (2010). The structure of this paper is as follows: in section 2, we give an evolving initial mass function form, and in section 3, we show the analysis method. The results are presented in section 4 and conclusions are shown in section 5. ", "conclusions": "In this Letter, we have presented that the redshift distribution of \\emph{Swift} GRBs with measured redshifts and calculated luminosities can be successfully fitted by the SFH with an evolving stellar IMF. It is widely considered by current theories that only massive stars with masses larger than the critical value can produce long GRBs. The evolving stellar IMF becoming increasingly top heavy at larger $z$ suggested by Dav\\'{e} (2010) can lead to more GRBs produced at high redshifts. Kistler et al. (2008) considered several possible reasons for the discrepancy between the \\emph{Swift} GRB rate and the SFH. They showed that the Kolmogorov-Smirnov test does not favor an interpretation as a statistical anomaly. Selection effects are also unlikely to cause an increased efficiency in detecting high-redshift GRBs. Although Kistler et al. (2008) have argued that alternative reasons are possible (e.g., evolution in the fraction of binary systems, an evolving IMF of stars, cosmic metallicity evolution), they did not give a quantitative analysis or a detailed discussion of the evolving IMF. We enlarged the GRB sample with 122 long GRBs and used a reasonable evolving IMF. The results in this paper indicate that the evolving IMF may explain the redshift distribution of \\emph{Swift} GRBs. If the redshift distribution of GRBs and SFH are well measured, GRBs would be used to probe the stellar IMF at high redshifts." }, "1101/1101.4225_arXiv.txt": { "abstract": "There are strong correlations between the three structural properties of elliptical galaxies -- stellar mass, velocity dispersion and size -- in the form of a tight ``fundamental plane\" and a ``scaling relation\" between each pair. Major mergers of disk galaxies are assumed to be a mechanism for producing ellipticals, but {\\bf semi-analytic galaxy formation models (SAM) have} encountered apparent difficulties in reproducing the observed slope and scatter of the size-mass relation. We study the scaling relations of merger remnants using progenitor properties from two SAMs. We apply a simple merger model that includes gas dissipation and star formation based on theoretical considerations and simulations. {\\bf Combining the SAMs and the merger model allows calculation of the structural properties of the remnants of major mergers that enter the population of elliptical galaxies at a given redshift.} Without tuning the merger model parameters for each SAM, the results roughly match the slope and scatter in the observed scaling relations and their evolution in the redshift range $z=0-3$. Within this model, the observed scaling relations, including the tilt of the fundamental plane relative to the virial plane, result primarily from the decrease of gas fraction with increasing progenitor mass. The scatter in the size-mass relation {\\bf of the remnants} is {\\bf reduced from that of the progenitors} because of a correlation between progenitor size and gas fraction at a given mass. ", "introduction": "\\label{sec:intro2} The merging of disk galaxies is one of the main hypothesized mechanisms for the formation of elliptical galaxies. Simulations have shown that mergers of disks with similar masses can effectively disrupt the ordered rotation in the disks and convert it into random velocity support, creating merger remnants that appear similar to observed elliptical galaxies \\citep{TT72, T77, Barnes92, MH94dsc}. Furthermore, the $\\Lambda$CDM cosmology predicts the hierarchical buildup of galaxies through a sequence of mergers. These results suggest that merging is a likely mechanism for the production of elliptical galaxies. However, observed ellipticals follow a number of scaling relations, including relatively tight relations between stellar mass and velocity dispersion, the Faber-Jackson relation \\citep{FJ}, and between size and stellar mass \\citep{Kormendy77}. Furthermore, observed ellipticals fall on a tight plane, the fundamental plane (FP), in the three-dimensional space of stellar mass, size, and velocity dispersion \\citep{Djorgovski87, Dressler87}. Recent studies of the Sloan Digital Sky Survey (SDSS) have provided excellent statistics on these scaling relations in the local universe \\citep{Shen03, Bernardi03a, Bernardi03b, Padmanabhan04, Gallazzi06, Shankar10b}, and studies using high redshift surveys have provided evidence for the evolution of these relations over cosmological time \\citep{Barden05, McIntosh05, Trujillo06,Trujillo07, Cimatti08,vanderWel08,Buitrago08, Saracco09, vanDokkum10, Fan10, Mancini10, Williams10}. If mergers are a major mechanism for the production of elliptical galaxies then they must be able to produce the correct scaling relations as well as the evolution of these scaling relations over time. Theoretical {\\bf high-resolution, hydrodynamical} studies have shown that simulations of gas-rich galaxy mergers are capable of reproducing the observed scaling relations of elliptical galaxies if the correct progenitor properties are used \\citep{Dekel06, RobertsonFP, HopkinsFP,Bournaud:2011a}. {\\bf These studies have been successful in creating high-redshift compact ellipticals from major mergers of gas-rich compact disk galaxies \\citep{Bournaud:2010a, Wuyts:2010b}.} This is a step toward verifying the production of scaling laws through mergers. However, current computing power only allows the simulation of relatively small numbers of mergers, and the space of possible merger initial conditions and progenitor properties is quite large. Furthermore, these simulations are not placed within a cosmological context, making it more difficult to explore in detail the origin and evolution of scaling relations. Currently, the primary theoretical tool for studying the evolution of statistical samples of galaxies over cosmological timescales is semi-analytic modeling (SAM) \\citep{kwg1993, Cole94, SP1999, Cole00, Galics03, Croton06, DeLucia06, Bower06, S08, Fontanot09, Benson10, BensonBower10, Cook10, Guo10}. These models combine dark matter halo merger trees with analytic recipes for populating the halos with galaxies. However, semi-analytic models (SAMs) do not currently incorporate realistic formulas for predicting the properties of the remnants of galaxy mergers including the effects of dissipation. The agreement between current SAMs and observed early-type scaling relations is not impressive. In particular, the observed size-mass relation of early-types is steeper than that of their potential late-type progenitors \\citep{Shen03}, and the scatter in the observed size-mass relation for early-types is remarkably small \\citep{Shen03, Nair10}. The disspationless merger models currently implemented within SAMs have thus far been unable to reproduce these features \\citep[e.g.,][]{Shankar10a,Guo10}. {\\bf Using a simple power law dissipation model \\citet{KhochfarSilk06} were able to reproduce the redshift-size evolution of elliptical galaxies, which, along with the high-resolution hydrodynamical simulations described above, suggests that dissipative effects may play an important role in determining elliptical galaxy scaling relations.} \\citet{remnants} recently developed a physically-motivated analytic model for predicting the stellar half-mass radii and central velocity dispersions of merger remnants \\citep{remnants}. The parameters in this new merger model were calibrated using a suite of {\\bf high-resolution hydrodynamical} galaxy merger simulations (see Section \\ref{sec:methods2}). Here we implement a simplified version of this model using post-processing of merger outputs from the SAMs developed by \\citet{Croton06}, based on the Millennium simulation \\citep{Springel05}, and \\citet{S08}. This results in a population of tens of thousands of merger remnants over a large range of redshifts ($0 R_{\\rm vac}$, where $R_{\\rm vac} = 12H_{0}^{2}\\Omega_{\\Lambda}$ is the cosmological vacuum curvature at the present time. Examples of viable $f(R)$ functions are \\cite{Hu:2007nk,st} \\begin{eqnarray} \\label{eq:a1}& & f(R) = -m^{2} {c_{1} (R/m^{2})^{2n} \\over 1+ c_{2}(R/m^{2})^{2n}} \\\\ \\label{eq:a2} & & f(R) = \\lambda R_{st} \\left[ \\left( 1 + {R^{2} \\over R_{st}^{2}}\\right)^{-d} - 1 \\right] \\end{eqnarray} \\noindent both of which can be expanded as \\begin{eqnarray} \\label{eq:i1} & & f(R) = -{R_{\\rm vac} \\over 2} + \\lambda R_{\\rm vac} \\left({R_{\\rm vac} \\over R} \\right)^{2s} + {\\cal O} \\left( \\lambda^{2} \\left({R_{\\rm vac} \\over R} \\right)^{4s} \\right) \\end{eqnarray} \\noindent for $R \\gg R_{\\rm vac}$, where $\\lambda$ and $s$ are the modified gravity parameters nontrivially related to the model specific parameters $m^{2},c_{1},c_{2},R_{st}, n,d$ (see also \\cite{ap,ts10,Cognola:2007zu,Nojiri:2006be,Fay:2007uy,Linder:2009jz,Bamba:2010ws,Nojiri:2008nt,Park:2010da,Elizalde:2010ts,Appleby:2009uf}). The fact that these functions reduce to expansions around a constant in regions of high curvature allow them to evade solar system tests and also reproduce the standard background cosmology with radiation and matter dominated epochs \\cite{Amendola:2006kh,Amendola:2006we}.\\footnote{The functions ($\\ref{eq:a1},\\ref{eq:a2}$) in fact still possess potential fine tuning issues in the very early Universe, see for example \\cite{st}. However, we only consider redshifts $z \\lesssim 1000$ and assume that these issues are ameliorated by modifying the $f(R)$ functional form at high energies.} Observational constraints on $f(R)$ models place lower bounds on the mass of the scalar field, at the curvature scale for which the observation is performed. In this work we will be concerned with cosmological constraints that can be placed on $f(R)$ models and associated modified gravity parameterisations. One can broadly classify cosmological probes in two categories: distance based probes (such as SNIa), which utilize the background evolution of the Hubble parameter, and structure based probes (such as weak lensing and galaxy clustering), which require knowledge of the evolution and scale dependence of density perturbations. \\\\ \\\\ \\noindent The paper proceeds as follows: \\\\ \\\\ In section \\ref{sec:1} we review the background evolution and perturbation equations for $f(R)$ models under the so called `quasi-static' approximation and explore a simple parameterisation to encapsulate them. In section \\ref{sec:paramaterisations} we show that the scalaron parameterisation that we initially describe is equivalent to common parameterisations of the Poisson equation $\\mu(a,k)$ and anisotropic stress $\\eta(a,k)$. Then, by comparing between scalaron models we highlight these forms to be less flexible than at first thought. We therefore suggest an improved parameterisation, which accurately represents viable $f(R)$ models in the literature. We discuss potential modified gravity signals in the expansion history in \\ref{sec:lum}, and in section \\ref{sec:3} we consider modified gravity effects in the CMB angular power spectrum and the linear and non linear matter power spectrum. In section \\ref{sec:4} and \\ref{sec:5} we examine how future CMB and weak lensing surveys (which probe the matter power spectrum) will constrain such models and argue for their complementarity to solar system constraints. We discuss our analysis and conclude in section \\ref{sec:6}. ", "conclusions": "\\label{sec:6} We have examined a general scheme for a class of viable gravitational $f(R)$ models; the scalaron framework. This is physically motivated and gives rise to distinctive observational features. For viable $f(R)$ models the background expansion is nearly identical to $\\Lambda$CDM, a fact reflected in a luminosity distance-redshift relationship that is practically indistinguishable from the standard cosmology. By solving the modified perturbation equations with an augmented {\\sc camb}, we also find that the CMB does not add information on the modified signal directly, because the scalaron mass is too large to affect the ISW substantially. The most significant observational effect for these models occurs in the matter power spectrum; specifically $P(a,k)$ acquires a redshift and scale dependent modification relative to GR. This signal can be used to constrain modified gravity parameters with weak lensing. By examining parameterisations of $f(R)$ gravity in detail we have argued that existing functional forms in the literature do not capture the essential features of $f(R)$ models at late times $z \\lesssim 1$. To accurately describe the behaviour of these models for $z \\lesssim 1$ we have directly parameterized the scalaron mass $M(a)$. Whilst our approach can be related to the standard $\\mu(a,k)$ and $\\eta(a,k)$ parameterization in the literature \\cite{Martinelli10,Bertschinger08,Zhao:2008bn,Daniel10} (and also the $\\gamma$ parameterization; see \\cite{Appleby:2010dx,Motohashi:2010tb,Fu:2010td}), the mass of the scalar field is an ideal function to parameterize, in the sense that it is independent of scale, only weakly dependent on cosmological parameters and has a clear physical interpretation. Using the general scalaron model ($\\ref{eq:107}$), we found that a future weak lensing probe such as Euclid will be able to place extremely tight constraints on deviations to gravity. We quote our final results as $M_{0} = 1.34 \\pm 0.62 \\times 10^{-30} [{\\rm h \\hspace{1mm} eV}] $, $\\nu = 1.50 \\pm 0.18$ for $l < 400$ and $M_{0} = 1.34 \\pm 0.25 \\times 10^{-30} [{\\rm h \\hspace{1mm} eV}] $, $\\nu = 1.50 \\pm 0.04$ for the $l < 10000$ (Planck + Euclid) combined analysis. Our forecast results highlight a number of interesting points, specifically; the importance of constraining these models with cosmological weak lensing data (compared to solar system tests), the value of using CMB data to break parameter degeneracies, and also the importance of modeling the nonlinear power spectrum. As stated previously, this regime is particularly hard to treat in any gravitational theory, and could bias the results if implemented blindly. While some approaches, such as higher order perturbation theory, can give us an indication of physics at nonlinear scales, it is evident that there is a huge demand for programmes of simulations and a better approach to fitting to them. A detailed study of this problem will be presented elsewhere. \\vspace{5mm}" }, "1101/1101.0630_arXiv.txt": { "abstract": "PTF/M-dwarfs is a 100,000-target M-dwarf planetary transit survey, a Key Project of the Palomar Transient Factory (PTF) collaboration. The survey is sensitive to Jupiter-radius planets around all of the target stars, and has sufficient precision to reach Neptunes and super-Earths for the best targets. The Palomar Transient Factory is a fully-automated, wide-field survey aimed at a systematic exploration of the optical transient sky. The survey is performed using a new 7.26 square degree camera installed on the 48 inch Samuel Oschin telescope at Palomar Observatory. Each 92-megapixel R-band exposure contains about 3,000 M-dwarfs usable for planet detection. In each PTF observational season PTF/M-dwarfs searches for Jupiter-radius planets around almost 30,000 M-dwarfs, Neptune-radius planets around approximately 500 M-dwarfs, and super-Earths around 100 targets. The full survey is expected to cover more than 100,000 targets over the next several years. Photometric and spectroscopic followup operations are performed on the Palomar 60-inch, LCOGT, Palomar 200-inch, MDM and Keck telescopes. The survey has been running since mid-2009. We detail the survey design, the survey's data analysis pipeline and the performance of the first year of operations. ", "introduction": "Thus far only a handful of planetary systems have been detected around stars with masses $< 0.5\\rm{M_{\\odot}}$, and very few lower-mass M-dwarfs have been searched for planets. The higher-mass M-dwarfs that have been probed up to now have a low Jupiter-mass planet companion frequency ($\\sim$2\\%; eg. \\citet{Johnson2007}). However, theorists predict a large population of lower-mass planets around all masses of M-dwarfs \\citep{Ida2005, Kennedy2008}, and recent microlensing and transit detection results (e.g. \\citet{Sumi2010, Charbonneau2009}) suggest that superEarths and mini-Neptunes may be very common around sub-solar-mass stars. PTF/M-dwarfs is a new transit detection survey targeted at 100,000 cool stars. The survey is a key project of the Palomar Transient Factory (PTF) collaboration (\\citet{Law2009} \\& \\citet{Rau2009}). PTF is a new wide field, multiple cadence transient search hosted at the 1.2m Samuel Oschin Telescope at Palomar Observatory (hereafter referred to as P48). Commissioning of the new 7.26-square-degree camera system was completed in August 2009, and PTF science operations are ongoing \\citep{Law2010}. \\begin{figure}[tb] \\plotone{law_n_fig1.eps} \\caption{\\label{ps}The cumulative PTF/M-dwarfs sample size in the first 9 months of operations, as a function of detectable planet radius. The planet sensitivity of each target is calculated on the basis of the achieved RMS photometric precision for that target and its photometrically-estimated spectral type. In this figure we assume that a planet detection requires a 2-sigma detection in each of at least three datapoints during a transit. This assumption is a conservative estimate that is improved with the utilization of full phase-folding detection techniques. Observation window effects are not included here.} \\end{figure} \\begin{figure}[tb] \\plotone{law_n_fig2.eps} \\caption{\\label{img}A typical PTF/M-dwarfs target field, with the planet-detection targets denoted by red points.} \\end{figure} ", "conclusions": "" }, "1101/1101.4319_arXiv.txt": { "abstract": "P Cygni is a prototype of the Luminous Blue Variables (or S Doradus variables), and the star displays photometric and emission line variability on a timescale of years (known as the ``short S Doradus phase'' variations). Here we present new high resolution H$\\alpha$ spectroscopy of P~Cyg that we combine with earlier spectra and concurrent $V$-band photometry to document the emission and continuum flux variations over a 24~y time span. We show that the emission and continuum fluxes vary in concert on timescales of 1.6~y and longer, but differ on shorter timescales. The H$\\alpha$ profile shape also varies on the photometric timescales, and we describe the observed co-variations of the emission peak and absorption trough properties. We argue that the episodes of photometric and emission brightening are caused by increases in the size of the emission region that are related to variations in wind mass loss rate and outflow speed. We find evidence of blueward accelerating, Discrete Absorption Components (DACs) in the absorption trough of the H$\\alpha$ profile, and these features have slower accelerations and longer durations than those observed in other lines. The DAC strengths also appear to vary on the photometric timescales, and we suggest that the propagation of the DAC-related wind structures is closely related to changes in the overall wind mass loss rate and velocity. ", "introduction": "Luminous Blue Variables (LBVs or S Doradus variables) are evolved, massive stars. LBVs are characterized by large mass loss rates and variability on multiple timescales. The two ``prototypical\" Galactic LBVs are $\\eta$ Carinae and P Cygni, and they probably represent different extremes of mass loss rate within the scheme of LBV evolution \\citep{IdG99}. One of the defining criteria of the LBVs is the observation of a large scale eruption, when the star brightens by several magnitudes. The quiescent times between these eruptions may last centuries. In addition to such rare, giant eruptions, these stars also display lesser photometric and spectroscopic variations on other timescales (e.g.~ Humphreys \\& Davidson 1994). \\citet{vg01} defines the S Doradus (SD-) phase to be the moderate, long-term, brightening and fading phases. There are two types of these phases, short and long, with similar characteristics. The short SD-phase is typically on the order of years ($\\lesssim$ 10 years), while the longer SD-phase is on a timescale of decades. These phases are thought to originate from changes in the star's photosphere, and both may have the same physical driving mechanism. These long-term variations are observed to differ from cycle to cycle, both in duration and amplitude \\citep{ster97}. P Cygni (HD 193237, HR 7763, Nova Cyg 1600) remains one of the most fascinating objects in the sky. It was discovered during its first recorded great eruption in 1600 by Willem Janszoon Blaeu, a Dutch chart-maker and mathematician. During this eruption, the star brightened to about 3rd magnitude for about six years, and then it faded from visibility by 1626. It rose again in 1654 to about the same maximal brightness, where it remained for five years. The star faded after this, and although its long-term variability is poorly documented, the star has been slowly brightening to its current magnitude of about 4.8 \\citep{IdG99}. The slow brightening may reflect evolutionary changes (de Groot \\& Lamers 1992; Lamers \\& de Groot 1992; Langer et al.\\ 1994). Long-term photometric monitoring of P Cygni began in the 1980s when \\citet{PW83}, \\citet{per88}, and \\citet{dG90} embarked on extended observing campaigns. Their observations showed that the variations often occur on three characteristic timescales: a short $\\sim$ 17 day variation similar to the $\\alpha$ Cygni type variations observed in hot supergiants, a $\\sim$ 100 day ``quasi-period\" similar to that observed in other LBVs, and a long-term cycle (years) attributed to a short-SD phase (de Groot et al.~2001; Percy et al.~2001). According to \\citet{IdG99}, comprehensive spectral monitoring of P Cygni was started by Luud (1967) and Markova (1993), among others. The first long-term spectroscopic monitoring campaign of P Cygni was presented in seminal papers by Markova et al.~(2001a,b). They found evidence of co-variability of the H$\\alpha$ emission line strength and Johnson $UBV$ photometry, indicating a short-SD phase with a quasi-period of $\\sim$7 years, although their observations did not fully cover two cycles. The variations were attributed to inversely correlated changes in effective temperature and radius, maintaining a nearly constant luminosity. A similar cycle time was found by \\citet{dG01}, who reported on photometric variations which were consistent with a timescale of 5.5 to 8.5 years. In addition to the large scale variations in emission strength, Markova (2000) found that there are at least four other kinds of line profile variability in the spectrum of P~Cygni. The most striking of these is the long documented appearance of blueward-migrating, absorption sub-features that are called Discrete Absorption Components (DACs: Israelian \\& de Groot 1999; Markova 2000). These are generally observed in low and intermediate excitation state lines in the optical (Markova 2000) and UV spectrum (Israelian et al.\\ 1996). They are frequently detected in the upper sequence of the H Balmer lines (principal quantum number $9\\le n \\le 15$; Markova 2000), but to our knowledge, DACs have not been reported before now for the absorption component of H$\\alpha$. DACs are often (but not always) narrow (FWHM $\\approx 10 - 15$ km~s$^{-1}$) and may be unresolved in low dispersion spectra. The DACs tend to appear over a radial velocity range of $-90$ to $-200$ km~s$^{-1}$ with an acceleration of $-0.1$ to $-0.6$ km~s$^{-1}$~d$^{-1}$. A recurrence timescale of $\\sim 200$~d is sometimes observed (Kolka 1983; Markova 1986a; Israelian et al. 1996; Kolka 1998). These accelerations are much slower and the timescales are much longer than those associated with DACs in the winds of O-stars (Kaper et al.\\ 1999). The DACs in the spectrum of P~Cygni may form in outward moving and dense shells (Kolka 1983; Lamers et al.\\ 1985; Markova 1986a; Israelian et al.\\ 1996), in spiral-shaped co-rotating interaction regions (CIRs; Cranmer \\& Owocki 1996; Markova 2000), or in dense clumps in the wind (L\\'{e}pine \\& Moffat 2008). In this paper, we present new high resolution H$\\alpha$ spectroscopy, which we combined with previous measurements by \\citet{mar1a} to explore the characteristics of P Cygni's short SD-phase. We also compare this with archival Johnson $V$ photometry and new observations obtained by AAVSO observers. Section 2 describes our observations. In Section 3, we present our analysis of long-term variations of the continuum and the H$\\alpha$ equivalent width. We describe the H$\\alpha$ profile morphology changes and DAC propagations in Section 4. Our discussion and conclusions are presented in Section 5. ", "conclusions": "The $V$-band and H$\\alpha$ variations we observe need to be interpreted in the context of current models for the star and its wind. Langer et al.\\ (1994) discuss the atmospheric properties of P~Cyg, and they argue that the star is in the LBV phase where the temperature and helium abundance are increasing, and the mass and luminosity are decreasing, as the star evolves towards the Wolf-Rayet phase. Langer et al.\\ emphasize the earlier conclusion from Pauldrach \\& Puls (1990) that the atmospheric parameters are close to a bi-stability point where the mass loss rate can change by an order of magnitude with small changes in radius and/or luminosity, which may explain the great eruptions observed in prior centuries. The atmospheric parameters are well established through a detailed quantitative spectroscopic analysis by Najarro et al.\\ (1997) and Najarro (2001), who find that He is overabundant and that the mass loss rate is high ($\\approx 2\\times 10^{-5}$ $M_\\odot$~y$^{-1}$ including wind clumping effects) and wind terminal velocity is low ($v_\\infty = 185$ km~s$^{-1}$). Najarro et al.\\ (1997) derive a systemic velocity of $\\gamma = -29$ km~s$^{-1}$, and thus our minimum measurement of $V_{\\rm min} = -215$ km~s$^{-1}$ is consistent with their estimate of $\\gamma - v_{\\infty} = -214$ km~s$^{-1}$ as this velocity measurement is related to $v_{\\infty}$. They estimate that the continuum forming radius is $76 R_\\odot$, which for a distance of 1.8~kpc implies an angular diameter of $\\theta = 0.39$ mas. On the other hand, Najarro et al.\\ (1997) predict that the emitting size of H$\\alpha$ will be much larger because of its greater optical depth. For example, their models show that there is a local maximum in the wind temperature distribution (presumably where the recombination processes that form H$\\alpha$ peak) near $r/R_\\star = 11$ (see their Fig.~5b). The corresponding angular size for H$\\alpha$ of $\\approx 4$~mas agrees well with the range of $3 - 7$~mas from H$\\alpha$ interferometry by Balan et al.\\ (2010). Thus, we need to keep in mind that the H$\\alpha$ variations reflect changes over a much larger spatial scale in the wind than those observed in the $V$-band flux. Variations in the H$\\alpha$ emission equivalent width are related to changes in both the mass loss rate and the wind velocity. In a very simplified approach, we can assume that most of the H$\\alpha$ flux originates in the optically thick region projected on the sky, $$f = 2 \\pi r_\\tau ^2 F(T)$$ were $r_\\tau$ is the boundary separating the optically thick and thin regimes, $F(T)$ is the monochromatic surface flux, and $T$ is the wind temperature at $r_\\tau$ (Najarro et al.\\ 1997). If we assume that the wind is approximately isothermal at this physical location (a reasonable choice: see Fig.~5b in Najarro et al.\\ 1997), then the emission flux variations are due to changes in the projected size of the optically thick region, $$\\triangle f / f = 2 \\triangle r_\\tau / r_\\tau.$$ Thus, we expect that the relative variations in angular size will be only half as large as the emission equivalent width variations, which is probably consistent with the lack of measurable size changes in the H$\\alpha$ interferometric measurements (Balan et al.\\ 2010). The H$\\alpha$ optical depth is dependent on the electron density squared since the emission is a recombination process. Thus, we expect that the optical depth unity boundary $r_\\tau$ will always be defined by the location in the wind with a specific characteristic density, $\\rho_\\tau$. We assume that $\\rho_\\tau$ has an approximately constant value so that the effective $r_\\tau$ boundary will vary as fluctuations in the wind mass loss rate and velocity define the radius where the density reaches $\\rho_\\tau$. According to the mass continuity equation, $r_\\tau$ is related to this density by $$r_\\tau ^2 = {\\dot{M}\\over {4 \\pi \\rho_\\tau v}}$$ where $\\dot{M}$ is the mass loss rate and $v$ is the wind velocity at the radial distance $r_\\tau$. We can differentiate the mass continuity equation to express the radius variation in terms of the changes in $\\dot{M}$ and $v$, $$2 r_\\tau \\triangle r_\\tau = {1\\over {4 \\pi \\rho_\\tau}} \\triangle [{\\dot{M}/ v}],$$ which we divide by $r_\\tau ^2$ to obtain $$2 {\\triangle r_\\tau \\over r_\\tau} = {{\\triangle [{\\dot{M}/ v}]} \\over {[{\\dot{M}/ v}]}}.$$ Since we argued above that the flux also varies as $r_\\tau ^2$, we can then use the relation above to re-write the fractional flux variation in terms of logarithmic changes in $\\dot{M}$ and $v$, $$\\triangle \\ln f = \\triangle \\ln \\dot{M} - \\triangle \\ln v .$$ Since we have observational data on the variations in emission strength and wind velocity, we can rearrange this equation to solve for the mass loss variations as a function of flux variations, $$\\triangle \\ln \\dot{M} / \\triangle \\ln f = 1 + \\triangle \\ln v /\\triangle \\ln f .$$ Puls et al.\\ (1996) present a much more detailed analysis of the dependence of the emission equivalent width on the wind parameters of hot, massive stars. However, in the limit of high optical depth, their expression for the emission flux (their eq.~41) leads to a similar relation, $$\\triangle \\ln \\dot{M} / \\triangle \\ln f = {3\\over4} (1 + \\triangle \\ln v /\\triangle \\ln f) .$$ We found in the previous section that the H$\\alpha$ equivalent width appears to vary inversely with two quantities related to wind dynamics, the H$\\alpha$ emission peak FWHM and the blue minimum flux velocity $V_r$(min). Figure 10 quantifies this relationship. The upper panel shows the inverse correlation between the emission peak FWHM and flux corrected H$\\alpha$ equivalent width, and if we take FWHM as a proxy for the wind speed, then a linear fit of natural logarithms of these measures gives $\\triangle \\ln v / \\triangle \\ln f = -0.66 \\pm 0.12$. We caution that the FWHM is also influenced by the absorption component of H$\\alpha$, and we showed above (Fig.~4) that the absorption component moves inward towards the line core when the emission is strong. Consequently, the apparent decrease in FWHM as the emission increases probably results both from a wind speed decrease and a blue wing decline due to blending with the absorption component. The minimum flux velocity is perhaps a more direct measurement of wind speed (at least in our line of sight), and we show in the lower panel of Figure~10 the co-variations of the difference between $V_r$(min) and the systemic velocity of P~Cyg, $\\gamma = -29$ km~s$^{-1}$ (Najarro et al.\\ 1997), as a function of the corrected equivalent width. A fit of the logarithmic slope here gives a smaller estimate of $\\triangle \\ln v / \\triangle \\ln f = -0.22 \\pm 0.04$. \\placefigure{fig10} If we adopt the minimum flux co-variation result as representative of the wind velocity component of variability, then the mass loss rate variation we derive from the relation above has an emission flux dependence of $\\triangle \\ln \\dot{M} \\approx 0.78 \\triangle \\ln f$. Omitting the bottom and top 10\\% of the distribution of $W_\\lambda$(corr), the derived range in emission strength in our observations of $\\pm 14 \\%$ probably implies mass loss rate changes of $\\pm 11\\%$. Markova et al.\\ (2001a) used the optically thick relation from Puls et al.\\ (1996) to arrive at an estimate of $\\pm 9\\%$ for the mass-loss variation amplitude. We showed above that the relation from Puls et al.\\ carries a factor of $3/4$ that is missing from our simple analysis, and if we use the Puls et al.\\ relation instead, we arrive at a mass loss rate variation of $\\pm 8\\%$, confirming the estimate from Markova et al.\\ (2001a). In this scenario, the H$\\alpha$ emission variations result from changes in the effective emission radius $r_\\tau$ caused by variations in the mass loss rate and wind velocity. During episodes when the mass loss rate is higher and the wind is slower, the projected size of the emission region increases leading to larger H$\\alpha$ emission flux. The same process probably causes the $V$-band variations, but the fractional radius variations must be smaller at the continuum forming radius because we found in last section that $\\triangle \\ln f [V] / \\triangle \\ln f [{{\\rm H}\\alpha}] = 0.16$ so that $\\triangle \\ln r_{\\tau V} = 0.16 \\triangle \\ln r_{\\tau {\\rm H}\\alpha}$, i.e., the continuum size variations are only $16\\%$ as large as those in H$\\alpha$. It is possible that the changes in the mass-loss rate are caused by a change in the luminosity of the star which would propagate through the wind, and be observed in both the $V$-band brightness and the emission line flux of H$\\alpha$. Finally we return to the relationship of the DACs to the short SD-phase variations. We found a trend with the DAC strength and the H$\\alpha$ emission equivalent width. We show in Figure~11 the temporal behavior of DAC quotient equivalent width $W_\\lambda$(DAC) (Table 1, column 10) along with scaled, running averages of the H$\\alpha$ equivalent width and $V$-band flux (Fig.~3). We rescaled the amplitude of the H$\\alpha$ flux by $\\triangle W_{\\lambda} {\\rm (DAC)} = \\triangle W_{\\lambda}{\\rm (corr)} / 20.46$ and the photometric light curve was rescaled by $\\triangle W_{\\lambda} {\\rm (DAC)} = \\triangle f_{V}\\times 22.56$. These curves were then shifted vertically to match the $W_{\\lambda}({\\rm DAC})$ points. We see that the DAC strength variations track both the H$\\alpha$ and $V$-band flux variations, suggesting that the DACs are related in some way to the short SD-phase changes. For example, we see that some of the strongest DACs were observed when the H$\\alpha$ emission was strong (around HJD 2,452,500; MJD 5.25$\\cdot 10^4$ in our plots) and the DACs were very faint or absent when the emission was weak (around HJD~2,450,500). \\placefigure{fig11} The DAC phenomenon is primarily observed in the UV wind lines of hot stars (Kaper et al.\\ 1999; Puls et al.\\ 2008), and, in fact, DACs are only rarely seen in the H$\\alpha$ profile where wind-related variations are usually due to changes in the dense and slower moving wind close to the star (Kaper et al.\\ 1997; Markova et al.\\ 2005). The only comparable DAC observed in H$\\alpha$ was in HD 92207 (Kaufer et al.\\ 1996) with a lifetime of $\\sim 150$ d and only one DAC observed, so the recurrence time is unknown. This star is cooler than P Cygni and is a ``normal\" supergiant, compared to the luminous blue variable nature of P Cyg. The DACs observed in the UV wind lines of O-stars are first seen at velocities of 0.2 to 0.4 $v_\\infty$ and then they migrate blueward to $v_\\infty$ on time scales of a day or so, exhibiting accelerations that are much slower than expected for the wind velocity law (Kaper et al.\\ 1999). Many of these same features are seen in the DACs in H$\\alpha$ for P~Cygni, although they occur on vastly longer time scales. For example, the wind flow time scales as $R_\\star / v_\\infty$, and while the wind gas will accelerate from $0.1 v_\\infty$ to $0.9 v_\\infty$ in 0.5~d for an O-star like $\\xi$~Per (Kaper et al.\\ 1999), it will take some 43~d for the wind of P~Cygni. However, this flow time is very short compared to the longevity of the DACs ($10^3$~d), so we are led to the same conclusion found for the O-stars, namely that the DACs represent some kind of perturbation in the wind through which the gas flows. Changes in wind velocity and/or mass loss rate can cause shocks and create structure in the wind, and these structures produce density enhancements and/or velocity plateaus that imprint DACs in the wind lines (Fullerton \\& Owocki 1992; Runacres \\& Owocki 2002; Puls et al.\\ 2008). Current theory suggests that the DACs are related to large scale spiral features in the wind known as co-rotating interaction regions that are formed at the intersection of fast and slower outflows, which develop from some inhomogeneity in the stellar photosphere (for example, pulsation or spots; Cranmer \\& Owocki 1996; Lobel \\& Blomme 2008). In these models, it is the slow transit of these equatorially centered regions across the photosphere that creates the DACs in the absorption cores but has little influence on the emission parts of the wind line (Dessart 2004; Lobel \\& Blomme 2008). However, in the case of P~Cygni, we find that emission parts do appear to strengthen when DACs are prominent (Fig.~11), and this indicates that the wind perturbation profoundly affects both wind gas surrounding the star and the wind gas projected against the photosphere. Thus, we suggest that the structures causing the DACs in P~Cygni may be more spherically symmetric than assumed in the geometry of the co-rotating interaction regions. In some models the seed perturbation occurs at a fixed longitude on the star, so that a new wind structure appears each time the star rotates (although for the best studied case of HD~64760, Lobel \\& Blomme 2008 argue that the originating spots must rotate some five times slower than the star in order to fit models to the observations). If this is the case for P~Cygni, then the 1700~d DAC recurrence time may be related to the star's rotational period. Markova (2000) found an upper limit of 100~d for the rotational period based upon estimates of the stellar radius and the projected rotational velocity. However, the line broadening in early supergiants may be dominated by turbulence rather than rotation (Howarth et al.~1997, 2007; Markova \\& Puls 2008), so a longer rotational period remains a possibility. However, regardless of the origin of the DACs, their close relation with the emission line and continuum flux variations (Fig.~11) suggests that much of the short-SD variability is caused by propagating structural perturbations in the outer atmosphere of the star. The discovery of the short SD-phase variations in P Cygni and its relationship to the wind velocity and DAC occurrence in the wind is a new observational result that warrants future investigation both for this star and other LBVs. The changing characteristics over these long timescales may eventually lead to a better understanding of the LBV stage of evolution and the underlying physics of their winds and circumstellar environments. We are currently pursuing a three year, spectroscopic monitoring program of Galactic and Magellanic Cloud LBVs. Such long-term observations will reveal if DACs are present in all LBVs and will show whether the variations found in P~Cygni are a general phenomena among LBVs." }, "1101/1101.4405_arXiv.txt": { "abstract": "Scalar-tensor (ST) gravity theories provide an appropriate theoretical framework for the variation of Newton's fundamental constant, conveyed by the dynamics of a scalar-field non-minimally coupled to the space-time geometry. The experimental scrutiny of scalar-tensor gravity theories has led to a detailed analysis of their post-newtonian features, and is encapsulated into the so-called parametrised post-newtonian formalism (PPN). Of course this approach can only be applied whenever there is a newtonian limit, and the latter is related to the GR solution that is generalized by a given ST solution under consideration. This procedure thus assumes two hypothesis: On the one hand, that there should be a weak field limit of the GR solution; On the other hand that the latter corresponds to the limit case of given ST solution. In the present work we consider a ST solution with negative spatial curvature. It generalizes a general relativistic solution known as being of a degenerate class (A) for its unusual properties. In particular, the GR solution does not exhibit the usual weak field limit in the region where the gravitational field is static. The absence of a weak field limit for the hyperbolic GR solution means that such limit is also absent for comparison with the ST solution, and thus one cannot barely apply the PPN formalism. We therefore analyse the properties of the hyperbolic ST solution, and discuss the question o defining a generalised newtonian limit both for the GR solution and for the purpose of contrasting it with the ST solution. This contributes a basic framework to build up a parametrised pseudo-newtonian formalism adequate to test ST negatively curved space-times. ", "introduction": "\\label{sec:1} The possibility that physics might differ in diverse epochs and/or places in the universe is a question of paramount importance to understand what are the limits of our present physical laws~\\cite{Will:2005va,Damour:2002vu,Martins:2009zz,Barrow:2005hw}. This issue is at present very much at the forefront of the debate in gravitational physics and cosmology\\footnote{For various perspectives on this issue see the other contributions in this volume} as a result of the observations of a possible variation of the fine structure constant $\\alpha_{em}$ at high redshifts ($z> 0.5$) by Webb et al~\\cite{Webb:2000mn}. These observations remind us that our physics is based on peculiar coupling constants that might also be evolutionary on the cosmological scale. Variations of fundamental constants are a common feature in the generalizations of Einstein's theory of general relativity (GR)~\\cite{Will:2005va}. Extensions of GR have not only been claimed to be unavoidable when approaching the Planck scale of energies, since gravitation is expected to be unified with all the other fundamental interactions, but they have also been advocated as an explanation for the late time acceleration of the universe recently unveiled by cosmological observations~\\cite{Lobo:2008sg,Bertolami:2007gv,Nunes:2009dj,Odintsov 2010}. Scalar-tensor (ST) gravity theories, in particular, provide an appropriate theoretical framework for the variation of Newton's gravitational constant, which is induced by the dynamics of a scalar-field non-minimally coupled to the space-time geometry. The experimental scrutiny of scalar-tensor gravity theories requires a detailed analysis of their post-newtonian features, and is encapsulated into the so-called parametrised post-newtonian formalism (PPN) ~\\cite{Will:2005va,Damour:2002vu,Martins:2009zz,Bohmer:2009yx}. This procedure assumes two hypothesis: On the one hand, that there should be a weak field limit of the GR solution; On the other hand that the latter corresponds to the limit case of a given ST solution. In the present work we investigate the impact of a hyperbolic geometry on the possible variation of Newton's constant $G$. This question has been somewhat overlooked in the past, and, as we will show in the present work, raises a fundamental question regarding the physical interpretation of the results. To address this issue we derive a new scalar-tensor solution with an hyperbolic threading of the spatial hypersurfaces\\cite{Lobo:2009du}. Our solution extends a general relativistic solution known as being of a degenerate class A2 for its unusual properties\\cite{Stephani:2003tm,Ehlers & Kundt 1962}. The latter GR solution is characterised by a threading of the spatial hypersurfaces by means of pseudo-spheres instead of spheres. It does not exhibit the usual weak field limit in the region where the gravitational field is static, because the gravitational field has a repulsive character. This absence of a weak field limit for the hyperbolic GR solution means that such limit is also absent for comparison with the ST solution, and thus one cannot barely apply the PPN formalism. To address the latter question, we believe that one should look at the perturbations of the general relativistic limit rather than of the absent newtonian weak field. At least this enables us to assess the effects of the variation of $G$. ", "conclusions": "We have considered a static solution with a pseudo-spherical foliation of space. We reviewed its exotic features, and derived the extended scalar-tensor solution. The fundamental feature of these solutions is the absence of a newtonian weak field limit. Indeed it is known that not all of the GR solutions allow a newtonian limit, and this is the situation here. However, assuming that the solutions of the Einstein field equations represent gravitational fields, albeit far from our common physical settings, it is possible to ascertain the implications of varying $G$ in the strong fields by comparing the ST to their GR counterparts. From the viewpoint of observations this relies on the future detection of gravitational waves. We conclude with a quotation from John Barrow ~\\cite{Barrow:1991fj} which seems appropriate here \\begin{quote} {\\em The miracle of general relativity is that a purely mathematical assembly of second-rank tensors should have anything to do with Newtonian gravity in any limit}. \\end{quote}" }, "1101/1101.3363_arXiv.txt": { "abstract": "We report on a new XMM-Newton observation of NGC 247 from December 2009. The galaxy contains a supersoft, ultraluminous X-ray source (ULX) whose spectrum consists of a thermal component with a temperature about 0.1~keV and a power-law tail with a photon index around 2.5. The thermal emission is absolutely the dominant component, contributing 96\\% of the total luminosity in the 0.3-10 keV band. Variability is detected at timescales of $10^2$~s and longer with a $\\nu^{-1}$ power spectrum. These properties are consistent with black hole binaries in the thermal state and suggest the presence of an intermediate mass black hole of at least 600 solar masses. However, the integrated rms power is much higher than typically found in the thermal state. An alternative explanation of the emission could be a photosphere with a radius about $10^9$~cm. A possible absorption feature around 1~keV is detected, which may be due to absorption of highly ionized winds. X-ray sources within the disk of NGC 247 have a luminosity function consistent with that found in low mass X-ray binaries. We confirm previous results that X-rays from the quasar PHL~6625 may be absorbed by gas in NGC 247, mainly at energies below 0.3~keV. ", "introduction": "Nonnuclear X-ray sources with luminosities higher than $3 \\times 10^{39}$~\\ergs\\ are classified as ultraluminous X-ray sources (ULXs). Many of them are variable and most likely powered by accretion onto black holes. However, their luminosities, assuming isotropic emission, are higher than the maximum seen in Galactic black hole binaries, suggesting the presence of intermediate mass black holes with masses of $10^2 - 10^4$~$M_\\sun$ \\citep{col99,kaa01,far09} or super-Eddington accretion onto stellar mass black holes of about 10~$M_\\sun$ \\citep{wat01,beg02}. Discovery of intermediate mass black holes would be important because they cannot be formed via core collapse of a single star with normal metallicity \\citep{bel10} and may shed light onto the formation of supermassive black holes in the early Universe \\citep{ebi01,vol10}. Dynamical measurement of the compact object mass for ULXs has been unfeasible to date \\citep{rob10}. Instead, indirect means such as X-ray spectroscopy may shed light on the black hole mass. At a fixed Eddington ratio, the inner temperature of the accretion disk scales with black hole mass as $T_{\\rm in} \\propto M^{-1/4}$ \\citep{mak00}. Application of this relation requires the black hole binary to be in the thermal state \\citep{mcc06,rem06}, in which the disk is believed to extend all the way to the last stable orbit around the black hole. However, the thermal state is rarely found in ULXs, the only known occurrences are M82 X41.4+60 \\citep{fen10} and M82 X37.8+54 \\citep{jin10} during outburst. X-ray sources with soft, $kT_{\\rm in} \\sim 0.1$~keV, thermal spectra may be good candidates for intermediate, $\\sim 10^3 \\, M_\\sun$, mass black holes in the thermal state. NGC 247 is a nearby spiral galaxy in the Sculptor Group, with a modest star formation rate estimated to be about 0.1~$M_\\sun$~yr$^{-1}$ by counting bright main-sequence stars \\citep{dav06} or based on H$\\alpha$ observations \\citep{fer96}. However, infrared and radio observations suggest a star formation rate 10 or 100 times lower \\citep{dav06}. The distance to NGC 247 has been measured independently by different groups using various means. Among them, the recent result from \\citet{gie09} using near-infrared observations of Cepheids is claimed to be the most accurate. We therefore adopt that distance of 3.4~Mpc in this paper. Previously, NGC 247 has been observed in the X-ray band with Einstein \\citep{fab92}, ROSAT \\citep{zan97,rea97,lir00}, and XMM-Newton \\citep{win06}. A bright, off-nucleus X-ray source (NGC 247 ULX = 1RXS~J004704.8$-$204743) was detected in all of these observations. ROSAT observations revealed a very soft spectrum with $kT = 0.12$~keV \\citep{rea97}. Later on, the source was observed to have brightened by a factor of 2 \\citep{lir00}, suggestive of an accretion powered system. In 2001, XMM-Newton detected the source at the highest flux observed to date, with an unabsorbed luminosity of $8 \\times 10^{39}$~\\ergs\\ (at 3.4~Mpc) in the 0.3-10 keV band, and the spectrum was still dominated by a soft thermal component with $kT = 0.12$~keV \\citep{win06}. These results imply that the source is a good candidate to be an intermediate mass black hole in the thermal state. Unfortunately, the 2001 XMM-Newton observation suffered strong background flares, resulting in little usable data and poor statistics in the extracted spectrum. Thus, we proposed a new XMM-Newton observation of NGC 247 to obtain a better spectral measurement. In this paper, we report results from the new observation about the ULX as well as other sources in NGC 247. ", "conclusions": "\\subsection{X-ray population in NGC 247} The XMM-Newton observation of NGC 247 obtained in December 2009 is, to date, the deepest X-ray observation of the galaxy and allows an X-ray source population study. As shown in Figure~\\ref{fig:lum}, sources inside the D25 ellipse of NGC 247 have a luminosity function well consistent with that of LMXBs. A population of X-ray sources dominated by LMXBs is reasonable given the low star formation rate of the galaxy. For sources outside of the D25 ellipse of the galaxy, the $\\log N (>S)$ versus $\\log S$ distribution is roughly consistent with that expected for AGN \\citep{bra05} (Figure~\\ref{fig:logns}) at fluxes above $2 \\times 10^{-14}$~\\ergcms. Below this flux, the distribution flattens significantly with decreasing flux. This may be caused by the decrease of sensitivity due to vignetting at large off-axis angles. Some of these sources have bright optical counterparts with an X-ray to optical flux ratio consistent with that of AGN. Therefore, we conclude that most sources outside of the D25 ellipse of NGC 247 are likely background objects. \\subsection{NGC 247 ULX} The observed X-ray spectrum of the ULX is completely dominated by a cool, thermal emission component with a temperature around 0.1~keV; a weak power-law tail with a photon index of about 2.5 is also evident but contributes only a small fraction of the flux. This spectral shape is in contrast with that of many other ULXs whose high energy component contributes the majority of the flux (more than 60\\%) and can be explained by an optically-thick corona \\citep{gla09}. In fact, Gladstone et al. (2009) fits can predict an intrinsically soft-dominated spectrum, but this is a model-dependent result, while the observed spectral shapes of their sources are always dominated by the emission from a hard tail. It is unlikely that the high energy component in NGC 247 ULX arises from an optically thick corona; otherwise, such a corona must have a tiny covering factor since most disk photons are unscattered. The spectrum, however, is consistent with that of black hole binaries in the thermal state \\citep{mcc06,rem06}, and may hint at a connection in nature. An important observational characteristic of the thermal state is that the disk luminosity varies roughly with the 4th power of the disk inner temperature, equivalent to a constant disk inner radius. This source has shown a consistent disk inner radius derived from two XMM-Newton observations, though the constraint from the 2001 observation is loose. For a positive identification of the thermal state, one also needs timing information. This ULX displays red noise with a $\\nu^{-1}$ form above the white noise, which is consistent with that seen in the thermal state. However, the integrated rms power is significantly higher than that seen in the thermal state. For Galactic black holes, the total rms power in the 0.1-10~Hz band is usually less than 0.075 in the thermal state, and only in the hard state is the rms power found to be greater than 0.15 \\citep{rem06}. However, the X-ray spectrum of the ULX is much too soft to be classified as in the hard state. If the ULX is in the thermal state, then the source would be an interesting candidate to be an intermediate mass black hole. The accretion disk in the thermal state is believed to extend all the way to the innermost stable circular orbit (ISCO) around the central black hole. The ISCO radius depends only on the black hole mass and spin, as $R_{\\rm ISCO} = 6GM/c^2$ for a non-spinning black hole or 6 times smaller for an extremely spinning object, where $G$ is the gravitational constant, $M$ is the black hole mass, and $c$ is the light speed. Assuming a hardening correction of 1.7 and that the maximum temperature occurs at 2.4 times the ISCO radius \\citep{mak00}, the ISCO radius can be expressed as $R_{\\rm ISCO} \\approx 1.2 R_{\\rm in}$, where $R_{\\rm in}$ is the apparent inner radius derived directly from the normalization of the MCD component. Adopting the parameters from the best-fitted model, MCD plus power-law with absorption edge, and assuming a face-on disk, we have $R_{\\rm in} = (0.5-2.2) \\times 10^4$~km in the 90\\% confidence region, corresponding to the ISCO radius of a $600-3000 M_\\sun$ Schwarzschild black hole, or a more massive black hole if spinning or the disk is tilted. The observed luminosity is about a few percent of the Eddington limit of such a black hole, which is consistent with Eddington ratios typically found in the thermal state. Similar results can be obtained with the Gaussian absorption line. Due to the lack of a characteristic feature in the power spectrum, we are unable to constrain the black hole mass via timing. The rms-mass relation derived for AGN \\citep{zho10} may have issues in extrapolation to lower masses. The faintness in the 2-10 keV band of the source prevents us from implementing the means proposed by \\citet{gon11} to estimate the black hole mass. A few other ULXs show a dominant soft, thermal emission component, and are classified as supersoft ULXs: M101 ULX-1 \\citep{pen01,kon05,muk05}, Antennae X-13 \\citep{fab03}, M81 ULS1 \\citep{swa02,liu08}, and NGC 4631 X1 \\citep{car07,sor09}. These sources also show extreme variability: the observed fluxes vary by up to a factor of a few tens, and the intrinsic luminosity changes may be as large as factors of $10^2-10^3$ \\citep{kon04,muk05,liu08}. Short-term variability has been detected in almost all of these sources at timescales down to a few ks \\citep{swa02,muk05,car07,liu08}. M101 ULX-1 also showed a power spectrum with a $\\nu^{-1}$ form at frequencies below $10^{-3}$~Hz \\citep{muk05} during its outburst in 2000 March with variability more pronounced at high energies (0.8-2 keV) than in the lower band (0.2-0.8 keV) \\citep{muk03}. For a direct comparison, we calculated the rms power of NGC 247 ULX in the same energy bands and found that the fractional rms is 0.80 in the 0.8-2 keV band and 0.44 in the 0.2-0.8 keV band. Similar to NGC 247 ULX, NGC 4631 X1 also exhibits an absorption edge at 1.01~keV with an optical depth of 2.0 in its high state \\citep{sor09}. These similarities suggest that NGC 247 ULX may be a new member of the supersoft, ultraluminous family. The other soft ULXs seem to keep a constant temperature during huge changes in luminosity, indicative of a dramatic change of the emitting surface area of the same order. \\citet{fab03} has pointed out that this would lead to difficulty in interpreting the emission as due to accretion power release within a few Schwarzschild radii of the black hole. For NGC 247 ULX, the four ROSAT and two XMM-Newton observations have detected a change by a factor of about 3 of the observed flux in 0.1-2 keV, which is significantly smaller than others. Therefore, the thermal state interpretation has not been ruled out for NGC 247 ULX; new, multiple observations are the key to revealing its nature. An alternative explanation of the soft, thermal emission is that it arises from a photosphere that has a radius of $10^9$~cm, inferred from the blackbody component. The shortest timescale of variation above white noise is about $10^2$~s (Figure~\\ref{fig:pow}), which places a constraint that the emitting area must be smaller than $10^{12}$~cm, compatible with the size obtained from the blackbody parameters. This photosphere is unlikely the surface of a white dwarf, as the observed luminosity and temperature are too high to be explained by steady nuclear burning on a white dwarf surface \\citep{sta04}. Plus, presence of the power-law component with a 0.3-10 keV luminosity of $1.2 \\times 10^{38}$~\\ergs\\ is unexpected in such a case. It has been suggested that supercritical accretion could drive massive outflows above the disk and form an optically-thick photosphere, with a temperature expected to be about 0.1 keV, which could shield the inner disk emission when viewed at a high inclination angle \\citep{ohs05,pou07}. In this picture, the supersoft ULXs and normal ULXs are the same objects with different viewing angles. This interpretation may be in accord with the absorption feature observed at $\\sim$1~keV, which is likely due to L shell transitions of highly ionized iron. Besides this ULX and NGC 4631 X1, less similar absorption features near this energy have been detected in the ULX NGC 1365 X1 \\citep{sor07} and in the quasar PDS 456 \\citep{ree03}, and are explained as a spectral signature of a massive outflow. The absence of absorption from low-Z elements like oxygen could be a consequence of high ionization. We conclude that the X-ray properties of the NGC 247 ULX are unlike any popular emission state known in black hole binaries. The source exhibits a spectrum like the thermal state, but shows variability much stronger than seen in the thermal state. A possible strong absorption feature is detected. This source and other supersoft ULXs may exhibit a new accretion state that is different from those seen in Galactic black hole binaries and the majority of ULXs. A variety of potential thermal emission components have been suggested to be present in the energy spectrum of ULXs, including the soft excess (0.1-0.4~keV) found in the majority of ULXs \\citep{fen05,sto06}, the high-temperature ($\\sim$keV) ULXs \\citep{mak00}, and supersoft ULXs. The properties of these thermal components are summarized and compared in Table~\\ref{tab:comparsion}, with NGC 247 ULX and M82 X41.4+60 in the thermal state highlighted. Among them, a disk interpretation is best established for M82 X41.4+60 because its thermal state is consistent with those of stellar-mass X-ray binaries in terms of timing and spectral properties and spectral evolution \\citep{fen10}. For the others, definitive evidence is lacking for positive determination of their nature, except that the disk explanation is ruled out for other supersoft ULXs other than NGC 247 ULX. Based on the timing information, it appears unlikely that the soft excess in 'normal' ULXs and the soft (dominant) spectral component in supersoft ULXs have the same physical origin. \\subsection{PHL 6625} PHL~6625 was identified as a radio-quiet background QSO at $z = 0.38$ based on a Mg~{\\sc ii} emission line \\citep{mar85}. A single power-law model with interstellar absorption provides an adequate fit to its X-ray spectrum. The parameters listed in Table~\\ref{tab:top} are obtained from fitting in the 0.3-10~keV energy range, and the hydrogen column density is very close to the Galactic value at $N_{\\rm H} = 2 \\times 10^{20}$~cm$^{-2}$, indicative of negligible absorption within NGC 247. However, if we include data down to 0.2~keV, two more energy bins, then the best-fit absorption is $N_{\\rm H} = 5.3_{-1.0}^{+1.2} \\times 10^{20}$~cm$^{-2}$. This is caused by the low flux between 0.2-0.3~keV, and is consistent with the results obtained from ROSAT observations, in which a total absorption column density $N_{\\rm H} = 5.4_{-2.1}^{+2.4} \\times 10^{20}$~cm$^{-2}$ was derived from fitting in the 0.1-2~keV band, and the excess absorption mainly takes effect at energies below 0.3~keV \\citep{elv97}. Those authors claimed a positive detection of absorbing materials in NGC 247 assuming a Galactic $N_{\\rm H} = 1.4 \\times 10^{20}$~cm$^{-2}$ based on 21~cm observations of neutral hydrogen. Here, we confirm this conclusion with tighter constraints, even assuming a higher Galactic absorption. Therefore, with future high-resolution spectroscopic observations, this QSO could be used to study the interstellar medium within NGC 247. For the 2009 data, the power-law photon index at the rest frame is $\\Gamma = 2.10_{-0.09}^{+0.13}$. We extracted spectra for this object from the 2001 XMM-Newton observation and found a harder spectrum with $\\Gamma = 1.64 \\pm 0.17$ at a flux level of $3.7 \\pm 0.6 \\times 10^{-13}$~\\ergcms\\ in 0.3-10~keV similar to that in 2009. The fluxes observed with XMM-Newton are about ten times lower than the ROSAT flux of $(5.9 \\pm 0.7) \\times 10^{-12}$~\\ergcms\\ in the 0.2-2.4 keV band \\citep{elv97}. No short-term X-ray variability has been detected for this object. \\subsection{Other individual sources} {\\bf Source 28}. --- The source is hard, and its spectrum can be described by a power-law with $\\Gamma = 1.27$, but the residuals vary systematically with energy. Adding an MCD component improves the residuals and results in $\\chi^2 / {\\rm dof} = 31.7/54$. The disk has a temperature at the inner radius of $0.35_{-0.10}^{+0.14}$~keV with a bolometric luminosity of $2 \\times 10^{37}$~\\ergs\\ assuming a face-on geometry, and the power-law component has a photon index $\\Gamma = 0.6_{-0.3}^{+0.2}$. The disk component has a fractional luminosity of 11\\% in the 0.3-10 keV band. If this source is a black hole binary, it is possibly in the hard state. {\\bf Source 34}. --- The source sits at the north end of the NGC 247 disk, and was bright during the 2001 XMM-Newton observation with $\\Gamma = 2.3_{-0.9}^{+1.5}$, $N_{\\rm H} = 1.9_{-0.8}^{+2.4} \\times 10^{21}$~cm$^{-2}$, and $f_{\\rm X} = 2.0 \\times 10^{-13}$~\\ergcms\\ in the 0.3-10 keV band, or $L_{\\rm X} = 4.6 \\times 10^{38}$~\\ergs\\ corrected for absorption. However, it became faint in the 2009 XMM-Newton observation, with a flux of $3.4 \\times 10^{-14}$~\\ergcms\\ in the 0.3-10 keV band, six times dimmer than in 2001. It is possibly source 3 in \\citet{fab92} from Einstein observations and source 1 in \\citet{rea97} from ROSAT observations. The ROSAT observations revealed a luminosity of $1.7 \\times 10^{38}$~\\ergs\\ corrected for Galactic absorption only (at 3.4~Mpc) in the 0.1-2 keV band, or $3 \\times 10^{38}$~\\ergs\\ corrected for all absorption. The luminosity in the same band from the 2001 XMM-Newton observation is about $6 \\times 10^{38}$~\\ergs. There is no bright optical counterpart found at the X-ray position. However, the source is associated with a possible star forming region, seen in the ultraviolet image from the 2001 XMM-Newton observation, with diffuse emission and star clusters nearby. Therefore, the source is likely an X-ray binary in NGC 247, and is perhaps an accreting black hole. {\\bf Source 36}. --- This source could be source 3 in \\citet{rea97}. The best-fitted column density is consistent with the Galactic value. Its flux estimated from the 2001 XMM-Newton observation is consistent with that in 2009. In the X-ray image, the source seems to show extended features and may consists of multiple objects, which can only be resolved with future Chandra observations. {\\bf Source 61}. --- Its X-ray spectrum cannot be adequately fitted by a single power-law model, with $\\chi^2 / {\\rm dof} = 80.9/60$; adding a Gaussian absorption line significantly improves the fit with $\\chi^2 / {\\rm dof} = 58.1/57$. The absorption line has a central energy of $0.9 \\pm 0.1$~keV and a width of $\\sigma = 0.22_{-0.16}^{+0.28}$~keV with an optical depth of $0.29_{-0.19}^{+0.58}$. The equivalent width is calculated to be around 0.30~keV. Its observed flux in 2001 was about twice of that in 2009. There is a bright optical source with $B = 18.83$ and $R = 18.58$ within 2\\arcsec\\ of the X-ray position, and the X-ray to optical flux ratio is estimated to be $-0.3$. As discussed above, this source is likely a background AGN. Due to the low absorption column density, the absorption line is likely produced intrinsic to the source. Similar absorption has been seen in AGN \\citep{blu02}. {\\bf Source 13}. --- This is the hardest source detected with XMM-Newton. It is obvious in the hard band (2-10~keV) but undetected in the soft band (0.3-2~keV), consistent with $N_{\\rm H} > 5 \\times 10^{22}$~cm$^{-2}$ for a power-law spectrum with $\\Gamma = 2$. The source lies about 7\\arcmin\\ west of the nucleus of NGC 247, and is likely a highly absorbed background AGN. {\\bf Sources 41, 50, 59, and 72}. --- These four sources are the softest besides the ULX and are only detected in the soft band (0.3-2 keV) with a hardness ratio around $-0.9$, which is consistent with a spectrum with a blackbody temperature $< 0.4$~keV or a power-law photon index $\\Gamma > 2.5$, assuming Galactic absorption. Sources 41 and 72 are outside of the D25 ellipse of NGC 247, but sources 50 and 59 are inside. Their X-ray luminosities are about $10^{36} - 10^{37}$~\\ergs. These sources do not appear to have bright optical counterparts. They could be canonical supersoft sources that are powered by nuclear burning on the surface of white dwarfs, or high redshift AGN. {\\it Facility:} \\facility{XMM-Newton}" }, "1101/1101.4891_arXiv.txt": { "abstract": "The theory for the formation of the first population of stars (Pop III) predicts a IMF composed predominantly of high-mass stars, in contrast to the present-day IMF, which tends to yield stars with masses less than 1 M$_{\\odot}$. The leading theory for the transition in the characteristic stellar mass predicts that the cause is the extra cooling provided by increasing metallicity and in particular the cooling provided at high densities by dust. The aim of this work is to test whether dust cooling can lead to fragmentation and be responsible for this transition. To investigate this, we make use of high-resolution hydrodynamic simulations. We follow the thermodynamic evolution of the gas by solving the full thermal energy equation, and also track the evolution of the dust temperature and the chemical evolution of the gas. We model clouds with different metallicities, and determine the properties of the cloud at the point at which it undergoes gravitational fragmentation. We follow the further collapse to scales of an AU when we replace very dense, gravitationally bound, and collapsing regions by a simple and nongaseous object, a sink particle. Our results suggest that for metallicities as small as 10$^{-5} \\rm Z_{\\odot}$, dust cooling produces low-mass fragments and hence can potentially enable the formation of low mass stars. We conclude that dust cooling affects the fragmentation of low-metallicity gas clouds and plays an important role in shaping the stellar IMF even at these very low metallicities. ", "introduction": "\\label{int} The first burst of star formation in the Universe is thought to give rise to massive stars, the so-called Population III, with current theory predicting masses in the range 20-150 M$_{\\odot}$ \\citep{2002Sci...295...93A, 2002ApJ...564...23B, 2007ApJ...654...66O, 2008Sci...321..669Y}. This contrasts with present-day star formation, which tends to yield stars with masses less than 1 M$_{\\odot}$ \\citep{2002Sci...295...82K, 2003PASP..115..763C}, and so at some point in the evolution of the Universe there must have been a transition from primordial (Pop. III) star formation to the mode of star formation we see today (Pop. II/I).\\\\ When gas collapses to form stars, gravitational energy is transformed to thermal energy and unless this can be dissipated in some fashion, it will eventually halt the collapse. Thermal energy can be dissipated by processes such as atomic fine structure line emission, molecular rotational or vibrational line emission, or the heating of dust grains. In some cases, these processes are able to cool the gas significantly during the collapse. This temperature drop can promote gravitational fragmentation \\citep{2004RvMP...76..125M, 2007prpl.conf..149B} by diminishing the Jeans mass, which means that instead of forming very massive clumps, with fragment masses corresponding to the initial Jeans mass in the cloud, it can instead form even more fragments with lower masses.\\\\ If the gas is cooled only by molecular hydrogen emission, numerical simulations show that the stars should be very massive \\citep{2002Sci...295...93A, 2002ApJ...564...23B, 2007ApJ...654...66O, 2008Sci...321..669Y}. This happens because the H$_2$ cooling becomes inefficient for temperatures bellow 200K and densities above $10^4 \\rm cm^{-3}$. At this temperature and density, the mean Jeans mass at cloud fragmentation is 1,000 times larger than in present-day molecular clouds, \\begin{equation}\\label{BEm} M_{\\rm frag} \\approx 700M_{\\odot} \\left(\\frac{T_{\\rm frag}}{200 \\rm K}\\right)^{3/2} \\left(\\frac{n_{\\rm frag}}{10^4 \\rm cm^{-3}}\\right)^{-1/2}, \\end{equation} for an atomic gas with temperature $T_{\\rm frag}$ and number density $n_{\\rm frag}$.\\\\ Metal line cooling and dust cooling are effective at lower temperatures and larger densities, and so the most widely accepted cause for the transition from Pop. III to Pop. II is metal enrichment of the interstellar medium by the previous generations of stars. This suggests that there might be a critical metallicity Z$_{\\rm crit}$ at which the mode of star formation changes.\\\\ The main coolants that have been studied in the literature are CII and OI fine structure emission \\citep{2001MNRAS.328..969B, 2003Natur.425..812B, 2006ApJ...643...26S, 2007MNRAS.380L..40F}, and dust emission. C and O are identified as the key species because in the temperature and density conditions that characterise the early phases of Pop. III star formation, the OI and CII fine-structure lines dominate over all other metal transitions \\citep{1989ApJ...342..306H}. By equating the CII or OI fine structure cooling rate to the compressional heating rate due to free-fall collapse, one can define critical abundances $[\\rm C/H] = -3.5$ and $[ \\rm O/H] = -3.0$\\footnote{$[\\rm X/ \\rm Y] = log_{10}(N_{\\rm X}/N_{\\rm Y})_{\\star} - log_{10}(N_{\\rm X}/N_{\\rm Y})_{\\odot}$, for elements X and Y, where $\\star$ denotes the gas in question, and where $\\rm N_X$ and $\\rm N_Y$ are the mass fractions of the elements X and Y.} for efficient metal line cooling \\citep{2003Natur.425..812B}. However, previous works \\citep{2009ApJ...696.1065J, 2009ApJ...694.1161J} show that this metallicity threshold does not represent a critical metallicity: the fact that metal-line cooling has a larger value than the compressional heating does not necessarily lead to fragmentation.\\\\ Dust-cooling models predict a much lower critical metallicity $(\\rm Z_{\\rm crit} \\approx 10^{-5} \\rm Z_{\\odot})$. The conditions for fragmentation in the low-metallicity dust cooling model are predicted to occur in high density gas, where the distances between the fragments can be very small \\citep{2000ApJ...534..809O, 2005ApJ...626..627O, 2002ApJ...571...30S, 2006MNRAS.369.1437S, 2010MNRAS.402..429S}. In this regime, interactions between fragments will be common, and analytic models of fragmentation are unable to predict the mass distribution of the fragments. A full 3D treatment, following the fragments, is needed.\\\\ Initial attempts were made by \\cite{2006ApJ...642L..61T,2008ApJ...676L..45T} and \\cite{2008ApJ...672..757C}. However, these treatments used a tabulated equation of state, based on results from previous one-zone chemical models \\citep{2005ApJ...626..627O}, to determine the thermal energy. This approximation assumes that the gas temperature adjusts instantaneously to a new equilibrium temperature whenever the density changes and hence ignores thermal inertia effects. This may yield too much fragmentation.\\\\ In this work, we improve upon these previous treatments by solving the full thermal energy equation, and calculating the dust temperature through the energy equilibrium equation. We assume currently that the only significant external heat source is the CMB, and include its effects in the calculation of the dust temperature.\\\\ ", "conclusions": "\\label{conc} In this paper we have addressed the question of whether dust cooling can lead to the fragmentation of low-metallicity star-forming clouds. For this purpose we performed numerical simulations % to follow the thermodynamical and chemical evolution of collapsing clouds. The chemical model included a primordial chemical network together with a description of dust evolution, where the dust temperature was calculated by solving self-consistently the thermal energy equilibrium equation% .\\\\ We performed three sets of simulations, two at low resolution and one at high resolution (Table \\ref{tsim}). All simulations had an initial cloud mass of 1000 M$_{\\odot}$, number density of 10$^5$ cm$^{-3}$, and temperature of 300K. We tested two different metallicities (10$^{-4} {\\rm Z}_{\\odot}$ and 10$^{-5} {\\rm Z}_{\\odot}$), and also the inclusion of small amounts of turbulent and rotational energies.\\\\ We found in all simulations that dust can effectively cool the gas, for number densities higher than $10^{11} \\rm cm^{-3}$. An increase in metallicity implies a higher dust-to-gas ratio, and consequently stronger cooling by dust. This is reflected in a lower temperature of the dense gas in the higher metallicity simulation.\\\\ For the low resolution case, we tested the effect of adding turbulence and rotation. These diminish the infall velocity, leading to different fluid elements undergoing different amounts of compressional heating. This lack of heating allows the gas to reach a lower temperature.\\\\ We found that the transport of angular momentum to smaller scales lead to the formation of a disk-like structure, which then fragmented into a number of low mass objects.\\\\ We conclude that the dust is already an efficient coolant even at metallicities as low as 10$^{-5}$ or 10$^{-4}Z_{\\odot}$, in agreement with previous works \\citep{2008ApJ...672..757C, 2010ApJ...722.1793O, 2002ApJ...571...30S, 2006MNRAS.369.1437S, 2006ApJ...642L..61T, 2008ApJ...676L..45T}. Our results support the idea that dust cooling can play an important role in the fragmentation of molecular clouds and the evolution of the stellar IMF.\\\\" }, "1101/1101.1961_arXiv.txt": { "abstract": "\\change{} We construct catalogues of present superstructures that, according to a $\\Lambda CDM$ scenario, will evolve into isolated, virialized structures in the future. We use a smoothed luminosity density map derived from galaxies in SDSS-DR7 data and separate high luminosity density peaks. The luminosity density map is obtained from a volume--limited sample of galaxies in the spectroscopic galaxy catalogue, within the SDSS-DR7 footprint area and in the redshift range $0.040$, the broad lines zero-lag, lag or lead the $\\gamma$-rays, respectively. It is applied to 3C 273, in which the lines and the radio emission have enough data, but the $\\gamma$-rays have not. We find $\\tau_{\\rm{ob}}<0$ and $\\tau_{\\rm{ob}}>0$ for the 5, 8, 15, 22 and 37 GHz emission relative to the broad lines H$\\alpha$, H$\\beta$ and H$\\gamma$. The lag may be positive or negative, however current data do not allow to discriminate between the two cases. The measured lags are on the order of years. For a given line, $\\tau_{\\rm{ob}}$ generally decreases as radio frequency increases. This trend most likely results from the radiative cooling of relativistic electrons. The negative lags have an average of $\\tau_{\\rm{ob}}=-2.86$ years for the 37 GHz emission, which represents that the lines lag the radio emission. The positive lags have $\\tau_{\\rm{ob}}=3.20$ years, which represents that the lines lead the radio emission. We obtain the radio emitting positions $R_{\\rm{radio}}=0.40$--2.62 pc and $R_{\\rm{radio}}=9.43$--62.31 pc for the negative and positive lags, respectively. From the constraint of $R_{\\rm{\\gamma}}\\la R_{\\rm{radio}}$ \\citep[e.g.][]{b24,b41}, we have $R_{\\rm{\\gamma}}\\la 0.40$--2.62 pc for the negative lags. For the positive lags, 4.67--$30.81 R_{\\rm{BLR}}$, i.e. $R_{\\rm{BLR}}R_{\\rm{BLR}}$ for BL Lacs. They used $R_{\\rm{BLR}}=10^{17}L^{1/2}_{\\rm{d,45}}$ cm to estimate $R_{\\rm{BLR}}$ for BL Lacs and FSRQs, where $L_{\\rm{d,45}}$ is accretion disc luminosity in units of $10^{45} \\/\\ \\rm{erg \\/\\ s^{-1}}$. It is appropriate to use $R_{\\rm{BLR}}=10^{17}L^{1/2}_{\\rm{d,45}}$ cm to estimate $R_{\\rm{BLR}}$ for FSRQs, but not for BL Lacs because this relation is derived from the type 1 AGNs. Thus it should be reliable that $R_{\\rm{diss}} R_{\\rm{BLR}}$ \\citep{b54,b79}. This is marginally consistent with $R_{\\rm{BLR}} R_{\\rm{BLR}}$ of \\citet{b54} and \\citet{b79}. These confirm the reliability of our results. Our previous works are applicable to the $\\gamma$-rays emitted from the regions in powerful blazars, where $R_{\\rm{\\gamma}}$ is not much larger than $R_{\\rm{BLR}}$ \\citep{b55,b56,b10}. The method proposed here can locate $R_{\\rm{\\gamma}}$ in the jet. The inner emitting regions with $R_{\\rm{\\gamma}}\\la R_{\\rm{BLR}}$ are likely the major contributor of the $\\gamma$-rays below 10 GeV, for that the $\\gamma$-rays above 10 GeV are subject to the photon-photon absorption due to the dense external soft photons at the inner regions \\citep[e.g.][]{b55,b56,b10}. The outer emitting regions with $R_{\\rm{\\gamma}}\\gg R_{\\rm{BLR}}$ are likely the major contributor of the $\\gamma$-rays above 10 GeV, for that these $\\gamma$-rays are not subject to the photon-photon absorption due to the thin external soft photons at the outer regions. For these possible $\\gamma$-ray emitters at large scales of kpc--Mpc \\citep{b8,b97,b96}, our works are not applicable. The most prominent features on VLBI images of jets in radio-loud AGNs are the radio core and bright knots in the jet \\citep{b42}. \\citet{b50} investigated the relation between AGN $\\gamma$-ray emission and pc-scale radio jets. They identified the pc-scale radio core as a likely location for both the $\\gamma$-ray and radio flares. A few hundreds of Schwarzschild radii, sub-pc-scale, is the preferred jet position where most of the dissipation occurs \\citep{b36,b38,b37}. \\citet{b77} suggested that the blazar emission zone is located at pc-scale distances from the nucleus. \\citet{b33} suggested that $R_{\\rm{\\gamma}}$ is at sub-pc scales. \\citet{b11} also suggested a sub-pc $\\gamma$-ray--emitting region from the central black hole. These previous findings support $R_{\\rm{\\gamma}}$ at sub-pc--pc scales. These $R_{\\rm{\\gamma}}$ of sub-pc--pc scales are consistent with those $R_{\\rm{\\gamma}}$ obtained in Cases B and D. These sub-pc--pc scale $R_{\\rm{radio}}$ obtained in Cases B and D are also consistent with the previous findings of the blazar emission zone and the dissipation zone. These agreements confirm the reliability of our results. It is recently advanced that the bulk of the $\\gamma$-rays is generated in regions of the jet at distances of tens of pc from the central black hole \\citep[e.g.][]{b76,b59}. For Case C, we obtain the outer emitting regions of 4.67--$30.810$ (Cases A, B and C), the broad lines zero-lag, lag and lead the $\\gamma$-rays, respectively. All cases are unified into equation (7). The method is applied to FSRQ 3C 273. Because the $\\gamma$-ray light curves are very sparsely sampled for 3C 273, it should be unreliable to employ them to estimate the time lags. Fortunately, it was suggested that $R_{\\rm{\\gamma}}\\la R_{\\rm{radio}}$ \\citep{b24,b41,b50,b5}. Thus $R_{\\rm{\\gamma}}$ could be constrained by $R_{\\rm{radio}}$ derived from the lags of the radio emission relative to the broad lines. The ZDCF method is used to analyze the correlations of the radio and infrared emission with the broad lines H$\\alpha$, H$\\beta$ and H$\\gamma$. The broad lines lag or lead the 5, 8, 15, 22 and 37 GHz emission (see Fig. 3). However, there is a lack of correlation between the infrared emission and the broad lines (see Fig. 4). The measured lags are on the order of years (see Table 1). The measured lags for a given line generally decrease as radio frequency increases (see Table 1). This trend most likely results from the radiative cooling of relativistic electrons (see Fig. 7). The measured negative lags have an average of $\\overline \\tau^{-}_{\\rm{cent}}=-2.86$ years for the 37 GHz emission relative to the broad lines. From $\\overline \\tau^{-}_{\\rm{ob}}=-2.86$ years, $R_{\\rm{BLR}}=2.70$ $\\rm{ly}$, $v_{\\rm{d}}=0.9$--$0.995c$, $\\theta=12^{\\rm{\\circ}}$--$21^{\\rm{\\circ}}$ and equation (7), we obtain $R_{\\rm{radio}}=0.40$--2.62 pc (Case B). These estimated $R_{\\rm{radio}}$ contain the typical size of $R_{\\rm{BLR}}=0.83$ pc, i.e. the radio emitting regions in Case B are around the BLR. Thus $R_{\\rm{\\gamma}}\\la 0.40$--2.62 pc for Case B and these inner emitting regions mostly satisfy $R_{\\rm{\\gamma}} \\la R_{\\rm{BLR}}$. For Case D, the special Case B, $R_{\\rm{\\gamma}}\\la R_{\\rm{radio}}=0.77$--0.81 pc and $R_{\\rm{\\gamma}} \\la R_{\\rm{BLR}}$. The measured positive lags have an average of $\\overline \\tau^{+}_{\\rm{cent}}=3.20$ years for the 37 GHz emission relative to the broad lines. From $\\overline \\tau^{+}_{\\rm{ob}}=3.20$ years, $R_{\\rm{BLR}}=2.70$ $\\rm{ly}$, $v_{\\rm{d}}=0.9$--$0.995c$, $\\theta=12^{\\rm{\\circ}}$--$21^{\\rm{\\circ}}$ and equation (7), we obtain $R_{\\rm{radio}}=9.43$--62.31 pc $\\gg R_{\\rm{BLR}}$ (Case C). Considering the zero-lag point and the constraint of $R_{\\rm{\\gamma}} \\la R_{\\rm{radio}}$, we have 4.67--$30.81 0$, $\\rho_{0} < 0$, ${\\bf E}_{\\parallel} = 0$ with electron acceleration; (ii) % ${\\bf \\Omega}\\cdot{\\bf B} < 0$, $\\rho_{0} > 0$, ${\\bf E}_{\\parallel} = 0$ with % positron and positive baryon acceleration or (iii) ${\\bf \\Omega}\\cdot{\\bf B} < % 0$, $\\rho_{0} > 0$ and ${\\bf E}_{\\parallel} \\neq 0$. Throughout this paper, for brevity, these conditions will be referred to as % cases (i), (ii) or (iii). There is no further possible case because free % emission of electrons is always possible at neutron-star surface temperatures. A previous paper (Jones 2010a, hereafter Paper I) reported some preliminary work % on the relation between these sets of boundary conditions and observed % phenomena, in particular, the existence of pulse nulls in a significant fraction % of radio pulsars. There seems to be no reason why neutron stars satisfying (i), % (ii) and even exceptionally (iii) should not exist and it is the purpose of the % present paper to continue the attempt to see if features of the predicted plasma % acceleration in these cases correlate with observed properties of subsets of the % pulsar population. There can be little doubt that in this context nulls and subpulse drift are important phenomena. Although there must be reservations about assigning undue % weight to the properties of a single neutron star, nulls in the isolated pulsar % PSR B1931+24 are informative. Kramer et al (2006) found that spin-down in the % on-state is approximately twice as fast as in the off-state. It is difficult to % see, in a pulsar of this age ($> 10^{6}$ yr), how the geometrical shape of the % acceleration volume or the magnitude of the current density ${\\bf J}$ within it % can change by a factor of this order in less than the rotation period $P$. % Time-varying fields near the light cylinder can further accelerate % ultra-relativistic particles of both signs and therefore the most obvious % explanation for the change in spin-down torque is that the particle and field % components of the magnetospheric momentum density and stress tensor near the light cylinder differ between on and off-states. A cessation of pair % creation during nulls would greatly change the charged particle number density % even though ${\\bf J}$ remains essentially unchanged. In this connection, it is % interesting that a recent paper by Lyne et al (2010) suggests quasi-periodic % switching between two different spin-down rates as the origin of the peculiar % timing residuals seen in many pulsars. The presence of nulls might be thought of as evidence that a pulsar is in the % final phase of evolution prior to complete cessation of radio emission. However, we suggest that this view is not consistent with a specific feature of % the 63 nulling pulsars listed in Tables 1 and 2 of the paper by Wang, Manchester % \\& Johnston (2007). The maximum potential difference $\\Phi _{max}$ available % for acceleration at the polar cap is proportional to $B_{d}P^{-2}$, in which % $B_{d}$ is the effective dipole field inferred from the spin-down rate. It is % approximately independent of the ratio of $B_{d}$ to the actual polar-cap field % $B$. The distribution of this quantity for all radio pulsars in the ATNF catalogue % (Manchester et al 2005) has a relatively sharp cut-off at $2.2\\times 10^{11}$ G % s$^{-2}$, equivalent to the existence of a well-defined death line in the % distribution of $B_{d}$ versus $P$. It might be expected that nulls should be seen only as a pulsar rotation slows % so that it approaches this value. But the form of the distribution of % $B_{d}P^{-2}$ for the nulling pulsars listed by Wang et al is broadly % indistinguishable from that of the whole ATNF catalogue and instead is % consistent with nulls being a long-term property of a certain sub-set of radio % pulsars. It is, of course just possible that the observed sharp cut-off is % itself a statistical fluctuation and that individual pulsar deaths actually % occur at values of $B_{d}P^{-2}$ throughout the whole distribution because some % other variable is involved, such as flux-line curvature. But we shall discount % this possibility and in the present paper adopt the view that pulsars with % boundary condition (ii) and, exceptionally, (iii) form the sub-set that null. A recent large-scale survey by Weltevrede, Edwards \\& Stappers (2006) has % revealed that subpulse drift is a common phenomenon. Under the assumption that % electron-positron pairs are the source of the coherent radio emission, it % implies that compact zones of pair creation exist and move in an organized way % within the open magnetosphere area of the polar cap. This motion has been % assumed to be an ${\\bf E}\\times{\\bf B}$ drift velocity under the case (iii) surface electric % field boundary condition, following the original paper of Ruderman \\& Sutherland % (1975). There have been many later papers on this subject and we refer to Gil, % Melikidze \\& Geppert (2003) for recent developments which have sought to refine % the calculation of the drift velocity. The problem is that the case (iii) % boundary condition is implausible as a common property of the pulsar population % and that cases (i) and (ii) have not hitherto provided any immediate and obvious % basis for the formation of organized subpulses. It was shown in Paper I that % short time-scale instability in the composition of accelerated plasma exists in % case (ii). In this paper, it is shown to be a plausible mechanism for subpulse % formation and drift. Under the assumption of an actual dipole field, pair creation is possible, in % principle, in all observed pulsars through the inverse Compton scattering (ICS) % of polar-cap photons by accelerated electrons or positrons (Hibschman \\& Arons % 2001, also see Fig. 1 of Harding \\& Muslimov 2002), but curvature radiation (CR) can % produce pairs only in those with high values of $B_{d}P^{-2}$. However, there is % a problem, noted by Harding \\& Muslimov, in the formation of a dense pair plasma % because the number of ICS pairs formed per primary electron or positron % accelerated can be smaller than unity (see also Fig. 8 of Hibschman \\& Arons). % They suggest that the pair density required for coherent radio emission may be % far smaller than previously thought. Although higher multipole field components % may be present in most pulsars, so increasing flux-line curvature, the existence % of this problem must remain a matter of concern perhaps even leading to doubts % about the role of pair plasma as the source of coherent radio emission. The existence of solutions of the electrostatic problem under boundary % conditions (ii) or (iii) is no more than a preliminary because we are concerned % in this paper with the presence of instabilities which arise from the reverse % electron flow at the polar cap. It might be questioned whether or not the charge % density on the surface separating open from closed magnetic flux-lines needed % for the condition ${\\Phi} = 0$ can be maintained in the presence of instability. % But we shall find that, even at short time-scales, instability principally % changes the nature of the particles accelerated rather than the current density % ${\\bf J}$ and the acceleration field. The instabilities considered here are % not, of course, a feature of case (i) in which electrons are the primary % component of ${\\bf J}$. Establishing instability in this case is a purely % electromagnetic problem which has been considered by Levinson et al (2005), % Melrose \\& Luo (2009) and Reville \\& Kirk (2010). Non-stationary flow in case % (iii) has been investigated by Timokhin (2010). The properties of the condensed matter surface are naturally important in cases % (ii) and (iii), in particular, the production of protons by the reverse flux of % electrons described previously (Paper I; see also Jones 1981). This is a % characteristic of electromagnetic showers that is usually of little practical % significance. The form of the shower depends principally on electron-photon % interactions but shower photons also interact directly with nuclei. The % cross-section is a maximum with the excitation of the giant dipole resonance % (GDR), a broad collective state, at a photon momentum $k\\approx 40$ mc. % (Electroproduction cross-sections are smaller by a factor of the order of the % fine-structure constant and can be neglected.) Photon track length per unit % interval of $k$ in a high-energy shower is $\\propto k^{-2}$ so that to a good % approximation, photoproduction of baryons can be assumed to occur entirely through GDR % formation. Known cross-sections and electromagnetic shower theory allow the % calculation of production rates, and we define $W_{p}$ as the number of protons produced per unit primary shower energy. Protons are initially of the order of % a few MeV but are very rapidly moderated to thermal energies without further % strong interaction, and then diffuse to the surface with a time delay that is of % prime significance for the stability of plasma acceleration. For further details we refer to Paper I, also to Jones (2010b) for the % cross-sections at high values of $B$ for processes of second order in % electron-photon coupling. In the early stages of acceleration, % partially-ionized atoms interact with the polar-cap blackbody radiation field % and this is the source of the reverse-electron flux that is considered here. It % has proved convenient to combine rates for this process with values of $W_{p}$ % by defining, for a particular pulsar, the parameter $K$ which is equal to the % number of protons produced per unit nuclear charge accelerated. Instabilities in solutions of the time-independent electrostatic problem % referred to above exist as a consequence of proton formation and we propose that % they are the basis for the nulling phenomenon and for subpulse formation and % drift. The complexity that arises is unfortunate because it limits what can be achieved % in terms of a physical theory of the acceleration process. Thus we shall be % able to expose the general properties attaching to cases (ii) and (iii) but are % not always able to give quantitative predictions. We shall show in Section 2 that general consideration of polar-cap parameters % rule out the possibility that the relatively numerous subsets of pulsars that % exhibit nulls and subpulse drift belong to case (iii). Therefore, sections 3 and % 4 of this paper are restricted to further consideration of case (ii) for which short time-scale instability in plasma acceleration has been % described previously in Paper I. Its properties are summarized in Section 3. The previous treatment of medium time-scales was, however, inadequate and % contained an error. The appropriate analysis is given here in Section 4. The % relations between these instabilities and the observed properties of nulls and % of subpulse drift are described in Section 5 and, in particular, the relation % beween the subpulse band separation periodicity $P_{3}$ and proton diffusion % time is given. \\section[]{Polar-cap parameters} Many of the polar-cap properties and parameters that will be required for the % arguments of Sections 3-5 have been given previously in Paper I. The properties % of the polar cap atmosphere are of particular importance and will be further % considered in Section 3. But it will be convenient to summarize the remainder % here with some additions. The most basic parameter is the ion number density $N$ of the condensed matter % at zero pressure. The magnetic dipole fields $B_{d}$ inferred from pulsar % spin-down rates vary over about five orders of magnitude, but the median value % for the 63 nulling pulsars listed by Wang et al is $2.8\\times 10^{12}$ G, which % is significantly larger than for the whole ATNF catalogue. It is also possible % that the actual polar-cap field $B$ is larger than $B_{d}$. For these reasons, we adopt the expression, \\begin{eqnarray} N = 2.6 \\times 10^{26}Z^{-0.7}B^{1.2}_{12}\\hspace{5mm} {\\rm cm}^{-3}, \\end{eqnarray} fitted to values computed by Medin \\& Lai (2006), primarily for $B_{12} > 10$, % where $B_{12}$ is the magnetic flux density in units of $10^{12}$ G and $Z$ is % the atomic number. The convenient unit of depth below the surface at $z = 0$ is % the radiation length, \\begin{eqnarray} l_{r} = 1.66Z^{-1.3}B^{-1.2}_{12}\\left(\\ln\\left(12Z^{1/2}B^{-1/2}_{12} \\right)\\right)^{-1}\\hspace{5mm}{\\rm cm} \\end{eqnarray} defined here in terms of the zero-field Bethe-Heitler bremsstrahlung % cross-section with screening factor modified for the neutron-star surface % density (see Paper I). The critical temperature above which the ion thermal emission rate is high % enough to maintain the case (ii) boundary condition is related to the cohesive % energy $E_{c}$ by $k_{B}T_{c} \\approx 0.025E_{c}$ (see the discussion of % atmospheric properties in Section 3). Cohesive energies have been calculated by % Medin \\& Lai as functions of $B$. For $Z = 26$ and $B_{12} = 10$, their value is % in good agreement with that obtained by Jones (1985) using a different % representation of the three-dimensional condensed matter state. In the interval $10 < B_{12} < 100$, their values can be fitted by the % expressions $E_{c} = 0.016B^{1.2}_{12}$ keV for $Z = 6$ and $E_{c} = % 0.16B^{0.7}_{12}$ keV for $Z = 26$ giving, \\begin{eqnarray} T_{c} & = & 4.6\\times 10^{3}B^{1.2}_{12}\\hspace{5mm}{\\rm K}, % \\hspace{1cm}Z=6\\nonumber \\\\ & = & 4.6\\times 10^{4}B^{0.7}_{12}\\hspace{5mm}{\\rm K},\\hspace{1cm}Z=26. \\end{eqnarray} These are to be seen in relation to other temperatures that are significant. Paper I contained a calculation of the rate of proton formation in the % electromagnetic showers formed by reverse electrons incident on the polar cap. % From the energy flux needed to produce a proton current density $J^{p} = % \\rho_{0}c$, we can infer a maximum steady-state polar-cap temperature, \\begin{eqnarray} \\bar{T} = \\left(T^{4}_{res} + \\frac{(-B\\cos\\psi)(1 - \\kappa)} {P\\sigma eW_{p}}\\right)^{1/4}\\hspace{5mm}{\\rm K}. \\end{eqnarray} In this expression, $W_{p}$ is the number of protons produced per unit primary % shower energy. We can approximate it initially by $W^{BH}_{p}$ which was % obtained in Paper I by using the zero-field Bethe-Heitler pair creation % cross-section with screening modified for the condensed matter density at the % neutron-star surface. Its values, given in Table 1 of that paper, can be % summarized conveniently in the intervals $B_{12} = 10^{1}-10^{2}$ and $Z = % 10-26$ by the expression \\begin{eqnarray} W^{BH}_{p} = 3.9\\times 10^{-4}\\langle Z\\rangle^{-0.6}_{sm} % B^{0.2}_{12}\\hspace{5mm}(mc^{2})^{-1}, \\end{eqnarray} in which the nuclear charge is the average in the region of the shower maximum. The angle $\\psi$ is that between ${\\bf \\Omega}$ and ${\\bf B}$. The % general-relativistic correction contained in the surface value of the Goldreich- % Julian charge density $\\rho_{0}$ is $\\kappa$ (see Muslimov \\& Harding 1997), and % $\\sigma$ is Stefan's constant. Equation (4) also contains a further temperature, $T_{res}$, which is the polar-cap temperature in the % absence of any reverse electron or photon energy flux (approximately the % observed whole-surface temperature corrected to the local proper frame). The % presence of $T_{res}$ assumes that there is a constant heat flow to the surface % driven by the very much higher temperature of the inner crust. With $T_{res} = % 0$ and $\\kappa = 0.15$, we find, \\begin{eqnarray} T_{max} = 5.1\\times 10^{5}\\langle % Z\\rangle^{0.15}_{sm}B^{0.2}_{12}\\left(\\frac{-\\cos\\psi}{P}\\right)^{1/4} \\hspace{5mm}{\\rm K} \\end{eqnarray} Equating this with $T_{c}$ allows us to estimate the minimum polar-cap magnetic % flux densities necessary to sustain the case (iii) boundary condition ${\\bf % E}_{\\parallel} \\neq 0$. These are, \\begin{eqnarray} B_{12} & = & 181\\left(\\frac{-\\cos\\psi}{P}\\right)^{1/4}\\hspace{1cm}Z = 6\\nonumber % \\\\ & = & 327\\left(\\frac{-\\cos\\psi}{P}\\right)^{1/2}\\hspace{1cm}Z = 26. \\end{eqnarray} Comparison with the median value of $B_{d12} = 2.8$ for the 63 nulling pulsars % listed by Wang et al shows that case (iii) can be widely realized only with % implausibly large deviations from a central dipole field. \\begin{table} \\caption{The table gives values of $W_{p}$ for high magnetic flux densities $B % \\geq B_{c}$ and nuclear charge $Z$. The effective total cross-sections % $\\sigma^{P}$ for pair creation in the GDR photon energy band are for transverse momenta below the lowest magnetic conversion threshold at $k_{\\perp} = 2$ mc and % are in units of barns. They are estimates obtained from Table 3 of Jones (2010b) by averaging over photon % polarization and over photon transverse momenta $k_{\\perp} = 1.0$ and $1.5$ mc, % and are equivalent to a mean free path for Coulomb pair creation lengthened by a % factor $\\eta_{p}$ compared with that for the modified zero-field Bethe-Heitler % cross-section. The Landau-Pomeranchuk-Migdal effect is not significant in the % GDR photon energy band} \\begin{tabular}{@{}rcccr@{}} \\hline $Z$ & $BB^{-1}_{c}$ & $\\sigma^{P}$ & $\\eta_{p}$ & $W_{p}$ \\\\ & & bn & & $(mc^{2})^{-1}$ \\\\ \\hline 10 & 1 & 0.27 & 1.14 & $2.3\\times 10^{-4}$ \\\\ & 3 & 0.071 & 2.95 & $3.4\\times 10^{-4}$ \\\\ & 10 & 0.017 & 6.3 & $3.6\\times 10^{-4}$ \\\\ 18 & 1 & 0.89 & 1.35 & $1.8\\times 10^{-4}$ \\\\ & 3 & 0.23 & 3.8 & $3.2\\times 10^{-4}$ \\\\ & 10 & 0.055 & 9.7 & $3.8\\times 10^{-4}$ \\\\ 26 & 1 & 1.86 & 1.45 & $1.3\\times 10^{-4}$ \\\\ & 3 & 0.48 & 4.2 & $2.7\\times 10^{-4}$ \\\\ & 10 & 0.11 & 11.3 & $3.4\\times 10^{-4}$ \\\\ \\hline \\end{tabular} \\end{table} However, the screening-modified zero-field Bethe-Heitler pair creation % cross-section is not obviously valid at magnetic flux densities of the order of % the critical field $B_{c} = 4.41\\times10^{13}$ G. Thus we have been obliged to % calculate the second-order bremsstrahlung and pair creation cross-sections using % Landau function solutions of the Dirac equation. The photoproduction of protons by giant-dipole resonance (GDR) decay is % determined by the total photon track length in the GDR band, centred on a % momentum $k = 40$ mc, which occurs almost entirely in the late stages of shower development. The track % length at these energies is limited by Coulomb and magnetic pair creation, also by Compton scattering the effect of which is almost always % to scatter the photon so that its transverse momentum component $k_{\\perp}$ % (perpendicular to the field) exceeds the threshold for magnetic conversion to % electron-positron pairs. We refer to Sections 2 and 3 of Paper I for a more % detailed description of these processes. Approximate values of the Coulomb pair creation cross-section at magnetic fields % $B\\geq B_{c}$ are given in Table 3 of Jones (2010b) and are the basis for the % values of $W_{p}$ given here in Table 1. These are not easily summarized as a % simple expression analogous with equation (5). The cross-section at $B = % 10B_{c}$ is at least an order of magnitude smaller the the modified zero-field % Bethe-Heitler cross-section but the effect on $W_{p}$ is not large because the % photon track length in this region is limited by Compton scattering. % Substitution into equation (4) gives values of $T_{max}$ that typically are % reduced by a factor of approximately $0.9$ from those of equation (6). But the % complexity of the second-order processes at $B\\sim B_{c}$ is such that we have % not reconsidered shower development and, specifically, have not allowed for the % production by cyclotron emission or Coulomb bremsstrahlung of photons with % $k_{\\perp}$ above the thresholds for magnetic conversion. This becomes % significant at $B = 10B_{c}$, as shown in Tables 2 and 4 of Jones (2010b), and % the high magnetic conversion transition rates reduce GDR-band photon track % length in the shower. The reduced $W_{p}$ values increase $T_{max}$. Owing to % this complexity at $B\\sim B_{c}$ there are inevitable uncertainties in our % estimates of $W_{p}$, but we believe that they do not seriously invalidate the % estimates of the minimum polar-cap magnetic flux densities needed to support the % case (iii) boundary condition and our conclusion that it can exist only in a % very small subset of pulsars. Our conclusions drawn from equation (3) are also % independent of the spectrum of the reverse electron-photon flux because $W_{p}$ % is almost completely independent of primary shower energy provided that is large % compared with the GDR energy. There appear to be no published estimates of the melting temperature of % condensed matter that are specific to very high magnetic fields. Consequently, % we are obliged to adopt the standard one-component Coulomb plasma expression % (see, for example, Slattery, Doolen \\& DeWitt 1980) which, with equation (1), % gives, \\begin{eqnarray} T_{m} = 1.0\\times 10^{4}Z^{2}_{v}Z^{-0.23}B^{0.4}_{12}\\hspace{5mm}{\\rm K}, \\end{eqnarray} in terms of an effective valence charge $Z_{v}$. This latter parameter % represents the fact that the deeply-bound Landau states certainly do not % participate significantly in the melting transition, but its estimation at % higher values of $Z$ is quite difficult. It is possible that the work of % Potekhin, Chabrier \\& Yakovlev (1997; see Fig. 1) could provide guidance % although it is at zero field and was directed toward a different problem. On % this basis, we assume for $Z = 26$ a value in the interval $Z_{v} = 10 - 15$. % In a typical polar-cap field, $B_{12} = 10$, the melting temperature is as low % as $T_{m} = 6\\times 10^{5}$ K for $Z_{v} = Z = 6$ but exceeds $10^{6}$ K for $Z % = 26$. Thus the state of the polar cap may be a sequence of melting and % solidifications. The order of magnitude of the shear modulus is a further % source of complexity. The standard (zero-field) expression for a body-centred cubic lattice (Fuchs 1936) can be adapted, with equation (1), to % give, \\begin{eqnarray} \\mu = 1.1\\times 10^{16}Z^{2}_{v}Z^{-0.93}B^{1.6}_{12}\\hspace{5mm}{\\rm erg} % \\hspace{2mm}{\\rm cm}^{-3}. \\end{eqnarray} However, the polar-cap gravitational constant is $g\\approx 2\\times 10^{14}$ cm % s$^{-2}$ so that any density inversion may well induce a form of Rayleigh-Taylor % instability. The final condensed-matter parameter that is important is the thermal % conductivity parallel with ${\\bf B}$, which is extremely high (see Potekhin % 1999). Thermal energy is deposited at shower maxima a distance $z_{p}$ below % the surface $z = 0$, which is defined as that separating condensed matter from % the atmosphere, and is then dissipated as blackbody radiation from the polar cap. The value of % $z_{p}$ in the high-density condensed matter of the neutron-star surface depends % on shower energy owing to the Landau-Pomeranchuk-Migdal (LPM) effect but for the % order of magnitude of the characteristic time we can assume $-z_{p} \\sim 10l_{r} % \\approx 10^{-3}$ cm using the radiation lengths $l_{r}$ given by equation (2) or in Table 1 of Paper % I. The characteristic time for shower energy input to produce a % surface-temperature fluctuation is then, \\begin{eqnarray} t_{cond} = \\frac{Cz^{2}_{p}}{2\\lambda _{\\parallel}} \\approx % 10^{-9}\\hspace{5mm}{\\rm s}, \\end{eqnarray} in which typical values of the parameters are the specific heat $C = 1.0\\times % 10^{12}$ erg cm$^{-3}$ K$^{-1}$ and the longitudinal coefficent of thermal % conductivity $\\lambda _{\\parallel} = 6\\times 10^{14}$ erg cm$^{-1}$ s$^{-1}$ K$^{-1}$. But for a surface temperature of $10^{6}$ K, the internal temperature gradient needed to conduct the radiated energy flux is extremely small, of the order of $10^{5}$ K cm$^{-1}$, within the condensed % matter at $z < 0$. In effect, heat is more easily conducted to greater depths % than radiated from the surface. Consequently, the Green function $G(z,z_{p},t)$ % giving the internal temperature distribution must be almost independent of % $z_{p}$ and very close to that for zero temperature gradient at the surface $z = % 0$. Thus a shower heat input $H\\delta(t)$ produces a temperature distribution % within the condensed matter at $z \\leq 0$, \\begin{eqnarray} T(z,t) = HG(z,t) \\approx \\frac{H}{(\\pi C\\lambda _{\\parallel}t)^{1/2}}\\exp\\left( \\frac{-Cz^{2}}{4\\lambda _{\\parallel}t}\\right) \\end{eqnarray} It is asymptotically $\\propto t^{-1/2}$ so that the polar-cap temperature % arising from a sequence of heating events, each producing a maximum temperature % $T_{max}$, has fluctuations away from $\\bar{T}$ whose magnitude is a function of % the time-scale concerned. Further discussion of this is deferred until our % consideration of observed polar-cap blackbody temperature in Section 5. \\section[]{Short time-scale instability} Under the assumption that the case (ii) boundary condition is satisfied, the % neutron-star surface has an atmosphere in local thermodynamic equilibrium (LTE) % with approximate scale height $l_{A} = (\\tilde{Z} + 1)k_{B}T/Mg \\sim 10^{-1}$ cm % at temperature $T = 10^{6}$ K, where $\\tilde{Z}$ and $M$ are the mean ion charge % and mass of the partially ionized atom. The expression for the chemical potential of an ideal Boltzmann gas gives the % LTE atmospheric number density at $z = 0$, \\begin{eqnarray} N_{A}(0) = \\left(\\frac{Mk_{B}T}{2\\pi \\hbar^{2}}\\right)^{3/2} {\\rm e}^{-E_{c}/k_{B}T}, \\nonumber \\end{eqnarray} for atmospheric temperatures such that $N_{A}(0) \\ll N$. Its order of magnitude % is $N_{A}(0) \\approx 10^{32}\\exp(-E_{c}/k_{B}T)$ cm$^{-3}$ for $M = 56m_{p}$. We estimate the critical temperature $T_{c}$ which must not be exceeded if the % case (iii) boundary condition is to be valid by equating the flux of ions in such an LTE atmosphere that are incident on the % neutron-star surface with the ion flux needed to give an open magnetosphere % current density equal to $\\rho_{0}c$. This gives $N_{A}(0) \\approx 10^{14}$ % cm$^{-3}$ and $k_{B}T_{c} = 0.025E_{c}$. But the extent to which the atmosphere % can be described as thin is very temperature-dependent. Thus for $T = 2T_{c}$ % and $Z = 26$, $N_{A}(0)$ is of the order of $10^{23}$ cm$^{-3}$ and the whole % atmosphere contains ions equivalent to about $10^{-1}l_{r}$. At a density $N_{A}(0) > 10^{23} - 10^{24}$ cm$^{-3}$, the Boltzmann gas estimate of $l_{A}$ assumed here is unreliable and it becomes necessary to allow for the interaction % of protons with the ion atmosphere. If instabilities in plasma acceleration with time-scales as short as$\\sim % 10^{-4}$ s are considered, the temperature at $z = 0$ is constant apart from % very small fluctuations, as shown by the Green function given at the end of the % previous Section. Consequently, the temperature distribution and number density % of ions in the atmosphere are also constant in time and the atmosphere is in % local thermodynamic equilibrium in the interval $0 < z < z_{1}$, where $z_{1}$ % is the top of the ion atmosphere, defined as the surface of last scattering. For % time-scales of the order of $1 s$, the temperature and depth of the atmosphere % can change appreciably but there is no doubt that it is always in local % thermodynamic equilibrium. The processes of proton formation by GDR decay followed by diffusion to the surface were described in Paper I, which assumed a % very thin atmosphere. Therefore, we need to consider here, in more detail, % diffusion in the atmosphere $0 < z < z_{1}$. The number of protons is very many % orders of magnitude smaller than that of ions so that the properties of the % atmosphere, its scale height and equilibrium electric field given by the % electrical neutrality condition, are determined solely by the latter. With the assumption of a single ion component we can determine the equilibrium % electric field within the LTE atmosphere in the presence of a gravitational % acceleration $g$ and find that the proton potential energy is, \\begin{eqnarray} \\left(1 - \\frac{A}{\\tilde{Z} + 1}\\right)m_{p}gz = \\alpha m_{p}gz \\end{eqnarray} for ions of charge $\\tilde{Z}$ and mass number $A$, and that this transports protons to $z > z_{1}$ from which region they are accelerated to % relativistic energies in preference to the ions. Proton diffusion at a rate % greater than is needed for a current density $J^{p} = \\rho _{0}c$ therefore cuts % off ion acceleration and produces a thin electrically neutral proton atmosphere % at $z > z_{1}$. There appears to be no reason why this atmospheric structure % should be disturbed by turbulent mixing. The proton average potential energy at $z < 0$ is close to $m_{p}gz$ but the % jump bias, $m_{p}ga_{s}/k_{B}T$ for ion separation $a_{s}$, is too small to be % significant because $m_{p}gz_{p}/k_{B}T \\ll 1$. The atmospheric proton density % at $z = 0$ is necessarily some orders of magnitude smaller than the density at % $z < 0$ owing to the density discontinuity at the surface. Thus diffusion to the surface at $z = 0$ is little changed by the presence of an % atmosphere that is not necessarily very thin. Movement of the protons through the ion atmosphere is effected by the chemical % potential gradient which is a function of ion number density. At high densities, % as in the solid at $z < 0$, this is determined principally by entropy but as the % density reduces, the potential energy given by equation (12) becomes the more % important factor. In effect, the motion changes from a random walk to a drift % velocity. The dividing ion density is given by the condition $\\lambda _{R} = N_{A}^{-1/3}$ and is $\\approx 4\\times 10^{22}$ cm$^{-3}$, below which it is appropriate to define a % drift velocity determined by the local value of the mean free path $\\lambda % _{R}$ derived from the total cross-section for Rutherford back-scattering, \\begin{eqnarray} \\bar{v} \\approx -\\alpha g\\lambda _{R}\\left(\\frac{m_{p}}{k_{B}T}\\right)^{1/2}, \\end{eqnarray} in which, \\begin{eqnarray} \\lambda _{R} = \\frac{1}{\\pi N_{A}} \\left(\\frac{k_{B}T}{\\tilde{Z}e^{2}}\\right)^{2}. \\end{eqnarray} The drift velocity is $\\bar{v} \\approx 2.6$ cm s$^{-1}$ at this density and $T = 10^{6} K$. The consequence for a $T = 2T_{c}$ atmosphere is that the diffusion time to $z = % z_{1}$ is increased, by some orders of magnitude, compared with the values % assumed in Paper I which were for diffusion through the condensed matter only. % Also, the change from random walk to drift velocity in the atmosphere has a % significant effect on the distribution of diffusion times because it removes % the long tail in the distribution that is a feature of random walks. Some effort has been made here to confirm that there is outward proton diffusion % because the process is of crucial importance. It was shown in Paper I that plasma acceleration is unstable, consisting of alternate bursts of proton and % ion acceleration, with time-scale determined by the time for proton diffusion % from $z = z_{p}$ to $z = z_{1}$. The polar-cap current density ${\\bf J}$ is % essentially time-independent for relativistic flow (as is the charge density) so % that there is no fluctuation in polar-cap electric field other than that arising % from the reverse flow of photo-electrons from accelerated ions (see Section 5). % The instability is in the nature of the plasma accelerated and remains % adequately described by the analysis given in Paper I which will not be repeated % here. The maximum acceleration potential difference, given a current density $J = % \\rho _{0}c$, is almost entirely dependent on the Muslimov \\& Tsygan (1992) % general-relativistic correction. There is a separate contribution arising from % ion inertia (Michel 1974) but it is important only at very low altitudes. % Neglecting this, the maximum potential difference is approximately \\begin{eqnarray} |\\Phi _{max}| \\approx\\frac{2\\pi^{2}\\kappa B_{d}R^{3}}{c^{2}P^{2}} \\approx % 7\\times 10^{12} \\frac{\\kappa B_{d12}}{P^{2}}\\hspace{5mm} {\\rm volts} \\end{eqnarray} and, unless restricted by pair creation at lower altitudes, is reached at an % altitude $z$ of the order of the neutron-star radius $R$ which is roughly two % orders larger than the polar-cap radius. We refer to Harding \\& Muslimov (2001, % 2002) for a complete account. We assume here that for this current density, % spontaneous pair creation by curvature radiation is not possible. Then the % primary source of any positron component in ${\\bf J}$ can only be the reverse % flow of photo-electrons from accelerated partially-ionized atoms as discussed in % Paper I. But there is an essential difference here in that protons in the very % low density region at $z\\sim z_{1}$ are almost completely ionized (here, we % refer to Fig.1 of Potekhin et al 2006) so that both track length and energy flux % of reverse-flow electrons from photo-ionized hydrogen atoms are negligibly % small. Thus positron production in any interval of proton acceleration is % negligible. This is also true for ions of low atomic number $Z \\sim 4-5$ which % are completely ionized, but for higher $Z$, the reverse flux of inward % accelerated electrons produces polar-cap ICS photons, as would outward % accelerated electrons in case (i). Pairs are produced by photons that are % scattered to transverse momenta above the magnetic conversion thresholds. The % only difference is that the photons are inward moving. But even if spontaneous % CR pair formation is not possible, positrons accelerated outward will produce CR pairs at higher altitudes, as in % case (i), though superimposed on a flux of ions. In case (i), the reverse flux of positrons must form protons which are not % accelerated but form an atmosphere at $z > 0$ whose equilibrium is defined by % various diffusion processes perpendicular to the magnetic flux. Backward moving % photons from the electron or positron showers are a further source of pair % creation in each of cases (i) - (iii). These arise from the decay of residual % nuclei following proton or neutron emission in GDR formation, and from % $(n,\\gamma)$ reactions. Those photons that are not absorbed in the more dense % part of the atmosphere can produce pairs if their transverse momenta exceed the % threshold. The LPM effect, which in the high-density condensed matter at $z < % 0$ is significant for electrons of more than $10$ GeV or photons of more than % $2$ GeV, is less important in the atmosphere within which there will be some GDR % formation. But we have been unable to make a satisfactory quantitative estimate % of pair formation through this process. \\section[]{Medium time-scale instability} Instability on time-scales some orders of magnitude longer than those of Section % 3 is also of interest and can appear as a fluctuation in the charge of nuclei % reaching the polar-cap surface which we shall define here to be always at $z = 0$ although it may move with respect to coordinates fixed at the centre of % the star. Ions of initial charge $Z_{a}$ move upward through the region of the shower % maxima at $z_{p}$ and, with nuclear charge reduced to $Z_{b}$ by GDR formation % and decay, enter the atmosphere at $z > 0$. We wish to find if there are % conditions under which the local average nuclear charge $Z_{0}(z)$ fails to be a % time-independent function of position with limits $Z_{a} \\geq Z_{0}(z) \\geq % Z_{b}$. The work of this Section replaces that of Section 4.1 in Paper I which % contains an error. The longer time-scales here enable us to assume the proton and ion current % densities $J^{p}$ and $J^{z}$ are the time averages of those described in the % previous Section. There is, of course, a distribution of discrete nuclear % charges in the shower region, but we shall work in terms of the average % $Z(z,t)$, the corresponding number density $N(z,t)$ given by equation (1), and % the velocity $v(z,t)$ with which nuclei approach the polar-cap surface at $z = % 0$. Proton formation by GDR decay occurs predominantly in the very late stages % of shower development, as explained in Section 1, and so we shall assume that it % is confined within limits $z_{a}$ and $z_{b}$ and is defined by the normalized % function $g_{p}(z)$, \\begin{eqnarray} \\int^{z_{b}}_{z_{a}}g_{p}(z)dz = 1. \\end{eqnarray} The physical basis for our study of the system is that the total numbers of % nuclei and of protons (bound or unbound) are conserved. Because both neutrons % and protons are produced in GDR decay, we can make the approximation of % neglecting the effect of $\\beta$-transitions. Therefore, the continuity equations are, \\begin{eqnarray} \\frac{\\partial N}{\\partial t} + \\frac{\\partial(Nv)}{\\partial z} = 0, \\end{eqnarray} and with neglect of the relatively short proton diffusion time so that within % medium time-scales all protons produced in the shower are assumed to be promptly % accelerated from the atmosphere at $z \\approx z_{1}$, \\begin{eqnarray} \\frac{\\partial (NZ)}{\\partial t} + \\frac{\\partial (NZv)}{\\partial z} = -g_{p}(z)J^{p}(t). \\end{eqnarray} Because the reverse electron flux from photo-electric ionization is the source % of the showers, we shall find it convenient here, as in Paper I, to introduce % the parameter $K$ which is a function of the atmospheric nuclear charge and is % the number of protons produced per unit positive nuclear charge accelerated. % With the representation of equation (1) in the form $N = CZ^{\\gamma}$, equations (17) and (18) can be combined to give, \\begin{eqnarray} C\\int^{z_{b}}_{z_{a}}dz\\left(Z^{\\gamma}\\frac{\\partial Z}{\\partial t} + vZ^{\\gamma}\\frac{\\partial Z}{\\partial z}\\right) = -J^{p}(t). \\end{eqnarray} It has the obvious time-independent solution, \\begin{eqnarray} Z_{a} - Z_{b} = KZ_{b}, \\end{eqnarray} and the time-independent velocity of nuclei is such that $Z^{\\gamma}_{0}(z)v_{0}(z)$ is independent of $z$. A natural fluctuation away from $Z_{0}$ would be of the form $Z(z,t) = Z_{0} + % \\delta Z(z,t)$ with, \\begin{eqnarray} \\delta Z = \\eta\\left(Z_{a} - Z_{0}(z)\\right){\\rm e}^{{\\rm i}\\omega t}, \\end{eqnarray} giving similar fluctuations in $J^{p}$, $N$ and $v$, in which $\\eta$ is % infinitesimal and independent of $z$ and $t$. From equation (17) we then have, \\begin{eqnarray} \\delta\\left(Z^{\\gamma}v\\right) = -{\\rm i}\\omega\\eta\\chi(z){\\rm e}^{{\\rm i}\\omega % t}, \\end{eqnarray} where, \\begin{eqnarray} \\chi(z) = \\gamma\\int^{z}_{z_{a}}dz^{\\prime}\\left(Z_{a} - Z_{0}\\right)Z^{\\gamma % -1}_{0}. \\end{eqnarray} Substitution into equation (19) with the retention of terms of first order in % $\\eta$ gives, \\begin{eqnarray} {\\rm i}\\omega\\eta {\\rm e}^{{\\rm i}\\omega % t}\\int^{z_{b}}_{z_{a}}dz\\left(Z^{\\gamma}_{0}\\left(Z_{a} - Z_{0}\\right) - % \\chi(z)\\frac{\\partial Z_{0}}{\\partial z}\\right) - \\nonumber \\\\ % Z^{\\gamma}_{0}v_{0}\\left(Z_{b} - Z_{a}\\right) \\eta {\\rm e}^{{\\rm i}wt} = - \\frac{1}{C}\\delta J^{p}(t). \\end{eqnarray} The current densities, averaged over short time-scales, are $J^{z} = NZv(0,t)$ % and, from the definition of $K$, $J^{p} = KJ^{z}$. The fluctuation away from the time-independent solution is $\\delta J = \\delta % J^{p} + \\delta J^{z} = 0$, which we assume to be maintained by the boundary conditions on $\\Phi$. To express $\\delta J^{p}$ in terms of $\\delta Z$ we represent the $Z$-dependence of $K$ in % the vicinity of $Z_{b}$ as a power law $K = K_{0}Z^{\\nu}(0,t)$. We find, after eliminating the velocity fluctuation $\\delta v(0,t)$, that \\begin{eqnarray} \\frac{1}{C}\\delta J^{p} = \\nu Z^{\\gamma}_{0}v_{0}\\frac{K}{K + 1}\\delta Z(0,t). \\end{eqnarray} The relationship with equation (24) is established by noting that, \\begin{eqnarray} \\delta Z(0,t) = \\delta Z(z_{b},t-\\tau) = \\eta\\left(Z_{a} - Z_{b}\\right){\\rm e}^{{\\rm i}\\omega(t-\\tau)}, \\end{eqnarray} where $\\tau$ is the time interval of nuclear movement from $z = z_{b}$ to $z = % 0$. From equations (24) - (26), the equation whose root $\\omega$ we require can % be expressed as, \\begin{eqnarray} \\frac{{\\rm % i}\\omega}{Z^{\\gamma}_{b}v_{0}(z_{b})}\\int^{z_{b}}_{z_{a}}dzZ^{\\gamma-1}_{0} \\left((1 + \\gamma)Z_{0} - \\gamma Z_{b}\\right)\\left(Z_{a} - Z_{0}\\right) \\nonumber \\\\ = - (Z_{a} - Z_{b}) - \\frac{\\nu K}{K + 1}(Z_{a} - Z_{b}) {\\rm e}^{-{\\rm i}\\omega \\tau}, \\end{eqnarray} in which the function $\\chi(z)$ has been removed by partial integration. The only assumptions we need make about the depth distribution of proton % formation in the late stage of shower development are that it is small outside % the interval $z_{a} < z < z_{b}$ and that $z_{b} - z_{a}$ is smaller than % $|z_{b}|$ though not necessarily so by as much as an order of magnitude. A % suitable function would be, \\begin{eqnarray} g_{p}(z)) = \\frac{2}{z_{b} - z_{a}}\\sin^{2}\\left(\\frac{\\pi (z - z_{a})} {z_{b} - z_{a}}\\right), \\end{eqnarray} in terms of which we could express $Z_{0}$ as an explicit function of $z$, which would be needed to obtain numerical values for the root $\\omega = \\omega % _{1} + {\\rm i}\\omega _{2}$ of equation (27). However, we are primarily concerned here not with growth rates but with the % boundary between stability and instability and so shall not proceed with % numerical solutions for $\\omega$. Fortunately, it is possible to obtain a sufficient condition for the existence % of instability independently of the form of $g_{p}$. Provided $-1 < \\gamma < % 0$, which is clearly satisfied, we can see directly from the values of the % integrand in equation (27) at the limits and from its lack of an extremum that, \\begin{eqnarray} \\int^{z_{b}}_{z_{a}}dz\\left((1 + \\gamma)Z_{0} - \\gamma Z_{b}\\right) \\left(\\frac{Z_{a} - Z_{0}}{Z_{0}}\\right) \\left(\\frac{Z^{\\gamma}_{0}}{Z^{\\gamma}_{b}}\\right) \\nonumber \\\\ < (z_{b} - z_{a})(Z_{a} - Z_{b}), \\end{eqnarray} and therefore that equation (27) can be replaced by the inequality, \\begin{eqnarray} {\\bf i}\\omega\\tau\\left(\\frac{z_{a} - z_{b}}{|z_{b|}}\\right) > -1 - \\frac{\\nu K}{K + 1} {\\rm e}^{-{\\rm i}\\omega\\tau}. \\end{eqnarray} Thus$\\omega _{1,2}$ satisfy the inequalities, \\begin{eqnarray} \\omega _{1}\\tau\\left(\\frac{z_{b} - z_{a}}{|z_{b}|}\\right) & > & \\frac{\\nu K}{K + 1}{\\rm e}^{\\omega _{2}\\tau}\\sin\\omega _{1}\\tau \\nonumber \\\\ \\omega _{2}\\tau \\left(\\frac{z_{b} - z_{a}}{|z_{b}|}\\right) & < & 1 + \\frac{\\nu K}{K + 1}{\\rm e}^{\\omega _{2}\\tau}\\cos\\omega _{1}\\tau. \\end{eqnarray} From the first of these, we can see that the real part of $\\omega$ must satisfy % $\\omega _{1}\\tau > \\pi/2$, provided the proton formation region in the shower is % sufficiently compact that, \\begin{eqnarray} \\frac{\\nu K}{K + 1}{\\rm e}^{\\omega _{2}\\tau} > \\frac{\\pi}{2}\\left(\\frac {z_{b} - z_{a}}{|z_{b}|}\\right). \\end{eqnarray} With greater compactness, $\\omega _{1}\\tau \\rightarrow \\pi$, as would be % expected. The second inequality gives the condition for $\\omega _{2} < 0$, that % is, fluctuation growth. At the threshold where $|\\omega _{2}\\tau| \\ll 1$, and % for $z_{b} - z_{a} \\ll |z_{b}|$, it is simply $\\nu K > K + 1$. The $Z$-dependence of $K$ was represented in Paper I as an approximate power law, $K = K_{0}Z^{\\nu}$ with $\\nu = 0.85$ for $Z \\geq 10$ but with the % reservation that this would certainly become invalid for an atmosphere of ions % with $Z \\sim 5$ which would be almost completely ionized. Therefore, we should % anticipate quite large values of $\\nu$ at small $Z$, sufficient to give the % unstable behaviour found here. Stability clearly depends on the energy flux % from photo-ionization. Low fluxes and moderate rates of proton formation in the % GDR region of the showers, such that the nuclear charge inferred from equation % (20) is $Z_{b} \\approx 10$, allow a stable time-independent progression of % nuclear charge as a function of depth below the polar-cap surface as given by % $Z_{0}(z)$. But larger values of $K$ and smaller $Z_{b}$ lead to instability. (The value of $Z_{a}$ is unimportant, % provided it is not small, and for order of magnitude estimations in this paper % we have assumed $Z_{a} = 26$ unless otherwise stated.) Unfortunately, our analysis of the instability of nuclear movement to the % polar-cap surface is limited to small fluctuations and does not extend to large % amplitudes. But it is not difficult to see the form it would take. An % atmosphere of high-$Z$ ions at time $t$ produces a high reverse-electron energy % flux which creates a layer of low-$Z_{b}$ ions in the shower maximum region % $z_{a} < z < z_{b}$. These flow toward the surface at $z = 0$ and form a low-$Z$ % atmosphere at time $t + \\tau$ which produces a low reverse-electron energy flux % and correspondingly large values of $Z_{b}$. The ions accelerated alternate % between high and low-$Z$ values. The basic unit of time is given by the time % $t_{rl}$ for the emission at the Goldreich-Julian current density of one % radiation length of ions, \\begin{eqnarray} t_{rl} = 2.1\\times 10^{5}Z^{-1}B^{-1}_{12}(- P\\sec \\psi) \\nonumber \\\\ \\left( \\ln \\left(12Z^{1/2}B^{-1/2}_{12}\\right)\\right)^{-1}\\hspace{5mm}{s}. \\end{eqnarray} The high-$Z$ intervals are subject to short time-scale instability as described % in Section 3, but low-$Z$ intervals have little or no reverse-electron flux and % therefore no possibility of significant positron acceleration and % electron-positron pair production. The instability outlined here is of quite simple form, but there are % complicating factors that have been mentioned earlier in the Section 2 % consideration of polar-cap parameters. Evaluations of the melting temperature % discussed following equation (8) show that there is every possibility that the % condensed matter state below the atmosphere may be liquid, or a solid close to % melting with a high self-diffusion rate. This would have no effect on short % time-scale instability but could complicate the behaviour of the system over % medium time-scales. The melting temperature is $Z$-dependent so that a density % inversion is possible, the upper layer of higher-$Z$ being either liquid or % solid. In the liquid case, we must anticipate Rayleigh-Taylor instability, % which may also exist in the solid case because its shear modulus (see equation % 9) may not be adequate to maintain mechanical stability. All these processes % are occurring at depths $|z_{p}| \\sim 10^{-3}$ cm but over a polar cap whose % radius is of the order of $10^{4}$ cm. Consequently, a further complication is % that different polar-cap areas are unlikely to be in phase with each other. % These problems are considered further, though necessarily in a qualitative way, % in Section 5. \\section[]{Nulls subpulse drift and polar-cap coherence} The time-dependent phenomena considered in Sections 3 and 4 are local and % one-dimensional because both shower depth $z_{p}$ and atmospheric scale height % are very many orders of magnitude smaller than the polar-cap radius. This % introduces the question of whether or not there is coherence over the whole polar-cap area. Both instabilities are dependent on the parameter $K$ which is a function of the surface nuclear charge $Z(0,t)$ and to a lesser % extent of surface temperature and acceleration field. For this reason alone, we % see no case for assuming complete coherence and anticipate that the polar cap we % describe has zones of proton and ion emission which cannot be stationary, the % total areas of each being determined, approximately, by the average value of % $K$. For neutron stars that are unable to support spontaneous pair production by % curvature radiation, the proton zones have no reverse electron flux, do not % support pair production, and hence merely produce an accelerated one-component % plasma as described in Section 3. But the ion zones will support ICS pair % production and so appear as moving sources within the polar cap. In this % Section, we shall attempt to compare possible forms this motion might take with % some of the observed phenomena in radio pulsars. The distinction between the average pulse profile and individual sub-pulses % within it was noted almost immediately following the discovery of pulsars (see, for example, Smith 1977). The amplitude and form of successive sub-pulses % can vary in times of the order of the rotation period and in some pulsars, % observed with higher resolution, sub-pulses have micro-structure of $10^{-4}- % 10^{-3}$ s time-scale. There are also more organized phenomena, and for recent % extensive surveys of these we refer to Wang et al (2007) in the case of pulse % nulls and to Weltevrede, Edwards \\& Stappers (2006) for sub-pulse drift. There % is a move toward a consensus (Lyne et al 2010) that quasi-periodic switching % between magnetospheric states with different spin-down rates is the basis of % mode-changing. These authors even suggest that it is the source of a large % component of pulsar timing noise. However, it is also true that the subpulse characteristics observed in a small number of pulsars are quite singular, but here discussion is confined to % the general features of subpulses. Given the properties of the medium time-scale instability described in Section % 4, it is natural to associate with nulls those intervals in which high-$Z$ ion % zones either do not exist or are confined to areas of the polar cap from which % radiation produced by the plasma is not visible to the observer. For % sufficiently low values of $Z_{b}$, ions are accelerated from the surface % completely ionized so that there is no reverse-electron flux and hence no pair % creation and radio emission. The current density is little changed in the open % sector of the magnetosphere but the absence of pair creation means that the % particle content near the light cylinder is quite different as are the % components of the momentum density or stress tensor on any spherical surface in % this region centred on the star. It is therefore not surprising that the % spin-down torque is reduced during the interval of a null, as has been observed % in PSR B1931+24 (Kramer et al 2006). Neutron stars with small $K$ such that % $Z_{b} \\approx 10$ are likely to have $\\nu < 1$ and so will have a steady-state % progressively reducing nuclear charge $Z_{0}(z)$, no medium time-scale % instablity, and no long-term nulls. But the questions about melting and % mechanical stability of the polar cap which were mentioned at the end of Section % 4 remain and the whole system is so complex and difficult to describe in % physical terms that we are unable to give useful quantitative predictions. % However, incomplete nulls, having a low but detectable level of emission should % be observed. It is also unsurprising that nulls are observed to be not % completely random (Redman \\& Rankin 2009). To some extent and unfortunately, the same remarks have to be made about the % effects of short time-scale instability. For temperatures $T > 2T_{c}$, the proton diffusion time is much longer than in the case of the very thin % atmosphere assumed in Paper I but remains difficult to calculate with complete % confidence. In order to describe the polar cap it is necessary to introduce % coordinates ${\\bf u}(z)$ on a surface perpendicular to ${\\bf B}$. As in Paper I, the proton current density at any point ${\\bf u}(0)$ can be % related to the ion current density through the definition of $K$, \\begin{eqnarray} J^{p}({\\bf u},t) + \\tilde{J}^{p}({\\bf u},t) = % K\\int^{t}_{-\\infty}dt^{\\prime}f_{p}(t - t^{\\prime})J^{z}({\\bf u},t^{\\prime}), \\end{eqnarray} in which $f_{p}$ is the normalized proton diffusion-time distribution. Without % the $\\tilde{J}^{p}$ term, whose significance is described below, this would be a % homogeneous Volterra equation of the second kind having no non-zero % square-integrable solution. (The time-dependence of $K$ arising from the % temperature dependence of the LTE ionic charge $\\tilde{Z}$ is neglected here.) The approximate expression for $f_{p}$ given in Paper I (equation 23) assumed a % random-walk diffusion through the condensed matter at $z < 0$ and so is not % valid for an atmosphere with properties given by equations (13) and (14). For $T > 2T_{c}$, most of the diffusion time is in the drift-velocity phase, in % which case the time distribution would be more suitably approximated by a % normalized gaussian function centred at $t - \\tau _{p}$ or, in the limit, by $f_{p}(t - t^{\\prime}) = \\delta(t - t^{\\prime} - \\tau _{p})$, where $\\tau _{p}$ is derived from equation (13) of Section 3. The quantity $\\tilde{J}^{p} = 0$ within intervals for which $J^{p} < 1$ and at % other times, when $J^{p} = J$, represents the storage of excess protons in the % atmosphere at $z > z_{1}$. The total current density $J$ is fixed by the % boundary conditions and at $z = 0$ differs little from the Goldreich-Julian % value. Thus $J^{p} = J$ and $J^{z} = 0$ until the instant at which the % atmosphere is exhausted and $J^{p}$ falls to some residual value $J^{p} < J$. % Then ion flow recommences almost immediately (in a time much shorter than the % growth time for spontaneous curvature radiation pairs) and continues until % proton diffusion grows sufficiently to re-form the proton atmosphere. The % durations of the time intervals for ion and proton emission are those for % electron-positron pair creation or otherwise, and labelled $\\tau _{ee}$ and % $\\tau _{gap}$ are given by, \\begin{eqnarray} J & = & K\\int^{\\tau _{ee}}_{0}dt^{\\prime}f_{p}(\\tau _{ee} - % t^{\\prime})J^{z}(t^{\\prime}) \\nonumber \\\\ \\tau _{gap} & = & \\frac{K}{J}\\int^{\\tau _{ee}}_{0}dt^{\\prime}J^{z}(t^{\\prime}) \\end{eqnarray} Therefore $\\tau _{ee}$ and $\\tau _{gap}$ both depend on the smallness of $f_{p}$ % for small values of $t - t^{\\prime}$. In the $\\delta$-function limit for % $f_{p}$, we have $\\tau _{gap} = K\\tau _{ee}$. Estimates of the parameter $K = K_{0}Z^{\\nu}$ were given in Paper I, but are repeated here, \\begin{eqnarray} K_{0} & = & 2.8\\langle Z_{sm}\\rangle^{-0.76}B^{0.62}_{12}T^{-1.0}_{6}, \\hspace{1cm}B_{12} > 1, \\nonumber \\\\ K_{0} & = & 1.6\\langle Z_{sm}\\rangle^{-0.76}T^{-1.0}_{6}, \\hspace{1cm} B_{12} < 1, \\end{eqnarray} with $\\nu = 0.85$ for $10 < Z < 26$. In these expressions, $T$ is not the local surface temperature but is an average % for radiation emitted over the whole polar-cap, and $B$ is the actual polar-cap % magnetic flux density. For $B_{12} \\gg 1$, the values of $K_{0}$ given need to % be modified, though not greatly, to allow for the high-$B$ values of $W_{p}$ % contained in Table 1. Guided by the estimates contained in equations (6) and (7), we anticipate that % except for a very small number of neutron stars, the surface temperature is at % all times $T > T_{c}$ so that the case (ii) boundary condition is always % maintained and ion emission is never temperature-limited. It is limited instead % by the fact that the proton atmosphere forms above the ions and the protons are % preferentially accelerated, as described in Section 3. The only effect of the % temperature variations described by equation (11) that occur as a result of % alternating proton and ion emission is to change the density of the LTE ion % atmosphere. It is also worth observing that the local nature of equation (34) is % not disturbed by the presence of ${\\bf E}\\times{\\bf B}$ drift above the polar % cap. Although this slightly displaces a reverse electron shower relative to the % point from which the ion was accelerated, it has no effect on ion emission, % which is not temperature-dependent. In view of the time-variation described by equations (34) and (35), is it % possible to imagine organized rather than chaotic motion of ion zones on the % polar cap? An example of chaotic motion would be the existence of very many % small zones without obvious organization. Two simple organized cases would be motion along a slot of constant rotational % latitude and circular motion at constant $u(0)$. With the Deshpande \\& Rankin % (1999) analysis of PSR B0943+10 in mind, we consider circular motion. Equations % (34) and (35) are local in ${\\bf u}$ and contain no information that can % determine an ion zone movement velocity $\\dot{\\bf u}$. The quantities that are essentially constant are $K$ and, apart % from the effect of varying LTE atmospheric temperature, $\\tau _{p}$. Thus the % circulation time for $n$ ion zones, is $\\hat{P}_{3} = nP_{3}$, where $P_{3}$ is % the band separation in the usual notation by which subpulse drift is described, and here is given by \\begin{eqnarray} P_{3} = \\tau _{gap} + \\tau _{ee} = (K + 1)\\tau _{p} \\end{eqnarray} in the $\\delta$-function limit of the diffusion function $f_{p}$. In this % system, the velocity $\\dot{\\bf u}$ is determined by $n$ for fixed values of $K$ % and $\\tau _{p}$. For the same reason, drift direction is also unspecified and there is nothing to % preclude a change in $n$ or a reversal following some disturbance to the % polar-cap surface such as might be a consequence of the kind of mechanical % instability briefly described in Section 4. The drift time is determined % principally by diffusion through the more dense layers of the atmosphere. From % equations (13) and (14) it is, \\begin{eqnarray} \\tau _{p} = \\int^{z_{1}}_{0}\\frac{dz}{\\bar{v}(z)} = \\left(\\frac{l_{A}}{\\alpha % g\\lambda _{R}(0)}\\right)\\left(\\frac{k_{B}T}{m_{p}}\\right)^{1/2}. \\end{eqnarray} In Section 3 we observed that atmospheric density at $z = 0$ is an exponential % function of $T > T_{c}$ so that its value is essentially unpredictable. % Therefore, it is likely that the density discontinuity at $z = 0$ is small and that the reverse-electron showers may be contained entirely within the % atmosphere. In this case, the lower integration limit in equation (38) must be % replaced by $z_{sm} > 0$. The diffusion time is then not directly a function of % $B$, but is dependent on surface gravity and on $\\tilde{Z}$. But its % distribution for all pulsars should be compact. The distribution in the values % of $K$, given by equation (36), is probably the more important source of the % observed spread in the values of $P_{3}$. Detailed calculation of $\\tau _{p}$ has not been attempted here. In particular, % our estimates of $\\lambda _{R}$ and $l_{A}$ are subject to some uncertainty. The fact that $P_{3}$ is constant for a particular pulsar whereas the % circulation time $\\hat{P}_{3}$ is dependent on $n$ could in principle allow % comparison with ${\\bf E}\\times {\\bf B}$ drift-velocity polar-cap models in which % $\\hat{P}_{3}$ is constant. It is also worth considering the effect of % mechanical instability in allowing leakage of protons (or low-$Z$ ions) to % limited areas of the surface as described at the end of Section 4. The excess % protons form a localized atmosphere of greater depth than in surrounding areas, % so that ion zones (with consequent pair creation) do not form there until it is % exhausted. The observer of radio emission probably sees plasma from no more % than a strip of polar-cap area at roughly constant rotational latitude so it is % possible that intervals of null emission would be be observed at fixed longitude in a sequence of subpulse bands. In a survey of $187$ pulsars, selected only by signal-to-noise ratio, Weltevrede et al (2006) found that reversals do occur and that roughly equal % numbers of pulsars have subpulse drift to smaller or greater longitudes. They % also comment that subpulse drift is so common a phenomenon that it cannot depend % on pulsar parameters having extraordinary values. The band separation $P_{3}$ is independent of radio frequency, and is not % correlated with period $P$, age $P/2\\dot{P}$, or with the inferred dipole field % $B_{d}$. Measured values of $P_{3}$ given for 77 pulsars in Table 2 of % Weltevrede et al are mostly contained within a single order of magnitude, $1 < % P_{3} < 10$ s although the distribution has a small tail extending to $\\sim 20$ % s. This is quite compact (for a neutron-star parameter) and is not inconsistent % with the interpretation of subpulse drift given here and with equation (37). The range of $P_{3}$, with estimates of $K$ found from equations (36), indicate % $\\tau _{p} \\sim 10^{-1} - 10^{0}$ s and $N_{A}(0) = 10^{23} - 10^{24}$ % cm$^{-3}$, values that are by no means unreasonable. It is necessary to compare our polar-cap model with two other sets of pulsar % observations. Given that protons are the major fraction of particles % accelerated and produce no electron-positron pairs, it is worthwhile considering % the radio luminosity $L_{\\nu}$. The order of magnitude of this can be expressed % as, \\begin{eqnarray} L_{\\nu}\\Delta\\Omega\\Delta\\nu \\approx \\frac{2\\pi^{2}R^{3}B_{d}}{ceP^{2}}\\epsilon \\approx \\left(\\frac{1.4\\times 10^{30}B_{d12}}{P^{2}}\\right)\\epsilon \\end{eqnarray} in terms of the neutron-star radius $R$ and the inferred dipole field $B_{d}$. From this we can estimate the energy $\\epsilon$ radiated into solid angle % $\\Delta\\Omega$ within bandwidth $\\Delta\\nu$ per unit charge (baryonic or % leptonic) accelerated at the polar cap. Using the 400 MHz luminosities listed in the ATNF catalogue (Manchester et al 2005) and a % bandwidth $\\Delta\\nu = 400$ MHz, we find by evaluating equation (37) for a small % sample (B0826-34, B0834+06, B0943+10, B0950+08, B1055-52, B1133+16, B1929+10) % that typical values are in the interval $\\epsilon = (40 - 400)\\delta\\Omega$ MeV. Even though the emission solid angle % is likely to be as small as $\\Delta\\Omega \\sim 10^{-2} - 10^{-1}$ sterad, there % must be some concern here because the number of ICS pairs produced per primary % accelerated positron (Harding \\& Muslimov 2002) is not large so that either the % conversion of electron-positron energy to coherent radio emission is efficient % or other plasma components are involved, as in the paper by Cheng \\& Ruderman % (1980). A second set of observations that are relevant are those of polar-cap blackbody % X-ray luminosities. The blackbody X-ray emission expected from a polar-cap of % ion and proton zones can be found from equation (11). The reverse-electron flux % from photo-ionization heats an ion zone at a rate $H_{0}$ within a time interval % $0 < t < \\tau _{p}$. An estimate of this can be obtained directly from the mean % electron energy per unit nuclear charge accelerated, given by equation (29) of paper I, and is, \\begin{eqnarray} H_{0} = 6.0\\times 10^{18}Z^{0.85}(0)B^{1.5}_{12}T^{-1}_{6}P^{-1}, \\end{eqnarray} in units of erg cm$^{-2}$ s$^{-1}$. In this expression, $T$ is not the local % surface temperature but is an average for radiation emitted over the whole polar % cap. The surface temperature derived from equation (11) with neglect of % radiative loss is \\begin{eqnarray} T(t) = \\frac{2H_{0}}{(\\pi C\\lambda _{\\parallel})^{1/2}}t^{1/2}, \\hspace{5mm}0 < % t < \\tau _{p}, \\end{eqnarray} and increases very rapidly, so that the limiting temperature, \\begin{eqnarray} T_{06} \\approx 0.6Z^{0.17}(0)B^{0.3}_{12}P^{-0.2}, \\end{eqnarray} derived from $H_{0}$ is reached at $t \\ll \\tau _{p}$. The cooling at $t > \\tau % _{p}$ is also rapid so that the major part of the X-ray luminosity is that of a % black body at $T_{0}$ and of area approximately equal to the canonical % dipole-field area $2\\pi^{2}R^{3}/cP$ divided by $K + 1$. There have been many attempts to measure the polar-cap blackbody temperature and % source area of a subset of radio pulsars. We refer to Zavlin \\& Pavlov (2004) % for B0950+08 ; De Luca et al (2005) for B0656+14 and B1055-52; % Tepedelenlio\\u{g}lu \\& \\\"{O}gelman (2005) for B0628-28; Zhang, Sanwal \\& Pavlov % (2005) for B0943+10; Kargaltsev, Pavlov \\& Garmire (2006) for B1133+16; Gil et % al (2008) for B0834+06; Misanovic, Pavlov \\& Garmire (2008) for B1929+10. Five pulsars (B0628-28, B0834+06, B0943+10, B1133+16, B1929+10) have source % areas one or two orders of magnitude smaller than the canonical area, but unfortunately, a systematic comparison with equations (40)-(42) is not possible % because in most instances the authors are able to say only that the observed % X-ray spectrum is consistent with the stated temperature and source area. The % quoted source temperatures are $\\sim 3\\times 10^{6}$ K and are larger than those % predicted by equation (42) unless it is assumed that $B \\gg B_{d}$. They are % also uncomfortably large in the context of ${\\bf E}\\times {\\bf B}$ % drift-velocity polar-cap models such as that developed by Gil, Melikidze \\& % Geppert (2003) although it must be conceded that this model allows an % interesting test of its validity (see Gil, Melikidze \\& Zhang 2007). However, the case (iii) surface electric-field boundary condition on which these models rely could be maintained only for ion cohesive energies of % $\\sim 10$ keV which in turn would imply actual fields two orders of magnitude or % more larger than the ATNF catalogue dipole field. ", "conclusions": "This paper is a continuation of a previous study (Paper I) of isolated neutron % stars with positive corotational charge density and surface electric-field % boundary condition ${\\bf E}\\cdot{\\bf B} = 0$ at the polar cap. The reverse flux % of electrons arising from photo-ionization of accelerated ions is incident on % the neutron-star surface and produces protons through formation and decay of the % giant dipole resonance in the later stages of electromagnetic shower % development. Protons are the major component of the accelerated plasma but we % find that a time-independent composition of ions, protons and positrons is % usually unstable. The consequences of these phenomena should be observable in % pulsars unless, of course, rendered nugatory by some factor not properly taken % into account here. But we emphasize that there are almost certainly other % sources of instability contributing to the complex behaviour observed in radio % pulsars that may be present in each of cases (i) - (iii). There are two instabilities which result in transitions between states of % different plasma composition above localized areas on the polar cap and it is % suggested here that these are the basis for the commonly observed phenomena of % nulls and of subpulse drift. The basic units of time are $\\tau _{p} \\sim % 10^{-1}$ s for the short time-scale instability and for the medium time-scale, % $t_{rl} \\sim 10^{2} - 10^{3}$ s. These instabilities are not primarily % electromagnetic in origin and will not be present in neutron stars with negative % corotational charge density and electron acceleration. Micropulses of $\\sim % 10^{-4} - 10^{-3}$ s duration are also present in some pulsar emissions but are % unlikely to be associated with the instabilities we consider here. There appears to be no reason why both signs of corotational charge density % should not be present in the isolated neutron-star population and it would be % interesting to see if there is any observational evidence for an associated % division of the ATNF catalogue pulsars. The transitions are principally in plasma composition and the current density at % the neutron star surface remains close to $J=\\rho _{0}c$. The electric % potential $\\Phi(z,{\\bf u}(z))$ in the open magnetosphere above the polar cap is % also almost unchanged. These transitions are, in principle, observable because % there are no processes by which electron-positron pair creation can occur in a % plasma containing only protons, whereas it does occur in the ion plasma. Thus % there will be no coherent radio emission in the proton phase if we assume, as is % usual, that a pair plasma is a necessary constituent for it. Both the medium time-scale instability and the mechanical instability of the % neutron-star surface mentioned in Section 4 provide a basis for pulse nulls, but % unfortunately, not one that is suitable for quantitative prediction. The short % time-scale instability is more interesting. We show that at any instant, the % polar cap is divided into moving ion and proton zones and consider organized % motion of these, specifically the circular movement of compact ion zones around % the magnetic pole. We make the hypothesis that such motions can occur % (following recent work by Deshpande \\& Rankin 1999) but have not attempted to show that they have long-term stability, to the extent % indicated by observation, against decay to a chaotic state. This is of some % interest because it shows that features usually considered to require the % surface-field boundary condition ${\\bf E}\\cdot{\\bf B} \\neq 0$ are possible under % the ${\\bf E}\\cdot{\\bf B} = 0$ condition. Actual polar magnetic flux densities exceeding $10^{14}$ G are % necessary for the former condition and it must be doubtful that such fields, two % orders of magnitude greater than the inferred (catalogue) dipole field, are % likely to exist in the very large number of pulsars observed by Weltevrede et al % (2006) to show subpulse drift. Our view of subpulse drift is also interesting % in that the motion is not an ${\\bf E}\\times {\\bf B}$ drift velocity as in the % model of Ruderman \\& Sutherland. The fixed parameters for a given neutron star % are the band separation $P_{3}$, and $K$ which is the number of protons produced % per unit ion charge accelerated. This latter parameter approximately determines % the ratio of the total areas of proton and ion zones on the polar cap at any % instant. Thus the circulation time $\\hat{P}_{3} = nP_{3}$ depends on $n$, the % number of ion zones, and is not necessarily constant as it would be in the ${\\bf % E}\\times{\\bf B}$ drift velocity model. The band separation $P_{3}$ is dependent % on $K$ and on the diffusion time $\\tau _{p}$ and so is independent of rotation % period $P$ and almost independent of $B$. We have assumed here, following the pair formation calculations of Hibschman \\& % Arons (2001) and Harding \\& Muslimov (2002), that spontaneous pair creation by % curvature radiation is not possible in most isolated neutron stars. Those with % a sufficiently large value of $|\\Phi _{max}| \\propto B_{d}/P^{2}$ to support % spontaneous CR pair creation, predominantly young high-field pulsars, are not % expected to show the instabilities we have considered here. These authors % actually considered case (i) in which the basis of plasma formation is an % electron current density close to $J = \\rho _{0}c$ at the surface accelerated to % energies sufficient for ICS or spontaneous CR pair formation. In case (ii), the % protons and ions form the basic current component with $|\\Phi _{max}|$ broadly % the same as in case (i) apart from the ion inertia term which is very small for electrons. Thus the condition for the spontaneous growth of CR pair % formation is very similar in cases (i) and (ii). Although the reverse electron % flux arising from CR pair formation in case (ii) is probably quite a small % fraction of $\\rho _{0}c$, it is likely that, for the large values of $|\\Phi % _{max}|$ that are necessary, there will be a proton atmosphere which is never % exhausted. In this instance, steady-state plasma formation and low-altitude % acceleration appear probable, though subject to possible instabilities at those % higher altitudes where the coherent radiation is formed. The same statement can % be made about all case (i) pulsars whether or not spontaneous CR pair formation % is supported. Microstructure with $10^{-4} - 10^{-3}$ s timescales is likely to % be a consequence of such higher-altitude instabilities. The reason for this % assumption is that the case (i) boundary condition on $\\Phi$ needs only an % electron density on the surface separating open from closed magnetospheres. The small electron cohesive energy means that there is no reason why this should not be maintained at all times, as is also % true for the Goldreich-Julian electron current density. For this boundary % condition, there is no obvious way in which non-electromagnetic time-constants % can influence plasma formation and acceleration at altitudes less than $\\sim % 10^{6}$ cm above the polar cap. It is possible, perhaps, to assign either this % boundary condition or case (ii) with spontaneous CR pair formation to those pulsars in the systematic study of Weltevrede et al (2006), % approximately one half of the total in number, that do not exhibit subpulse % drift." }, "1101/1101.5912_arXiv.txt": { "abstract": "We present an analysis of observations made with the Arcminute Microkelvin Imager (AMI) and the Canada-France-Hawaii Telescope (CFHT) of six galaxy clusters in a redshift range of 0.16--0.41. The cluster gas is modelled using the Sunyaev--Zel'dovich (SZ) data provided by AMI, while the total mass is modelled using the lensing data from the CFHT. In this paper, we: i) find very good agreement between SZ measurements (assuming large-scale virialisation and a gas-fraction prior) and lensing measurements of the total cluster masses out to $r_{200}$; ii) perform the first multiple-component weak-lensing analysis of A115; iii) confirm the unusual separation between the gas and mass components in A1914; iv) jointly analyse the SZ and lensing data for the relaxed cluster A611, confirming our use of a simulation-derived mass-temperature relation for parameterizing measurements of the SZ effect. ", "introduction": "We employ two independent methods in a pilot study to investigate mass distributions of six galaxy clusters selected to cover a range of redshifts and merging states: \\begin{enumerate} \\item{Weak gravitational lensing (see e.g. \\citealt{2001PhR...340..291B} for a review), in which images of background objects are distorted by a mass lying along the line-of-sight, and can be used to directly probe the cluster mass distribution.} \\item{The Sunyaev--Zel'dovich effect (SZ; see e.g. \\citealt{SZ/Bir99} and \\citealt{carlstrom02} for reviews), the inverse-Compton scattering of the Cosmic Microwave Background (CMB) by the hot cluster gas, which effectively measures the gas pressure. The dark matter content is normally assessed by combining the SZ information with an X-ray measurement of gas temperature and using the assumption of hydrostatic equilibrium; but neither the X-ray temperature nor this assumption are necessary if we assume the cluster is virialised and we incorporate a sensible prior on the ratio of gas to dark matter \\citep{malak-param2010}.} \\end{enumerate} SZ and weak lensing have a natural complementarity, as they both have the potential to measure the distributed outskirts of clusters, with no strong bias towards concentrations of gas or mass. Also possible is the determination of the gas fraction of a given cluster by calculating the gas mass from the SZ and the total mass from lensing. It is expected that thousands of galaxy clusters will be detected by new SZ surveys performed by the Atacama Cosmology Telescope (ACT: \\citealt{ACT11}), the South Pole Telescope (SPT: \\citealt{SPTGalaxyClusters}) and \\textit{Planck} \\citep{PlanckMission}; the last has produced the first all-sky SZ catalogue, the first release of which is available in \\cite{PlanckESZ}. Upcoming large-area multi-wavelength optical surveys such as the the Dark Energy Survey \\citep{DES2005WP} and eventually the Large Synoptic Survey Telescope survey \\citep{LSSTOverview} will improve photometric redshift measurements of galaxies and release lensing-quality data over tens of thousands of square degrees of sky. Combining SZ and lensing measurements of very large samples of galaxy clusters may allow us to model their internal physics well enough for cosmological applications \\citep{HoekstraLensingReview}. The Arcminute Microkelvin Imager (AMI) is a radio interferometer that has made observations of hundreds of known galaxy clusters to measure their gas masses and structures via the SZ. A limitation of the large-area SZ survey instruments is their inability to resolve the morphology of cluster gas, so with these instruments it is difficult to examine the cluster dynamical state. Interferometric arrays such as AMI and the Combined Array for Research in Millimeter-wave Astronomy \\citep{CARMA} allow the examination of structures on scales between the high resolution of X-ray instruments and the somewhat lower resolution of the SZ survey instruments, as well as covering the northern sky. This study concerns a small selection of clusters of known X-ray structure, observed by AMI, for which publicly-available CFHT data were accessible at the time the AMI observations were made. As a pilot study for a future SZ-lensing comparison with a larger sample, we examine here six clusters with: a redshift range of 0.16--0.41, a range of masses and varying degrees of merging activity. Given the depth of the optical observations (see Section~\\ref{sec:GL}), clusters at higher redshift would be more difficult to observe as the field galaxy selection would likely be contaminated with foreground galaxies, and with only two optical bands we would be unable to easily reduce this contamination. Below a redshift of about 0.08, AMI starts to resolve out the cluster gas; this sample of clusters should not be affected by either of these issues, allowing us to examine the agreement of lensing and SZ mass measurements and the effect of the cluster dynamical states. AMI is limited to observing above Declination $20^{\\circ}$ and is effectively limited to measuring clusters of internal gas temperatures $\\gtrapprox2$\\,keV. Clusters of lower temperature may be detectable but would require very high (unphysical) electron densities to produce detectable SZ flux. A data reduction pipeline was developed to extract a weak-lensing catalogue from background galaxies in the CFHT fields. This allowed the measurement of the total matter distributions of the galaxy clusters, which could then be directly compared with gas measurements from the SZ. The SZ observations and data reduction are described in Section~\\ref{sec:SZ}, and weak lensing in Section~\\ref{sec:GL}. The Bayesian analysis, including sampling parameters and priors, is described in Section~\\ref{sec:modelling}. We present notes on each cluster and its available data in Section~\\ref{sec:results}, discuss the ramifications of these results in Section~\\ref{sec:discussion} and outline our conclusions in Section~\\ref{sec:conclusions}. Throughout, we assume a `concordance' $\\Lambda$CDM cosmology, with $\\Omega_{\\mathrm{m,0}}=0.3$, $\\Omega_{\\Lambda,0}=0.7$, and $H_{0}=70$km\\,s$^{-1}$Mpc$^{-1}$ (and thus $h=H_{0}/100=0.7$). All coordinates are J2000 epoch and all optical magnitudes use the AB system. ", "conclusions": "\\label{sec:conclusions} \\begin{enumerate} \\item{Using AMI SZ data that measure out to $r_{200}$, and thus probe scales comparable to those of weak-lensing observations, we have used a fast, Bayesian analysis to produce posteriors for useful parameters for six clusters, on the assumptions of an isothermal $\\beta$-model to SZ data and an NFW model for the lensing data;} \\item{Of the four clusters for which we have both weak lensing and SZ data, we find that the mass estimates are in very good agreement, and that we may need to improve our field-galaxy selection for clusters at $z>0.4$ for a larger sample;} \\item{We perform the first multiple-component weak-lensing analysis of A115 and discover significant substructure;} \\item{We confirm the unusual separation between the gas and mass components in A1914;} \\item{For A611, the most relaxed cluster in the sample, we have carried out a joint weak-lensing and SZ analysis with the gas fraction as a free parameter, and find that $f_{\\mathrm{gas}}=0.23\\pm0.10h^{-1}$.} \\end{enumerate}" }, "1101/1101.0476_arXiv.txt": { "abstract": " ", "introduction": "\\label{sec:intro} Type Ia supernova (SNIa) observations imply an acceleration of the cosmic expansion if general relativity is valid on cosmological scales and if the universe is homogeneous and isotropic on scales larger than 200Mpc. If we abandon one of these assumptions, however, other explanations become possible. One approach of such a nature is the so-called modified gravity that abandons general relativity. The other is the local void model, which was first proposed by Tomita \\cite{2000ApJ...529...26T}\\cite{2000ApJ...529...38T}, Goodwin et al. \\cite{1999astro.ph..6187G}, and Celerier \\cite{2000A&A...353...63C}, independently. This model abandon the second assumption, which is often called the Cosmological Principle or the Copernican Principle, and assumes that we are around the center of a low density spherically symmetric void and that the spacetime is well described by the Lema\\^{i}tre-Tolman-Bondi (LTB) model \\cite{1933ASSB...53...51L}\\cite{1934PNAS...20..169T}\\cite{1947MNRAS.107..410B}. In this model, the cosmic expansion rate decreases outward on each constant time slice, which produces an apparent acceleration of the universe when observed along the past light cone. Although this model violates the Cosmological Principle and requires an accidental situation concerning our location in the universe, it does not require any dark energy or a modification of gravity theory. Further, as far as the redshift-luminosity distance relation obtained by the SNIa observations is concerned, this model can reproduce the observational results with any accuracy because it contains at least one arbitrary function of the radius (see, e.g., \\citen{2008PThPh.120..937Y}). Actually, it has passed all observational tests so far. Therefore, it is of crucial importance to find observational tests that enable us to discriminate this void model from spatially homogeneous models employing dark energy or modified gravity theory, in order to establish the necessity of dark energy or a modification of gravity. There have been proposed various tests for that purpose so far\\cite{Tomita.K2009,Moss.A&Zibin&Scott2010}: CMB anisotropies on large angular scales, radial BAO\\cite{2008JCAP...04..003G}, spectral distortions of CMB\\cite{2008PhRvL.100s1302C}, the kinematic SZ effect\\cite{2008JCAP...09..016G}, and estimates of the Hubble rate \\cite{2010MNRAS.405.2231F}. Among these, the oldest and simplest one is to observe the effect of the inhomogeneity on CMB temperature. This effect was first estimated by Alnes and Amarzguioui\\cite{AA2006} for a special class of void models\\cite{AAG2006}. Recently, this analysis was extended to a wider class of models by Kodama, Saito and Ishibashi with the helps of analytic formulas for the dipole and quadrupole moments of the CMB temperature anisotropy for an off-center observer in a general spherically symmetric universe \\cite{2010PThPh.124..163K}. Although this type of test provides a strong constraint on the allowed range of our distance from the void center, it is not so decisive because lower moments including the dipole and quadrupole are significantly affected by the cosmic variance. One possible way to circumvent this weakness is to extend the analysis of the off-center CMB anisotropies to polarisation. As is well known, in the spatially homogeneous cosmology, inhomogeneities of the matter (galaxy) distribution modify the sky pattern of CMB polarisation by gravitational lensing and produce B-modes from E-modes\\cite{1996ApJ...463....1S}\\cite{1998PhRvD..58b3003Z}\\cite{2000PhRvD..62d3007H}\\cite{2003PhRvD..67h3002O}\\cite{2006PhR...429....1L}. In the local void model, the central observer detects no such effect because of the spherical symmetry, in spite of the strong inhomogeneity of the model. However, an off-center observer can detect the gravitational shear field through the observation of B-modes. Therefore, in this paper, we calculate the gravitational lensing effect on the CMB temperature and polarisation anisotropies for an off-center observer in the local void model and estimate the correlations among the temperature, the E-mode and the B-mode anisotropies. The paper is organised as follows. First, in the next section, we review the basic matters on CMB polarisation that are relevant to the present paper. Then, in Section \\ref{sec:gl}, we perturbatively solve the geodesic equations in the LTB model to find the change in the propagation direction due to the gravitational lensing effect for null rays close to the central past light cone. This determines the shift vector on the sky for an off-center observer that represents the difference between the observed direction of a light ray and the corresponding direction of the last scattering point in the fiducial spatially homogeneous model in a gauge in which the fiducial universe has the same spherical symmetry as the LTB background. Next, in Section \\ref{sec:polLTB}, we carefully examine how the polarisation evolves from the last scattering surface to the present time in the LTB spacetime with the help of the collisionless Boltzmann equation. In particular, we show that the temperature and polarisation anisotropy of CMB in the LTB model can be expressed by the same formula in terms of the shift vector as in the FLRW universe background. Putting these results together, in Section \\ref{sec:result}, we give general formulas for the CMB temperature and polarisation anisotropy produced by the void inhomogeneity and for the correlations among the harmonic components of the temperature, the E-mode polarisation and the B-mode polarisation. Finally, Section \\ref{sec:sum} is devoted to summary and discussions. ", "conclusions": "\\label{sec:sum} In this paper, we have developed a formulation to calculate the gravitational lensing effects on the CMB temperature and polarisation for an observer close to the center of a spherically symmetric void described by the LTB model. In particular, we have derived explicit expressions for the correlations among the anisotropies of temperature and polarisation induced by gravitational lensing in terms of an integration of known geometrical quantities along the central past light cone. With the helps of these formulas, we have found that for an off-center observer in the local void, there appear nonzero correlations between $T$ and $B$ and between $E$ and $B$ that are diagonal in the harmonic coefficient expression in the leading order with respect to the observer offset distance. Similar correlations arise if there exists a quintessence-type axionic field with mass in the range from $10^{-29}\\r{eV}$ to $10^{-33}\\r{eV}$, but in this case the magnitudes of correlations have different dependence on $\\ell$. Hence, if the correlations suggested by our result are detected in future B-mode measurement experiments with high precision, they would provide a clear signal showing the existence of a local void. We have also given some preliminary numerical estimations of the gravitational lensing effect. The results indicate that the B-mode amplitude produced by the lensing is around $10^{-3}$ times that of E-modes if we take into account the constraint on the off-center distance on the observer from the dipole anisotropy of the CMB temperature\\cite{2010PThPh.124..163K}. This is because the effect is proportional to the distance of the observer to the symmetry center. Thus, we need next generation B-mode experiments to use this effect to test the viability of the local void model. However, the results also show that the gravitational lensing effect is very sensitive to the void profile and sizes. Hence, when it is detected, it is useful to specify a model. Because the lensing effect becomes larger for smaller voids in general, it may be also used to detect anomalously large voids, which are much smaller than the standard void size ($\\sim$ 1Gpc) in the local void model but still statistically rare in the $\\Lambda$CDM model. In the present paper, we have only considered the local void model, but the method developed here can be applied to more general inhomogeneous models, such as realistic universe models with voids outside us. Such extensions are under investigation." }, "1101/1101.4639_arXiv.txt": { "abstract": " ", "introduction": "\\label{sec:intro} Over 20 years ago,~\\citet{weid89} asked, ``Do asteroids have satellites?'' Today, binary systems have been discovered in every dynamical class of small Solar System bodies from near-Earth asteroids to Mars-crossers and main-belt asteroids, among the Jupiter Trojans, and in the outer Solar System in the Kuiper belt. Beginning with the Galileo spacecraft's serendipitous discovery in 1993 of tiny satellite Dactyl orbiting (243) Ida~\\citep{chap95,belt96} while on its cruise to Jupiter and continuing with the success of radar, lightcurve photometry, ground-based adaptive-optics imaging, and Hubble Space Telescope imaging, reviewed by~\\citet{merl02},~\\citet{rich06}, and~\\citet{noll08}, over 180 small Solar System bodies are suspected to be binary or multiple systems.\\footnote{The value of more than 180 suspected binary or multiple systems is taken from http://www.johnstonsarchive.net/astro/asteroidmoons.html and based on references therein. Lists of binary and multiple asteroid parameters are available from the Ondrejov Asteroid Photometry Project (http://www.asu.cas.cz/$\\sim$asteroid/binastdata.htm) and the Planetary Data System (http://sbn.psi.edu/pds/asteroid/EAR$\\_$A$\\_$COMPIL$\\_$5$\\_$BINMP$\\_$V3$\\_$0/data/).} In this work, we build upon the ideas presented in~\\citet{weid89} to illustrate the tidal end states of binary asteroid systems and discuss the inherent material strength of the bodies implied by tidal evolution. The Keplerian orbit of the two components in a binary system about the center of mass allows one to determine the mass of the system and, with an estimate of the component sizes, the (assumed equal) densities of the bodies. Density estimates combined with taxonomic classifications hint at the internal structure of asteroids~\\citep{brit02}. When compared to analog meteorites with known bulk densities, low densities found for rocky asteroids imply that, rather than monolithic bodies, asteroids are aggregate bodies made up of many fragments with varying degrees of porosity. Aggregates can range from low-porosity fractured or shattered bodies to porous, cohesionless rubble piles with ``gravitational aggregate'' acting as a catch-all term for bodies comprised of many fragments, independent of porosity, that are held together by their collective gravity~\\citep{rich02}. The tidal evolution of these binary systems is intimately tied to the internal structure and material strength of the bodies involved, and we will exploit this dependence to estimate the combined effect of rigidity and energy dissipation in the bodies to determine if asteroids tend to have solid, fractured, or shattered interiors. Hints of the internal structure of asteroids also come from the most probable binary formation mechanisms for each population: near-Earth binaries likely form through rotational disruption via spin-up by thermal torques due to sunlight~\\citep[the YORP effect;][]{bott02,wals08yorp} or close planetary encounters~\\citep{rich98,wals06}, 1- to 10-km-scale binaries in the main belt may also form through YORP spin-up~\\citep{prav07}, 100-km-scale main-belt binaries likely form through sub-catastrophic impacts~\\citep[\\textit{e.g.,}][]{durd04}, and Kuiper-belt binaries may have formed very early on via gravitational collapse~\\citep{nesv10} or later on via collisions or a flavor of dynamical capture~\\citep[see][for a review]{noll08}. Binary formation through rotational disruption at the very least implies a fractured interior, and if satellite formation is due to the accretion of smaller particles spun off the parent body, as modeled by~\\citet{wals08yorp}, part or all of the parent body is likely an aggregate of smaller pieces and structurally weak. For sub-catastrophic impacts, the satellite is either a solid shard or an aggregate made up of impact ejecta while the parent body may remain solid or fractured rather than shattered. Thus, we anticipate that examination of tidal evolution in binary asteroid systems will find that large main-belt binaries are consistent with solid or fractured rock, while near-Earth and small main-belt binaries are consistent with structurally weaker, shattered or porous gravitational aggregates. Section~\\ref{sec:angmom} introduces the important dynamical quantities of binary systems, angular momentum and energy. In Section~\\ref{sec:synch}, we derive and illustrate for a range of angular momenta the fully despun, double synchronous orbits that are end states of tidal evolution and then discuss stability limits and energy regimes in Section~\\ref{sec:stablimit}. Using the known properties of binary asteroids in the near-Earth region and the main belt, in Section~\\ref{sec:material} we constrain the material properties of the bodies based on tidal evolution. ", "conclusions": "\\label{sec:disc} We have derived and plotted the locations of the fully synchronous tidal end states for all mass ratios and angular momenta and find that only binary asteroid systems with nearly equal-mass components such as (90) Antiope and (69230) Hermes reside in these end states with all other systems with smaller mass ratios undergoing a lengthy tidal evolution process. Though (90) Antiope and (69230) Hermes are excellent examples of the stable fully synchronous tidal end state, these systems can evolve to their current states so rapidly for reasonable values of the product of rigidity $\\mu$ and tidal dissipation function $Q$ that their material properties are difficult to constrain. Relying instead upon the binary systems with smaller mass ratios, we find among 100-km-scale main-belt binaries values of the material strength $\\mu Q$ that, for tidal evolution over the age of the Solar System, are consistent with solid or fractured rock as one might expect for binaries created via sub-catastrophic collisions. Binaries formed in the near-Earth region via spin-up that have evolved over the typical 10 Myr dynamical lifetime require material strengths orders of magnitude smaller than their large main-belt counterparts as one might expect for gravitational aggregates of material compared to essentially monolithic bodies. If near-Earth binaries formed in the main belt prior to injection to the near-Earth region, the time available for tidal evolution can lengthen and $\\mu Q$ can increase in a similar manner. The discovery of many small main-belt binaries with rapidly spinning primaries and angular momenta similar to the near-Earth binaries lends some credence to the idea that near-Earth binaries, or some fraction of near-Earth binaries, could form in the main belt prior to delivery to the near-Earth region and have binary ages that are limited by collisional lifetimes rather than the dynamical lifetime in the near-Earth region. An older age for some near-Earth binaries, \\textit{i.e.}, 100 Myr versus 10 Myr, would help reconcile some especially small values of $\\mu Q$ with the model of gravitational-aggregate structure of~\\citet{gold09}. If small main-belt binaries are up to 1 Gyr old, the maximum collisional lifetime of their secondaries, their material strengths are more consistent with solid rock than gravitational aggregates. Younger ages similar to the near-Earth binaries would reduce $\\mu Q$, but leave an older binary population unaccounted for either because of detection bias, collisional destruction, or another destruction mechanism. Smaller $\\mu Q$ values for the small main-belt binaries would be more consistent with the gravitational-aggregate structure implied by the likely formation mechanism of YORP spin-up and would allow tidal spin-down to overcome continued YORP spin-up of the primary, which could account for the slightly slower rotation rates of small main-belt primaries compared to near-Earth primaries. Beyond the question of age, our ignorance of the inherent $Q$ of gravitational aggregates and whether they can be described in such idealized terms are significant caveats in this approach to understanding the material properties of asteroids. Especially in the near-Earth region, the assumption that tides are the dominant method of mutual-orbit expansion may also be tenuous. The binary YORP effect, mass lofting, and close planetary encounters could conspire to also expand or contract the mutual orbits, either assisting or hindering the expansion of the mutual orbit by tides. A combination of effects altering the semimajor axis would skew the resulting $\\mu Q$ when considering tides alone; multiple methods of orbital expansion working in concert would give the illusion of a small $\\mu Q$ value rather than the true material strength of the body. It is difficult to disentangle different orbit-expansion mechanisms on top of not knowing the precise age of a binary system. If we take tides to be the dominant orbit-expansion mechanism, then the assumption of the binary age $\\Delta t$ is easily the largest source of error in the calculation of $\\mu Q$ for a binary system causing an order-of-magnitude error in $\\mu Q$ for an order-of-magnitude error in $\\Delta t$. Other sources of error, those summarized in Section~\\ref{sec:other} and also from ignoring tides raised on the secondary, will typically cause errors of less than a factor of two. In our examination of tidal evolution, a few observations have hinted at the presence of a mechanism (or mechanisms) in addition to tides that evolves the semimajor axis of binary asteroid systems. Following~\\citet{gold09}, using their model for gravitational-aggregate structure and assuming the actual strength of the BYORP torque on the mutual orbit is $10^{-3}$ times its maximum possible value,\\footnote{A similar factor for the YORP torque was found to be $4\\,\\times\\,10^{-4}$ for asteroid (54509) YORP (formerly 2000 PH$_{5}$) by~\\citet{tayl07} and~\\citet{lowr07}.} it is found that BYORP should dominate the semimajor-axis evolution of all near-Earth binary systems with synchronous secondaries. For small main-belt binaries with unequal-mass components, BYORP and tides could contribute comparably to the semimajor-axis evolution, and even for large main-belt binaries, BYORP could be important under certain circumstances. Among the near-Earth binaries, the systems most susceptible to BYORP at some point in their evolution are 2004 DC and 2003 YT$_{1}$, precisely the systems whose asynchronous secondaries, eccentric mutual orbits, and small and worrisome $\\mu Q$ values led us to suspect a BYORP component to their evolution in Section~\\ref{sec:nea}. Large main-belt binary systems with mass ratios above $\\sim$0.01 are dominated by tides at their current separations. Large main-belt binaries with the smallest size ratios are more susceptible to BYORP, and we find that those that plot furthest to the right of the 4.5 Gyr curve in Fig.~\\ref{fig:mba} and, hence, have the smallest calculated $\\mu Q$ values: (45) Eugenia, (130) Elektra, and (702) Alauda, are most likely to be dominated by the BYORP effect at this time, which would lead to artificially low $\\mu Q$ values compared to the rest of the binaries in Table~\\ref{tab:mba}. In the past, at smaller separations, these low mass-ratio systems among the large main-belt binaries would have been tide-dominated, which would allow for synchronization of the secondary with a rough transition to BYORP-dominated semimajor-axis evolution occurring around 6$R_{\\rm p}$ if the strength of the BYORP effect is correctly estimated. Subsequent expansion by BYORP and tides would then produce the currently observed systems where assuming tide-dominated evolution finds a smaller value for $\\mu Q$ than is necessarily true. This fact is related to point \\#3 in Section~\\ref{sec:other}: the calculation of $\\mu Q$ is most sensitive to the evolution within 10\\% of the current separation such that if BYORP is a significant contributor to the semimajor-axis expansion in this range, then $\\mu Q$ will be skewed to a lower value than is appropriate for the material. If the strength of the BYORP torque is $10^{-5}$ times its maximum possible value instead, the majority of near-Earth binaries and the three large main-belt systems mentioned would remain BYORP-dominated in their current configurations. The small main-belt binaries would be tide-dominated currently making it more difficult for destruction by BYORP to account for the apparent lack of older systems. If the strength of the BYORP torque is reduced even further to $10^{-7}$ times its maximum possible value, then only 2004 DC and 2003 YT$_{1}$ (and possibly 2000 DP$_{107}$) would remain BYORP-dominated. Given that BYORP could have a significant impact on many of the binary systems we have considered, perhaps the best example for the calculation of $\\mu Q$ is (22) Kalliope, which has a small enough mass ratio that the system is still tidally evolving, unlike (90) Antiope or (69230) Hermes, yet not so small that BYORP could be an issue. Given the slew of sources of uncertainty in the calculation of material properties via tidal evolution, Kalliope's $\\mu Q$ should lie within a factor of a few of $3.3 \\times 10^{12}$ N\\,m$^{-2}$ based on a 4.5 Gyr evolution, which implies stronger material properties than Phobos, but is reasonable for a monolithic or fractured asteroid. For a tide-dominated binary system, assuming an age equal to the age of the Solar System provides a confident order-of-magnitude upper bound on the value of $\\mu Q$. One way to improve on this upper bound is to limit the evolution time by using collisional lifetimes (as discussed for near-Earth and small main-belt binaries) or by finding binaries in collisional families. If we assume the binary was formed as the result of the catastrophic disruption of a parent body that produced a family of asteroids (\\textit{i.e.}, EEBs, escaping ejecta binaries, as described by~\\citet{durd04}, then the binary evolves only over the age of the family. Changing the evolution time by some factor will affect the value of $\\mu Q$ in the same manner. Of the main-belt binaries listed in Tables~\\ref{tab:mba} and~\\ref{tab:smba}, about a dozen (19\\%) are believed to be part of the collisional families determined by~\\citet{zapp95}, smaller than the value of roughly 30\\% found among all asteroids~\\citep{marz99,bend02}. Those families the binaries are associated with that have age estimates, $\\textit{i.e.}$, Flora~\\citep{nesv02flora}, Themis~\\citep{marz95,marz99,nesv05}, and Eos~\\citep{vokr06eos}, are believed to be of order 1 Gyr old so that limiting the ages of the binaries to the age of the families rather than the age of the Solar System does not change the upper bound of $\\mu Q$ by more than a factor of a few. If one could instead identify binary systems among much younger groups, such as the Karin or Veritas clusters, both of which are younger than 10 Myr~\\citep{nesv02karin,nesv03}, the ages of the binaries, and thus their $\\mu Q$ values, would be far better constrained." }, "1101/1101.0256_arXiv.txt": { "abstract": "{Most of the X-ray background (XRB) is generated by discrete X-ray sources. It is likely that still unresolved fraction of the XRB is composed from a population of the weak sources below the present detection thresholds and a truly diffuse component. It is a matter of discussion a nature of these weak sources.} {The goal is to explore the effectiveness of the nearest neighbor statistics (NNST) of the photon distribution for the investigation of the number counts of the very weak sources.} {All the sources generating at least two counts each induce a nonrandom distribution of counts. This distribution is analyzed by means of the NNST. Using the basic probability equations, the relationships between the source number counts $N(S)$ and the NNST are derived.} {It is shown that the method yields constraints on the $N(S)$ relationship below the regular discrete source detection threshold. The NNST was applied to the medium deep \\chandra\\ pointing to assess the source counts $N(S)$ at flux levels attainable only with the very deep exposures. The results are in good agreement with the direct source counts based in the \\chandra\\ Deep Fields (CDF).} {In the next paper of this series the NNST will be applied to the the CDF to assess the source counts below the present flux limits.} ", "introduction": "The X-ray background (XRB) is mostly generated by discrete extragalactic sources (e.g. \\citealt{lehmann01, kim07}, and references therein). Thus, the source counts provide the essential information on the constituents of the XRB and the X-ray $N(S)$ relationship has been a subject of numerous studies for the last $40$ years. The individual point-like source is detected if a number of counts within a specified area exceeds the assumed threshold. Size of the detection box is defined by the Point Spread Function (PSF), while the detection threshold is usually selected to minimize number of false detections and at the same time to maximize number of the real sources. The detection threshold is typically set at the level of $4-5\\, \\sigma$ above the local average count density. A presence of weaker sources, below the formal detection threshold, is manifested by the increased fluctuations of the count distribution as compared to the fluctuations expected for the random counts. A common approach to assess counts of sources weaker than the detection limit is based on the count density fluctuation analysis. To quantify signal generated by the discrete sources one should determine the intensity distribution $P(D)$, i.e. the histogram of the number of pixels as a function of the number of counts. The observed function $P(D)$ is then compared with the functions obtained from the simulated count distributions \\citep[e.g.][]{hasinger93,miyaji02}. It is assumed that the simulated source counts represent the actual source distribution if the model $P(D)$ function mimics the observed histogram. A contribution of point sources to the count distribution one can estimate also using the auto-correlation function (ACF). Since the integral of the ACF is directly related to the second moment of the $P(D)$ distribution, \\citep[e.g.][]{soltan91} both methods are closely related. An innovative method to assess the number of weak sources was proposed by \\citet{georgakakis08}. In their approach the count distribution in the detection cell is explicitly expressed as a sum of the source and background counts. As a result of a rigorous application of the Poisson statistics, a flux probability distribution is derived as a function of the total and background counts observed in the detection cell. This probability distribution combined with the adequately defined sensitivity map of a given observation is then used to estimate the the source number counts. The count fluctuations are proportional to the source intensities. Thus, the observed fluctuation amplitude is dominated by the sources just below the detection threshold set for the individual objects, whereas it is only weakly sensitive to the fainter sources which produce smaller number of counts. One should note, however, that every source which produces more than one count generates deviation from the random count distribution. In the present paper we investigate the efficiency of the nearest neighbor statistics (NNST) for the weak source analysis and we show that the NNST is a powerful tool to estimate the source counts, $N(S)$, down to very low flux levels. We apply this technique to one of the \\chandra\\ AEGIS fields with the exposure of $465$\\,ks. The NNST allows us to obtain the $N(S)$ relationship extending down to $2\\times 10^{-17}$\\,erg\\,cm$^{-2}$s$^{-1}$ and $7\\times 10^{-17}$\\,erg\\,cm$^{-2}$s$^{-1}$in the $0.5-2$\\,keV and $2-8$\\,keV energy bands, respectively, i.e. a factor of $5-10$ below the standard sensitivity threshold corresponding to this exposure. Since the present count estimates are contained within the flux range covered by the direct source counts derived from the deepest \\chandra\\ fields (such as CDFS), the effectiveness of the NNST method could be directly assessed. The organization of the paper is as follows. In the next section, the method and all the relevant formulae are presented. Then, in Sect.~\\ref{observations}, the observational material is described and the computational details including questions related to the PSF are given. Results of the calculations, i.e. estimates of the source counts below the nominal sensitivity limit are presented in Sect.~4. The results are summarized and discussed in Sect.~5. Prospects for the application of the NNST to the deep \\chandra\\ fields are presented. ", "conclusions": "We have estimated the source number counts in the S band down to $\\sim\\!2\\cdot 10^{-17}$\\,cgs using the merged data with the integrated exposure time of $\\sim\\!465$\\,ks. This flux level is below the standard detection threshold for individual sources in the deepest \\chandra\\ exposures of $2$\\,Ms \\citep{kim07}. Our slope estimate below $S = 10^{-16}$\\,cgs fits perfectly the actual discrete source counts determined using such deep observations. It shows the NNST potential as an effective tool in the investigation of the extremely weak source population. In the second paper of this series we plan to apply the NNST method to the \\chandra\\ Deep Fields. With a $2$\\,Ms exposure the NNST will allow to assess number counts down to $\\sim\\!4\\cdot 10^{-18}$\\,cgs in the $0.5-2$\\,keV band and to $\\sim\\!2\\cdot 10^{-17}$\\,cgs in the $2-8$\\,keV band, although in the latter case the expected accuracy of our estimate might be quite low. The constraints obtained for the H band are not restrictive. This is because the contribution of the non X-ray counts increases with energy and the data become strongly contaminated by the particle background which effectively ``dilutes'' the counts concentrations produced by the sources. In the next paper some prospects to improve the S/N ratio above $2$\\,keV will be explored." }, "1101/1101.0898_arXiv.txt": { "abstract": "We detected recent star formation in nearby early-type galaxies located in low density environments, with {\\it GALEX} Ultraviolet (UV) imaging. Signatures of star formation may be present in the nucleus and in outer rings/arm like structures. Our study suggests that such star formation may be induced by different triggering mechanisms, such as the inner secular evolution driven by bars, and minor accretion phenomena. We investigate the nature of the (FUV-NUV) color vs. Mg$_2$ correlation, and suggest that it relates to ``downsizing'' in galaxy formation. ", "introduction": "Although early-type galaxies (ETGs hereafter) are considered the fossil record of the process of galaxy evolution, there is growing evidence that ``rejuvenation'' episodes may occur in their star formation history. In this context, the halo mass of ETGs is the main driver for their evolution, however also the environment plays a significant role \\citep{Clemens06,Clemens09}. According to these studies, the galactic nuclei (SDSS fibers sample the central 3\\arcsec) of ETGs in low density environments (LDE hereafter) are about 20\\% younger than those of their cluster counterparts. Wide-field (1.25 degrees diameter), deep far-UV (1344 $-$ 1786 \\AA) and near UV (1771 $-$ 2831 \\AA) imaging from {\\it GALEX} (see for details Bianchi these proceedings) is greatly contributing to this view. Using {\\it GALEX} data, e.g. \\citet{Schawinski07} found that 30\\% of massive ETGs show ongoing star formation, and that this fraction is higher in LDE than in the high-density environments. \\citet{Rampazzo07} and \\citet{Marino09} showed that ETGs with shell structures (indicative, according to simulations, of recent accretion episodes) host a ``rejuvenated'' nucleus \\citep[see also][for optical spectroscopic studies]{Longhetti00}. Similar results have been obtained by \\cite{Jeong09} for the {\\tt SAURON} galaxy sample \\citep{deZeeuw02}. We are performing a comprehensive, multi wavelength study of 65 nearby ETGs, a large fraction of which show ionized gas emission, predominantly located in LDE \\citep[][A07 and A10 hereafter]{Annibali07,Annibali10}. The sample is composed of 70\\% of elliptical and 30\\% of S0s \\citep[see both][]{RC3,RSA}. Anyway, from the kinematic point of view, about 68\\% of our ETGs have fast rotator characteristics in the $\\epsilon$ vs. $V/\\sigma_e$ plane (see Appendix A in A10). Here we present the GALEX view for a sub-sample of 40 ETGs, out of the 65 in the original sample, available in the NASA archive. The UV spectral region is sensitive to even small amounts of recent star formation, and is thus effective in unveiling possible ``rejuvenation'' episodes. For a detailed presentation of the GALEX FUV and NUV photometry we refer to \\citet{Marino10}. We review our multi-wavelength approach in quest of signatures of recent star formation in ETGs. Finally we combine the analysis of the UV photometry with optical line-strength indices and show preliminary results. \\begin{figure}% \\includegraphics[width=6.7cm]{f1a.ps} \\includegraphics[width=6.7cm]{f1b.ps} \\includegraphics[width=6.7cm]{f1c.ps} \\caption{GALEX false color image (FUV blue; NUV yellow) of NGC 3258, NGC 5813 and of the Seyfert~2 IC~5063 (adapted from \\citet{Marino10}). The ellipse marks the isophote at $\\mu_B=25$~mag~arcsec$^{-2}$ (D$_{25}$). Optical line-strength indices indicate that NGC~3258 has a luminosity weighted age of 4.5$\\pm$0.8 Gyr, while NGC 5813 is an old ETG with an age of 11.7$\\pm$1.6 \\citep[see][]{Annibali07}} % \\label{fig1} \\end{figure} \\begin{figure}% \\begin{center} \\includegraphics[width=6.2cm]{f2a.eps} \\includegraphics[width=6.0cm]{f2b.eps} \\includegraphics[width=6.0cm]{f2c.eps} \\caption{Optical \\citep[see details in][]{Rampazzo05,Annibali06} and MIR spectra of NGC~3258 and NGC~5813 galaxies imaged in Figure~1. The top panel shows the nuclear spectra of the two galaxies, fluxes are shifted arbitrarily, in the range 3750 $\\leq \\lambda \\leq$ 7250 \\AA. In the mid panel we show the position of the two galaxies in the typical BTP diagram (see for details A10) which indicates the LINER/Composite nature of the two galaxy nuclei. The bottom panel shows the {\\it Spitzer}-IRS MIR spectra. Notice the presence of PAHs in the spectrum of NGC 3258 at odds with NGC 5813. A cold dust component is also visible in the NGC 3258 spectrum } % \\end{center} \\label{fig2} \\end{figure} \\begin{figure}% \\includegraphics[width=\\columnwidth]{f3a.ps} \\includegraphics[width=\\columnwidth]{f3b.ps} \\caption{ B-band (top panel) and {\\it GALEX} FUV images (bottom panel) of MCG-05-07-001, a southern polar ring galaxy. In the FUV band (the square encloses the field of view of the B-band image), only the ring and the nucleus of the galaxy are visible: the old main body of the galaxy disappears \\citep[see][for details]{Marino09}} % \\label{fig3} \\end{figure} \\begin{figure}% \\includegraphics[width=\\columnwidth]{f4.eps} \\caption{ The {\\it GALEX} FUV emission (red areas) is superposed to the residual of the optical image of NGC~1210 after the main body of the galaxy has been subtacted. The residuals show a wide system of shells and a complex central fine structure. A ring of HI emission, which perfectly superposes to the {\\it GALEX} FUV, has been detected by \\citet{Schminovic01} \\citep[see][for details]{Marino09} } % \\label{fig4} \\end{figure} ", "conclusions": "" }, "1101/1101.3549_arXiv.txt": { "abstract": "We present a statistical parallax analysis of low--mass dwarfs from the Sloan Digital Sky Survey (SDSS). We calculate absolute $r$-band magnitudes ($M_r$) as a function of color and spectral type, and investigate changes in $M_r$ with location in the Milky Way. We find that magnetically active M dwarfs are intrinsically brighter in $M_r$ than their inactive counterparts at the same color or spectral type. Metallicity, as traced by the proxy $\\zeta$, also affects $M_r$, with metal poor stars having fainter absolute magnitudes than higher metallicity M dwarfs at the same color or spectral type. Additionally, we measure the velocity ellipsoid and solar reflex motion for each subsample of M dwarfs. We find good agreement between our measured solar peculiar motion and previous results for similar populations, as well as some evidence for differing motions of early and late M type populations in $U$ and $W$ velocities that cannot be attributed to asymmetric drift. The reflex solar motion and the velocity dispersions both show that younger populations, as traced by magnetic activity and location near the Galactic plane, have experienced less dynamical heating. We introduce a new parameter, the independent position altitude (IPA), to investigate populations as a function of vertical height from the Galactic plane. M dwarfs at all types exhibit an increase in velocity dispersion when analyzed in comparable IPA subgroups. ", "introduction": "Low--mass dwarfs (0.08 $\\msun$ $< \\cal{M} <$ 0.8 $\\msun$) dominate the stellar content of the Milky Way \\citep{2010AJ....139.2679B}. However, the study of these stars in large numbers was only recently realized with the advent of surveys such as the Sloan Digital Sky Survey \\citep[SDSS;][]{2000AJ....120.1579Y} and the Two Micron All Sky Survey \\citep[2MASS;][]{2006AJ....131.1163S}. These surveys, with their precise, multi--band photometry and (in the case of SDSS) spectroscopic coverage have led to observational catalogs of unprecedented size. Currently, the largest spectroscopic database of M dwarfs \\citep[][hereafter Paper I]{west10} has produced spectral types, radial velocities (RVs) and chromospheric activity estimates (as traced by Balmer series emission) for over 70,000 M dwarfs. SDSS photometry of millions of stars has been used to measure the field luminosity and mass functions of M dwarfs \\citep{covey08,2010AJ....139.2679B}, the structure of the Milky Way's thin and thick disks \\citep{2008ApJ...673..864J,2010AJ....139.2679B}, and the properties of wide, common proper motion binaries \\citep{2010AJ....139.2566D}. The ubiquity and long main sequence lifetimes of low--mass stars \\citep{1997ApJ...482..420L} make M dwarfs ideal tracers of nearby Galactic structure and kinematics \\citep[e.g.,][]{ 1993ApJ...409..635R,1995AJ....110.1838R,2007AJ....134.2418B, 2009AJ....137.4149F, 2010AJ....139.2679B}. Despite the recent advances enabling the study of these stars in a broad Galactic context, one fundamental parameter remains elusive: distance. Accurate trigonometric parallaxes, and the resulting absolute magnitude and distance measurements, remain difficult to obtain due to the intrinsic faintness ($L \\lesssim 0.05 L_{\\odot}$) of low--mass stars. The largest trigonometric parallax survey to date was performed by the Hipparcos satellite \\citep{1997yCat.1239....0E,2007A&A...474..653V}. However, due to the relatively bright magnitude limit of the Hipparcos sample ($V \\lesssim 12$), M dwarfs observed by Hipparcos are saturated in SDSS photometry, hindering the construction of a color--absolute magnitude relation (CMR) in native SDSS filters. A few parallax estimates of SDSS M dwarfs have been obtained \\citep[e.g.,][]{2002AJ....124.1170D, 2004AJ....127.2948V}, but these measurements are observationally taxing and take years to complete, and are thus limited in number ($\\sim$ 20). Current (\\citealp[][]{2010arXiv1008.0648R}; Faherty, private communication) and future parallax studies, such as GAIA \\citep{2001A&A...369..339P} will add valuable observations to this regime. Since trigonometric parallaxes for the vast majority of SDSS M dwarfs are not available, alternate means to estimate their absolute magnitudes must be employed. \\cite{2002AJ....123.3409H} and \\cite{2005PASP..117..706W} used SDSS--2MASS cross--matching to create $M_J$--spectral type relations, and combined with average colors as a function of spectral type, to construct SDSS CMRs. Golimowski et al. (private communication) and \\cite{bochanskithesis} used $u^{\\prime} g^{\\prime} r^{\\prime} i^{\\prime} z^{\\prime}$ photometry of $\\sim 200$ nearby M dwarfs with accurate trigonometric parallaxes, and transformed from the $u^{\\prime} g^{\\prime} r^{\\prime} i^{\\prime} z^{\\prime}$ system of the SDSS photometric telescope to the native SDSS $ugriz$ system using the relations of \\cite{2007AJ....134.2430D}. Other CMRs applicable for warmer stars in the SDSS filter--set have been derived using binaries \\citep{Sesar08} and averages of existing CMRs \\citep{2008ApJ...673..864J}. Using a kinematic model for motions of stars in the Milky Way, \\cite{2010ApJ...716....1B} derived a distance scale and CMR for main sequence stars observed by SDSS. Other studies have used 2MASS--SDSS color transformations \\citep{2009MNRAS.396.1589B} and synthetic photometry \\citep{2007AJ....134.2398C} to derive CMRs in the $ugriz$ filter system. One method for constructing CMRs in the native SDSS system that has not yet been utilized is the classical statistical parallax method. Using proper motions and radial velocities of a large number stars from a homogeneous population, an estimate of the average absolute magnitude, the velocity ellipsoid, and the Sun's peculiar motion can be obtained. This method has a rich history in the astronomical literature \\citep[e.g.,][]{1965BAN....18...71V,1971MNRAS.151..231C,1983veas.book.....M,1998ApJ...506..259P}, and has previously been employed for RR Lyraes \\citep[e.g.,][]{1986ApJ...302..626H,1986MNRAS.220..413S,1996AJ....112.2110L,1998A&A...330..515F,1998ApJ...506..259P} and Cepheids \\citep[e.g.,][]{1991ApJ...378..708W}. In addition to absolute magnitude estimates, the statistical parallax technique provides a measurement of the velocity ellipsoid of the stellar population and the reflex motion of the Sun. The velocity ellipsoid is described by dispersions and directions along three principal axes. Low--mass stars make excellent tracers of the local Galactic potential \\citep{1997PASP..109..559R,2007AJ....134.2418B,2009AJ....137.4149F} as constrained by the observed velocity dispersions. The solar peculiar motion determination complements studies using Hipparcos proper motions of nearby, young stars \\citep{Dehnen98}, SDSS observations of M dwarfs \\citep{2009AJ....137.4149F} and the Palomar-Michigan State University survey of nearby M dwarfs \\citep[PMSU;][]{1995AJ....110.1838R,1996AJ....112.2799H}. In this paper, we use the large spectroscopic sample of M dwarfs described in \\citetalias{west10} to determine absolute magnitudes in the native SDSS filters and examine kinematics using the statistical parallax method. The observations employed in our study are described in \\S \\ref{sec:obs}. We detail our statistical parallax analysis in \\S \\ref{sec:method}. Our results are presented in \\S \\ref{sec:results}, with conclusions following in \\S \\ref{sec:conclusions}. ", "conclusions": "\\label{sec:conclusions} We present a statistical parallax analysis of the most recent catalog of M dwarfs identified with SDSS spectroscopy \\citepalias{west10}. Our sample was subdivided on many criteria, to explore both the intrinsic changes in low--mass dwarfs due to metallicity and magnetic activity, and to investigate their kinematics. We have demonstrated that $r-z$ color is a much better proxy for absolute magnitude than spectral type, and suggest that photometric parallaxes are the preferred method to determine absolute magnitudes and distances for M dwarfs. Some interesting trends with $M_r$ were revealed in our analysis. First, magnetically active M dwarfs were shown to be intrinsically brighter at the same spectral type or color than their inactive counterparts. Eclipsing binary studies have demonstrated that magnetic activity inflates a star's radius \\citep{2006Ap&SS.304...89R, 2007ApJ...660..732L, 2008A&A...478..507M, 2010ApJ...718..502M}, but the effect on effective temperature and luminosity is less constrained. Our results suggest that activity may increase the radius {\\it and} luminosity of active low--mass stars. Metallicity, which influences the luminosity of a star, was also explored. We divided the sample using the $\\zeta$ parameter \\citep{2007ApJ...669.1235L} as a proxy for metallicity. Our results are similar to those observed for higher mass stars: low--metallicity M dwarfs are dimmer at a given spectral type (or color) than high metallicity stars. We have quantified this effect for M dwarfs in the SDSS photometric system. To isolate the effects of metallicity and activity, we separated active and inactive stars for the same $\\zeta$ as a function of color and spectral type. Activity still brightened stars at the same $\\zeta$, however the effect diminished at smaller $\\zeta$ (lower metallicity). The statistical parallax analysis also allowed us to investigate the reflex solar motion and velocity dispersions for each subsample. The more distant, early--type stars, which are presumably older, have the largest reflex solar motion, particularly in the $V$ direction, which we attribute to the increased asymmetric drift. The active and inactive stars exhibit expected behavior with the active star populations having smaller mean motions relative to the Sun. The inactive, late--type M dwarfs, which we identify as an older population due to their lack of activity (ages $\\gtrsim 4$ Gyrs, \\citealp{2008AJ....135..785W}) have velocity dispersions similar to early--type M dwarfs, which we identify as old due to their greater vertical distance from the Galactic plane. Thus, the activity and dynamical heating age indicators give consistent results. When the velocity dispersions are analyzed as a function of vertical height or independent position altitude (IPA), all stars exhibit increasing dispersion at increasing height above the Plane. As astronomy enters a new era of large photometric surveys, such as PanSTARRS \\citep{2002SPIE.4836..154K} and LSST \\citep{2008arXiv0805.2366I}, it will be vital to develop techniques for estimating the activity and metallicity of low--mass stars from photometry alone. There have already been efforts to characterize metallicity using the color--color distributions of M dwarfs in SDSS (\\citealp{2009AIPC.1094..545L}; \\citetalias{west10}), but those data must be calibrated with spectroscopic observations. The work presented in this paper highlights the need to determine these fundamental parameters, since they affect the estimated distance to each star. Finally, the importance of the dual spectroscopic and photometric nature of SDSS cannot be overstated. The large spectroscopic samples of M dwarfs it has acquired have enabled many novel investigations, including this one. Significant spectroscopic followup of the next generation of surveys should be a high priority. We thank Adam Burgasser, Jacqueline Faherty, Rob Simcoe and Kevin Covey for helpful discussions. We thank Neill Reid, Kelle Cruz and Richard Gray for making their nearby stellar spectra available to us. JJB personally acknowledges Roy Halladay for inspiration and motivation throughout this work. JJB thanks the financial support of Adam Burgasser and Kevin Luhman. We also gratefully acknowledge the support of NSF grants AST 02-05875 and AST 06-07644 and NASA ADP grant NAG5-13111. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/. The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington." }, "1101/1101.1857_arXiv.txt": { "abstract": "We briefly review the turbulent mean-field dynamo action in protoneutron stars that are subject to convective and neutron finger instabilities during the early evolutionary phase. By solving the mean-field induction equation with the simplest model of $\\alpha$-quenching we estimate the strength of the generated magnetic field. If the initial period of the protoneutron star is short, then the generated large-scale field is very strong ($> 3 \\times 10^{13}$G) and exceeds the small-scale field at the neutron star surface, while if the rotation is moderate, then the pulsars are formed with more or less standard dipole fields ($< 3 \\times 10^{13}$G) but with surface small-scale magnetic fields stronger than the dipole field. If rotation is very slow, then the mean-field dynamo does not operate, and the neutron star has no global field. ", "introduction": "\\label{aba:sec1} A protoneutron star (PNS) is a very hot ($T\\sim 10^{11} \\; K$), rapidly rotating, lepton rich object that has been formed from the collapse of a massive stellar progenitor. It is believed that lepton and negative entropy gradients generates hydrodynamical instabilities which can play a significant dynamical role in the early stage of the PNS evolution. (Grossman et al. 1993, Bruenn \\& Dineva 1996, Miralles et al. 2000). While convective instability is presumably connected to the entropy gradient, the so-called neutron-finger instability is instead generated by a negative lepton gradient. The latter is due to dissipative processes which are rather fast in PNSs and it grows on a timescale $\\sim 30-100$ ms, that is one or two orders of magnitude longer than the growth time of convection (Miralles et al. 2000). Turbulent motions caused by hydrodynamic instabilities in combination with rotation make turbulent dynamo one of the most plausible mechanism of the pulsar magnetism. The character of turbulent dynamo depends on the Rossby number, $Ro = P/ \\tau$, where $P$ is the PNS spin period and $\\tau$ the turnover time of turbulence. If $Ro \\gg 1$, the effect of rotation on turbulent motions is weak and the mean-field dynamo is inefficient. The small-scale dynamo can be operative, however, even at very large Rossby numbers. If $Ro \\leq 1$ and the turbulence is strongly influenced by the rotation, then the PNS can be subject to a large-scale mean-field dynamo action. Note that the small-scale dynamo can still operate in this case. The Rossby number is typically large in the convective region, $Ro \\sim 10-100$, and the mean-field dynamo is not likely to work in this region. On the contrary, except very slowly rotating PNSs, the Rossby number is of the order of unity, $Ro \\sim 1$, in the neutron-finger unstable region (Bonanno et al. 2003,2005) where turbulent motions are slower, and the turbulence is strongly modified by rotation. This favors the efficiency of mean-field dynamos in the neutron-finger unstable region. This dynamo mechanism is then very different from the one proposed by Thompson \\& Duncan (1993) who argued that only small-scale dynamos can operate in most PNSs. \\def\\figsubcap#1{\\par\\noindent\\centering\\footnotesize(#1)} \\begin{figure}[b]% \\begin{center} \\parbox{2.1in}{\\epsfig{figure=a2.ps,width=1.8in}} \\hspace*{4pt} \\parbox{2.1in}{\\epsfig{figure=cm.ps,width=1.8in}} \\caption{On the left panel is depicted a typical field configuration for $\\alpha^2$-dynamo while in the right panel is depictied a field configuration in the presence of differential rotation. Color levels are for the toroidal field, while dashed lines represents poloidal field lines. }% \\label{fig1.2} \\end{center} \\end{figure} ", "conclusions": "" }, "1101/1101.1250_arXiv.txt": { "abstract": "We present results from mid-IR spectroscopic observations of two young supernova remnants (SNRs) in the Large Magellanic Cloud (LMC) done with the {\\it Spitzer Space Telescope}. We imaged SNRs B0509-67.5 and B0519-69.0 with {\\it Spitzer} in 2005, and follow-up spectroscopy presented here confirms the presence of warm, shock heated dust, with no lines present in the spectrum. We use model fits to {\\it Spitzer} IRS data to estimate the density of the postshock gas. Both remnants show asymmetries in the infrared images, and we interpret bright spots as places where the forward shock is running into material that is several times denser than elsewhere. The densities we infer for these objects depend on the grain composition assumed, and we explore the effects of differing grain porosity on the model fits. We also analyze archival {\\it XMM-Newton} RGS spectroscopic data, where both SNRs show strong lines of both Fe and Si, coming from ejecta, as well as strong O lines, which may come from ejecta or shocked ambient medium. We use model fits to IRS spectra to predict X-ray O line strengths for various grain models and values of the shock compression ratio. For 0509-67.5, we find that compact (solid) grain models require nearly all O lines in X-ray spectra to originate in reverse-shocked ejecta. Porous dust grains would lower the strength of ejecta lines relative to those arising in the shocked ambient medium. In 0519-69.0, we find significant evidence for a higher than standard compression ratio of 12, implying efficient cosmic-ray acceleration by the blast wave. A compact grain model is favored over porous grain models. We find that the dust-to-gas mass ratio of the ambient medium is significantly lower than what is expected in the ISM. ", "introduction": "Supernova remnants (SNRs) provide a laboratory to study various aspects of interstellar medium (ISM) evolution across the whole electromagnetic spectrum. With the advent of high spatial resolution telescopes in both the X-ray and infrared (IR) regimes, it is possible to probe the interaction of the rapidly moving shock wave with the dust and gas of the surrounding ambient medium and to study the ejecta products of the SN itself. The expanding shockwave sweeps up and heats gas to $10^{6}-10^{9}$ K, causing it to shine brightly in X-rays. Dust grains embedded in the hot, shocked plasma are collisionally heated, causing them to radiate at IR wavelengths, and slowly destroying them in the process via sputtering. Because the physical processes behind X-ray and IR emission are related, a combined approach to studying SNRs using both energy ranges can reveal more information than either could on its own. To better characterize IR emission from SNRs, we conducted an imaging survey with the {\\it Spitzer Space Telescope} of $\\sim 40$ known SNRs in the Large and Small Magellanic Clouds. The Clouds were chosen because of their known distance and relatively low Galactic IR background. Subsequently, we obtained IR spectra of several of the SNRs in the LMC using the Infrared Spectrograph (IRS) on {\\it Spitzer}. We report here on IRS observations of two of these, SNRs B0509-67.5 (hereafter 0509) and B0519-69.0 (hereafter 0519). Both are remnants of thermonuclear SNe \\citep{smith91,hughes95} and have fast, non-radiative shocks (several thousand km s$^{-1}$) \\citep{tuohy82,ghavamian07}. There is no evidence in either for slower, radiative shocks. They are both located in the LMC, at the known distance of $\\sim 50$ kpc. In addition, both have ages determined from light echoes \\citep{rest05}, with 0509 being 400 $\\pm 120$ and 0519 being 600 $\\pm 200$ years old. They are nearly identical in size, having an angular diameter of $\\sim 30$'', which corresponds to a physical diameter of $\\sim 7.3$ pc. In \\citet{borkowski06} hereafter Paper I, we used 24 and 70 $\\mu$m imaging detections of both SNRs to put limits on the post-shock density and the amount of dust destruction that has taken place behind the shock front, as well as put a limit on the dust-to-gas mass ratio in the ambient medium, which we found to be a factor of several times lower than the standard value for the LMC of $\\sim 0.25$\\% (Weingartner \\& Draine 2001, hereafter WD01). These limits, however, were based on only one IR detection, at 24 $\\mu$m, with upper limits placed on the 70 $\\mu$m detection. Both remnants were detected with {\\it Akari} \\citep{seok08} at 15 and 24 $\\mu$m (0519 was also detected at 11 $\\mu$m). Seok et al. applied single-temperature dust models to {\\it Akari} data, deriving warm dust masses of 8.7 $\\times 10^{-5}$ $M_\\odot$ and $3.6 \\times 10^{-4}$ $M_\\odot$ in 0509 and 0519, respectively. With full spectroscopic data, we can place much more stringent constraints on the dust destruction and dust-to-gas mass ratio in the ISM. We also explore alternative dust models, such as porous and composite grains. Additionally, we examine archival data from the Reflection Grating Spectrometer (RGS) onboard the {\\it XMM-Newton} X-ray observatory. The high spectral resolution of RGS allows us to measure the strength of lines in X-ray spectra. We can use post-shock densities derived from IR fits to predict the strength of lines arising from shocked ambient medium, and can thus derive relative strengths of ejecta contributions to oxygen lines. Doing this requires knowledge of the shock compression ratio, $r$, defined as $n_{H}/n_{0}$ (where $n_{H}$ and $n_{0}$ are postshock and preshock hydrogen number densities, respectively), which for a standard strong shock (Mach number $\\gg$ 1) is 4. However, shock dynamics will be modified \\citep{jones91} if the shock is efficient at accelerating cosmic rays, and $r$ can be increased by a factor of several. We use several representative values of $r$ in modeling X-ray line strengths. ", "conclusions": "We present mid-IR spectral observations of two young SNRs in the LMC, 0509-67.5 and 0519-69.0, as well as analysis of archival high-resolution X-ray spectra. By fitting dust heating and sputtering models to IR spectra, we can determine post-shock gas density. We derive post-shock densities of 0.59 and 6.2 cm$^{-3}$ in 0509 and 0519, respectively, using a compact grain model. Porous grains require a density higher by a factor of a few. By assuming a value for the shock compression ratio, $r$, we can infer the pre-shock density, swept gas mass, strength of oxygen Ly$\\alpha$ and K$\\alpha$ X-ray lines, and dust-to-gas mass ratio for the ambient ISM. Derived values for the pre-shock density of the ISM vary greatly depending on the model used and compression ratio assumed, but for the sake of comparison, values listed in Tables 2 \\& 3 are comparable to densities inferred from various line-of-sight HI column densities through the LMC. These column densities vary from 10$^{20}$ cm$^{-2}$ to a few $\\times 10^{21}$ cm$^{-2}$, which, assuming an LMC depth of 1 kpc, yields densities of $\\sim 0.05-1.5$. For standard strong shocks, $r$ is 4, but modification by cosmic rays would increase this by an unknown amount. We report values of the quantities above for several plausible values of $r$. In principle, this method could be used in reverse to determine $r$. This would require X-ray emission from shocked ambient medium that is well-separated from ejecta emission; thus, X-ray lines could be modeled as solely arising from material at cosmic abundances. We believe this method can be used for older, larger remnants in the LMC, such as DEM L71. Based on our analysis, a significant fraction of the O line emission seen in the X-ray spectrum of 0509 arises from the ejecta. If grains are highly porous, then the ejecta contribution is less, but a contribution is still required because the line flux ratios observed still do not match the model ratios. Both standard shocks and cosmic-ray modified shocks with higher $r$ values can provide acceptable fits for 0509. The data for 0509 are inconclusive in favoring either compact or porous grains. Compact grains require nearly all of the oxygen seen in X-ray spectra to be coming from shocked ejecta, while porous grains would significantly lower the ejecta contribution, implying a higher contribution from the forward-shocked material. In 0519, there is significant evidence for higher than standard compression, and little contribution from ejecta to O X-ray lines. A shock speed of $<$ 2300 km s$^{-1}$ is favored for this object from hydrodynamical simulations, in order to reproduce the observed size, age, and X-ray line widths. Compact grains are favored for this remnant. The standard value of $r=4$ is possible only if absorption to the remnant is significantly higher than has been reported. The column density from the LMC would need to be at least $4 \\times 10^{21}$ cm$^{-2}$ to bring oxygen lines in the model down to below observed values. \\citet{ghavamian07} derive an upper limit for 0519 of $E(B-V) \\le 0.11$, implying an upper limit to the HI column density of $\\sim 1.6$ $\\times 10^{21}$ cm$^{-2}$ \\citep{cox06}. We derive dust-to-gas mass ratios that are lower by a factor of several than what is generally expected in the ISM. Since the general properties of dust in the ISM come from optical/UV absorption line studies averaged over long lines of sight, it is possible that there exist large local variations on smaller scales. The ratios presented here probe parsec scales. Another potential explanation is that the sputtering rate for dust grains is significantly understated in the literature. However, since the deficit of dust is about an order of magnitude, sputtering rates would also have to be increased by this amount to account for the discrepancy. Also, because sputtering rates typically assume a compact grain, further work in this field will be needed to determine if these rates are appropriate for porous grains. Extremely high values for $r$ ($>$ 50) could bring dust-to-gas mass ratios up to more typical values, but such values are ruled out by radio observations. A possible observational bias could exist, in that SNRs that are easy to study in the IR (i.e., well separated from sources of IR confusion) could be found preferentially in low-dust regions of the ISM. Studies such as this will benefit greatly from the increases in both spatial resolution and sensitivity of future generations of telescopes, such as the {\\it James Webb Space Telescope}. Being able to spatially separate the dust spectra right behind the shock from that further inside the shell will be crucial to reducing some of the uncertainties listed above. Longer wavelength observations with the {\\it Herschel Space Observatory} would be useful for detecting or placing upper limits on any cold dust that may be present in the remnant." }, "1101/1101.1299_arXiv.txt": { "abstract": "We analyze the relation between the mass of the central supermassive black hole (\\mbh) and the number of globular clusters (\\ngc) in elliptical galaxies and bulges as a ramification of the black hole fundamental plane, the theoretically predicted and observed multi-variable correlation between \\mbh\\ and bulge binding energy. Although the tightness of the \\mbh--\\ngc\\ correlation suggests an unlikely causal link between supermassive black holes and globular clusters, such a correspondence can exhibit small scatter even if the physical relationship is indirect. We show that the relatively small scatter of the \\mbh--\\ngc\\ relation owes to the mutual residual correlation of \\mbh\\ and \\ngc\\ with stellar mass when the velocity dispersion is held fixed. Thus, present observations lend evidence for feedback-regulated models in which the bulge binding energy is most important; they do not necessarily imply any `special' connection between globular clusters and \\mbh. This raises the question of why \\ngc\\ traces the formation of ellipticals and bulges sufficiently well to be correlated with binding energy. ", "introduction": "\\label{s:intro} There are now well-established correlations between the mass of supermassive black holes (SMBHs) and properties of their host galaxies, such as bulge luminosity, mass, light concentration, and velocity dispersion \\citep{kr95, magorrian98, ferrarese00, gebhardt00, tremaine02, marconi03, gultekin09}. This suggests that the physical mechanism driving growth of the SMBH also plays a key role in forming the bulge (for spiral galaxies) or galaxy (for ellipticals). Analytical estimates \\citep{sr98,burkert01,hh06}, as well as numerical simulations \\citep{dsh05,sdh05,cox06_kinematics,robertson06_msigma,croton06,johansson09} with simple prescriptions for SMBH accretion have demonstrated the plausibility of this inference by matching the expected slopes of these correlations. Regardless of the detailed feedback prescription, these models predict that SMBHs grow until reaching some critical mass, where the energy and/or momentum released by feedback expels material from the nucleus. As such, they robustly predict that the ``true'' correlation should be between SMBH mass and a quantity such as the binding energy or potential well depth of material in the bulge. \\citet{hopkinsfp_07a} show that the observed correlations with different variables, and importantly their scatter and systematic deviations from the relations, can be understood as the projections of a single fundamental dependence. This relation is approximated closely by a multi-variable correlation, a black hole fundamental plane (BHFP). \\citet{ar07} confirmed this in a sample of ellipticals and spiral bulges using dynamical models of bulge potentials, and \\citet{fm09} did so with simple proxies such as $M_{BH} \\propto E_b \\sim M_* \\sigma^2$. Additional correlations have been found between SMBH mass and dark matter halo mass, as well as the number \\ngc\\ of globular clusters (GC) in the host galaxy \\citep{sf09, bt10, harris10}. In particular, \\citet[][hereafter BT10]{bt10} argued that \\ngc\\ is a better predictor of \\mbh\\ than the velocity dispersion $\\sigma$, citing a smaller intrinsic scatter and a residual correlation between \\ngc\\ and \\mbh\\ in elliptical galaxies even after accounting for the median $M_{BH}-\\sigma$ correlation, suggesting a fundamental link between the accretion of gas by the SMBH and the formation of a galaxy's globular cluster system. \\citet[][hereafter HH10]{harris10} extended the sample by making reasonable estimates of \\ngc\\ from the literature in galaxies with \\mbh\\ measurements. In this letter, we illustrate that the above link can be understood as a consequence of the BHFP relation combined with a residual correlation between \\ngc\\ and the bulge's stellar mass $M_*$ at fixed $\\sigma$. Rather than suggesting a single ``best'' correlation between \\mbh\\ and a single galaxy parameter, the BHFP implies that the best predictor of SMBH mass is some combination thereof. For example, \\mbh\\ has a positive correlation with the bulge's stellar mass even at fixed $\\sigma$. Although the number of globular clusters in a particular galaxy, like \\mbh, is a complex function of the galaxy's formation history, there exists a similar positive residual correlation between \\ngc\\ and $M_*$, so that the resulting \\ngc--\\mbh\\ residuals (fixing $\\sigma$) will be positively correlated. In \\S\\ref{s:data} we describe a sample of 32 elliptical galaxies from \\citet{peng08} with auxiliary data compiled in \\citet[][and subsequent papers]{hopkins08_mr}. In \\S\\ref{s:correlations} we fit separately the relations $M_*$--$\\sigma$ and \\ngc--$\\sigma$ in these galaxies to establish the residual correlation between \\ngc\\ and $M_*$. Then we combine this residual slope with knowledge of the \\mbh--$M_*$ correlation at fixed $\\sigma$ from the BHFP, and calculate the residual correlation and scatter expected between \\ngc\\ and \\mbh. We summarize and conclude in \\S\\ref{s:conclusions}. ", "conclusions": "\\label{s:conclusions} We have shown that the number of globular clusters in elliptical galaxies exhibits a residual dependence on $M_*$ at fixed $\\sigma$, implying that the bulge binding energy ($\\sim M_* \\sigma^2$) is a better indicator of \\ngc\\ than $\\sigma$ or $M_*$ alone. The same was shown to be true for \\mbh\\ by \\citet{hopkinsfp_07a}, as these parameters constitute a formulation of the BHFP. Thus the apparent power of \\mbh--\\ngc\\ versus \\mbh--$\\sigma$ owes to the fact that \\ngc\\ and \\mbh\\ are both tracers of the same fundamental property such as the bulge binding energy. This resolves several puzzling aspects of the previous interpretation of the data. As BT10 themselves point out, there cannot be a direct causal correlation between \\ngc\\ and \\mbh, since most of the GC mass is at very large radii and has never had any interaction with the galaxy nucleus. Moreover, while most GCs likely formed at very high redshift, the final mass of the SMBH is sensitive to its growth via gas accretion at $z\\lesssim2$ \\citep[e.g.][]{hh06,hrh07,hopkins08_ell}. However, this naturally predicts that \\ngc\\ should serve reasonably well as a mass tracer, so that the dependence of \\mbh\\ on $M_*$ and formation time leads to a surprisingly tight but expected \\ngc--\\mbh\\ correlation. The same arguments explain the result in \\citet{hmt09}, who show that the observed \\mbh\\ is sensitive to the entire galaxy baryonic mass -- i.e. perhaps the mass traced by \\ngc\\ is the same as the mass that actually sets the escape velocity and potential well depth at $R=0$, rather than just the stellar mass enclosed in a small radius around the BH, which can vary widely in systems of similar \\mbh. Such a relation between \\ngc\\ and global galaxy mass or luminosity has been demonstrated \\citep[e.g.][]{hvdb81, mclaughlin99}, and this trait supports the idea that \\ngc\\ and \\mbh\\ are connected indirectly by a more fundamental galaxy property. This also naturally explains why HH10 find that the relation breaks down for S0 galaxies. These galaxies are structurally different than ellipticals and may have different formation histories \\citep{larson80}, so \\ngc\\ and the total stellar mass may not faithfully trace the bulge binding energy. Since S0's are not particularly discrepant in \\mbh--$\\sigma$ \\citep[e.g.][and previous works]{gultekin09}, this suggests that the \\ngc--bulge relation is the connection that weakens for these systems. HH10 also find no statistically significant correlation in spirals: although three out of the four spirals from HH10 lie on the \\ngc--\\mbh\\ relation, there simply isn't yet enough data to know for sure if this relation persists for spiral bulges. However, the underlying BHFP relation ties \\mbh\\ to the binding energy and explains its residual correlations with bulge parameters, even for these disky galaxies where \\ngc\\ possibly deviates. This alone suggests that the BHFP, not \\mbh--\\ngc\\ or \\mbh--$\\sigma$, is the `more fundamental' correlation. \\citet{hopkinsfp_07b} showed that the existence of a black hole fundamental plane is a robust prediction of numerical simulations of gas-rich mergers that include the effects of gas dissipation, cooling, star formation, and black hole accretion and feedback. The present work shows that this local and widely expected correlation between supermassive black hole mass and bulge binding energy in feedback-regulated scenarios, combined with a similar correlation for \\ngc, can account for the observed \\ngc--\\mbh\\ relation and its scatter. The interesting question raised by such a correlation is {\\em not} why \\ngc\\ correlates tightly with \\mbh, since this is indirect, but why \\ngc\\ correlates tightly with galaxy binding energy/potential well depth. Some such correlation is expected and observed \\citep{mclaughlin99, blakeslee99, peng08}: an example is that systems at fixed velocity dispersion with higher stellar mass have accreted or formed more stars, likely including globular clusters. But that the \\ngc--bulge relation should be so tight, and include both metal-rich and metal-poor populations, may support the inferences by BT10 and HH10 (and references therein) that the formation of globular cluster systems and growth of supermassive black holes in elliptical galaxies are driven by a common galaxy property." }, "1101/1101.1911_arXiv.txt": { "abstract": "{Baryon Acoustic Oscillations (BAO) are a feature imprinted in the density field by acoustic waves travelling in the plasma of the early universe. Their fixed scale can be used as a standard ruler to study the geometry of the universe.} {BAO have been previously detected using correlation functions and power spectra of the galaxy distribution. In this work, we present a new method for the detection of the real-space structures associated with this feature. These baryon acoustic structures are spherical shells with a relatively small density contrast, surrounding high density central regions.} {We design a specific wavelet adapted to the search for shells, and exploit the physics of the process by making use of two different mass tracers, introducing a specific statistic to detect the BAO features. We show the effect of the BAO signal in this new statistic when applied to the $\\Lambda$ -- Cold Dark Matter ($\\Lambda$CDM) model, using an analytical approximation to the transfer function.We confirm the reliability and stability of our method by using cosmological $N$-body simulations from the MareNostrum Institut de Ci\\`encies de l'Espai (MICE).} {We apply our method to the detection of BAO in a galaxy sample drawn from the Sloan Digital Sky Survey (SDSS). We use the `Main' catalogue to trace the shells, and the Luminous Red Galaxies (LRG) as tracers of the high density central regions. Using this new method, we detect, with a high significance, that the LRGs in our sample are preferentially located close to the centres of shell-like structures in the density field, with characteristics similar to those expected from BAOs. We show that stacking selected shells, we can find their characteristic density profile.} {We have delineated a new feature of the cosmic web, the BAO shells. As these are real spatial structures, the BAO phenomenon can be studied in detail by examining those shells.} ", "introduction": "\\label{sec:intro} Before recombination, the energy of photons is high enough to avoid the formation of neutral hydrogen atoms. This means that baryons and photons are coupled through Compton scattering and electromagnetic interaction between protons and electrons, forming a plasma. In this fluid two phenomena act in opposite directions: gravitational forces tend to compress the plasma around high density regions, while radiation pressure tends to dilute any such over-density. The combination of both in the presence of any initial inhomogeneity give rise to acoustic waves propagating in the baryon-photon plasma. This phenomenon ends abruptly at the epoch of recombination, when the temperature drops enough to allow hydrogen atoms to form, and therefore radiation decouples from the baryons. Baryon acoustic oscillations (BAO) are therefore due to the propagation of these sound waves in the baryon-photon plasma in the early universe \\citep{pee70a,hu97a,eis98a,bas09a}. Any primordial over-density in the early universe produces a spherical acoustic wave in the baryon-photon plasma, travelling outwards: the radiation pressure drags the baryons that are coupled to the photons, and compensates the gravity force that pulls all matter towards the centre. Dark matter, however, is totally decoupled from the photons, and therefore its density at the centre continues growing. About $380,000$ years after the Big Bang, temperature drops so that photons and baryons decouple, and the scale of the baryon shells freezes. After this time, both the central over-density and the shell grow gravitationally, accreting both dark matter and baryons. The result at late times is a large over-density at the position of the original perturbation, surrounded by a faint spherical shell at a fixed co-moving scale \\citep{eis07a}. The BAO scale is fixed by the sound horizon at decoupling: it is the distance that the expanding acoustic shells can travel before decoupling. It has been accurately measured by the study of the anisotropies in the Cosmic Microwave Background (CMB) to be \\citep{kom08a} $r_s = 153.3 \\pm 2.0 \\, \\mathrm{Mpc} = 110.4 \\pm 1.4 \\hMpc$ (where we take $h = 0.72$, \\citealp{fre01a})\\footnote{$h$ is the Hubble constant in units of $100 \\kms\\, \\mathrm{Mpc}^{-1}$}. Therefore, this scale, once measured, could be used as a standard ruler to measure the Hubble expansion rate with redshift $H(z)$ and the angular diameter distance $D_A(z)$ \\citep{coo01a,bla03c,seo05a}. The BAO should appear as a series of damping wiggles in the matter power spectrum, with the locations of the peaks and throughs in $k$-space being a function of $r_s$ and other cosmological parameters \\citep{eis98a}. All the harmonics sum up to the same peak in the galaxy correlation function $\\xi(r)$ at the scale $r_s$, and therefore it could seem more appropriate to use this statistic for the detection of the BAO feature on the available galaxy redshift surveys encompassing large volumes of the universe \\citep{san08b}. The first detection (claiming a $\\sim 3\\sigma$ level) was reported in the analysis of the correlation function \\citep{eis05a} of the Sloan Digital Sky Survey (SDSS) \\citep{yor00a} Luminous Red Galaxies (LRG) sample \\citep{eis01a}, and later in the power spectrum \\citep{col05a} of the 2-degree Field Galaxy Redshift Survey (2dFGRS) \\citep{col01a}. But certainly this is a controversial topic. \\citet{cab10a} are not finding such level of detection using a data set twice as large in volume and in number of galaxies. They do not claim this result to be in contradiction with the standard $\\Lambda$CDM model, but to be a consequence of insufficient data. One of the arguments in \\citet{cab10a} is the fact that mixing model selection with parameter determination can lead to some confusion in the interpretation of the results and their significance. Different authors are using different criteria to assess the significance of ther BAO detection. For example, when \\citep{eis05a} affirm that the baryon signature was detected at 3.4 $\\sigma$ (or at 3.0 $\\sigma$ when including only data points between 60 and 180 $h^{-1}$ Mpc) they are comparing their results of the SDSS-LRG correlation function with the expected for the best-fit pure CDM model and different BAO models. The best BAO detection up to now \\citep{per10a} was obtained studying the combined power spectrum of LRG and `Main' \\cite{str02a} samples of SDSS, together with the 2dFGRS sample, and is at the $\\sim 3.6\\sigma$ level. The authors explicitly state that since this number is obtained comparing to an arbitrary smooth model, the significance cannot be directly compared with the one reported in \\citep{eis05a}. This is a clear example of different authors using different ways to assess the significance of their results that in practice are not comparable. \\citet{hut06a} calculated the redshift space power spectrum of the SDSS-LRG sample drawn from the Data Release 4. He concludes that BAO models are favored by $3.3 \\sigma$ over the corresponding models without any oscillatory behavior in the power spectrum. \\citet{per07a} detected BAOs in the clustering of the combined 2dFGRS and SDSS main galaxy samples, and use their measurements to constrain cosmological models, in particular a given combination of the angular diameter distance $D_A(z)$ and the Hubble parameter $H(z)$. \\citet{cab08a,cab08b} studied the LRGs anisotropic redshift space correlation function $\\xi(\\sigma,\\pi)$, where $\\pi$ is the line-of-sight or radial separation and $\\sigma$ is the transverse separation. Moreover, \\citet{gaz08b} have shown how to constrain $H(z)$ using the correlations in the radial direction. \\citet{kaz10a} found similar results for the correlation measurements and uncertainties, but manifest disagreement in the interpretation of the results regarding the detection of a line-of-sight baryonic acoustic feature. More recent studies \\citep{mar08a,cab08a,san09a,kaz09a} have confirmed this detection in the last Data Release (DR7, \\citealp{aba08a}) of the SDSS-LRG, containing twice as many galaxies as the original sample, although the observed peak is in these cases wider than that observed in the original detection -- an issue that needs further explanation. These measurements of the BAO scale at a low redshift, combined with other cosmological probes, have been used to put stringent constraints on the values of cosmological parameters \\citep{teg06a,per07a,san09a,per10a,rei09a,kaz10a}. While \\citet{bas10a} argue that low-level detections may not be sufficient to robustly estimate the cosmological parameters, \\citet{cab10a} show instead that it is still possible --assuming a model-- to locate the BAO position with data providing very low significant BAO detection. It is important, therefore, to find evidence of BAO in the galaxy distribution based on complementary methods. A step further is to search for real structures in the galaxy distribution that are responsible for the BAO feature in these second order statistics. The detection of these structures would be a confirmation of the existence of the baryon acoustic phenomenon. Moreover, if we are able to localize these structures in configuration space, this would allow us to study in more detail the properties of BAO. In this paper, we introduce a new method for the detection of BAO, which is closely tied to the underlying physics of the process, and apply it to a sample drawn from the SDSS catalogue. This method (described in Section~\\ref{sec:method}) is based on analyzing directly the 3D galaxy distribution using a very specific wavelet function (which we called `BAOlet'), which is specially well suited to search for BAO features. The method makes use of two different tracers, one to map the overall density field (including the BAO shells), and the other to locate the position of the largest overdensities, which should correspond to centres of the shells. As we study directly the galaxy distribution in configuration space, this method also allows us to identify regions of space where the BAO signal is stronger or fainter. We describe the expected signal in the $\\Lambda$CDM model in Section~\\ref{sec:theory}, using both analytical prediction and a $N$-body simulation catalogue. We describe the samples used in the case of SDSS in Section~\\ref{sec:data}. In Section~\\ref{sec:results}, we show the results obtained in this case. We also make a test to assess the significance of these results, and explore the implications of this analysis regarding the localization of BAO structures. Finally, we summarize our conclusions and discuss possibilities for future work in Section~\\ref{sec:conc}. ", "conclusions": "\\label{sec:conc} In summary, we have designed a new method for the detection of baryon acoustic oscillations in the galaxy distribution and for the localization, in configuration space, of the structures responsible for them. This method is based on the use of a specially designed wavelet applied directly on the density field. Our approach also relies on the use of two different tracers: one for the overall density distribution, and the other for the central overdensities of the baryon acoustic structures. After testing the method with simulations, we applied this method to the detection of baryon acoustic structures in a sample drawn from the SDSS. In this case, we used galaxies from the `Main' catalogue to trace the overall density field, and galaxies from the LRG catalogue to trace the location of massive dark matter haloes. We clearly detect BAO in the sample providing a confirmation of the detection obtained previously using general two point statistics (the power spectrum and correlation function). In fact, our approach provides an independent method for the detection. Finally, we showed how this method allows us, through the use of $W_{\\rm max}(\\mathbf{x})$, to localize in configuration space the actual structures responsible for the BAO signal obtained. This is a consequence of using a wavelet acting directly on the density field. We illustrate the utility of this approach by showing the density distribution stacked around a set of centres known to show the BAO feature given their $W_{\\rm max}$ value. Recent works have proposed alternative methods to study the BAO based on wavelets \\citep{xu10a, tia11a}. In particular, \\citeauthor{tia11a} use a Mexican hat wavelet function with two parameters, conceptually similar to ours. They use it to search for a peak in the two point correlation function of the `Main' SDSS sample, obtaining a detection with a $p$-value $p = 0.002$ (equivalent to $3.1\\sigma$ in the Gaussian case). As in our case, this shows the utility of using the `Main' sample to reduce the shot noise in the calculation and to obtain significant detections. However, these works apply the wavelet to the measured two point correlation function, instead of directly to the density field. In this way, the use the capabilities of the wavelets to characterize accurately the BAO signal (in terms of radius and width), but they are not able to get any information about the localization of these structures in space. The use of wavelets directly on the density field isolates valuable information about the baryon acoustic structures that is hidden in the standard two point statistics. In particular it gives us information, through the coefficients $W_{R,s}(\\mathbf{x})$, to localize regions in the sampled volume giving the largest or lowest signal. We expect that this new method for studying BAO will be of much use for ongoing or planned surveys, such as the WiggleZ Survey \\citep{dri10a}, the Baryon Oscillation Spectroscopic Survey \\citep[BOSS,][]{eis11a}, or the Physics of the Accelerating Universe (PAU) Survey \\citep{ben08a}, which will cover a much larger volume than studied here, and will explore higher redshifts." }, "1101/1101.4019_arXiv.txt": { "abstract": "Recent measurements by the Pierre Auger Observatory suggest that the composition of ultra-high energy cosmic rays (UHECRs) becomes dominated by heavy nuclei at high energies. However, until now there has been no astrophysical motivation for considering a source highly enriched in heavy elements. Here we demonstrate that the outflows from Gamma-Ray Bursts (GRBs) may indeed be composed primarily of nuclei with masses $A \\sim 40-200$, which are synthesized as hot material expands away from the central engine. In particular, if the jet is magnetically-dominated (rather than a thermally-driven fireball) its low entropy enables heavy elements to form efficiently. Adopting the millisecond proto-magnetar model for the GRB central engine, we show that heavy nuclei are both synthesized in proto-magnetar winds and can in principle be accelerated to energies $\\gtrsim 10^{20}$ eV in the shocks or regions of magnetic reconnection that are responsible for powering the GRB. Similar results may apply to accretion-powered GRB models if the jet originates from a magnetized disk wind. Depending on the precise distribution of nuclei synthesized, we predict that the average primary mass may continue to increase beyond Fe group elements at the highest energies, possibly reaching the $A \\approx 90$ (Zirconium), $A \\approx 130$ (Tellurium), or even $A \\approx 195$ (Platinum) peaks. Future measurements of the UHECR composition at energies $\\gtrsim 10^{20}$ eV can thus confirm or constrain our model and, potentially, probe the nature of GRB outflows. The longer attenuation length of ultra-heavy nuclei through the extragalactic background light greatly expands the volume of accesible sources and alleviates the energetic constraints on GRBs as the source of UHECRs. ", "introduction": "\\label{intro} The origin of Ultra-High Energy Cosmic Rays (UHECRs) is one of the great mysteries in high energy astrophysics (e.g.~\\citealt{Blandford&Eichler87}; \\citealt{Nagano&Watson00}). UHECRs are generally thought to originate from extra-galactic distances: the break observed in the cosmic ray spectrum at $\\sim 3\\times 10^{18}$ eV (the `ankle') is often interpreted as the energy beyond which the Galactic magnetic field can neither isotropize UHECRs nor appreciably prolong their residence time in the Galaxy (e.g.~\\citealt{Hillas05}). An extra-galactic origin is also suggested by the cut-off observed in the spectrum above $\\sim 6\\times 10^{19}$ eV (\\citealt{Abraham+08}; \\citealt{Abbasi+08}), which is generally interpreted as the result of UHECRs (protons or heavy nuclei) interacting with the cosmic microwave background (CMB) and other sources of extragalactic background light (EBL). This is the `GZK' effect, initially proposed for protons by \\citet{Greisen66} and \\citet{Zatsepin&Kuzmin66}. Only a handful of astrophysical sources are plausible sites for accelerating UHECRs because the requirements on the magnetic field, compactness, and energy budget are stringent. Commonly discussed candidates can be divided into persistent and transient sources. Persistent sources include powerful relativistic jets from Active Galacti Nuclei (AGN; e.g.~\\citealt{Mannheim&Biermann92}; \\citealt{Berezinsky+02}; \\citealt{Farrar&Gruzinov09}; \\citealt{Dermer+09}; \\citealt{Takami&Horiuchi10}), weaker AGN jets (e.g.~\\citealt{Honda09}; \\citealt{Peer+09}), and galaxy clusters (e.g.~\\citealt{Inoue+07}; \\citealt{Kotera+09}). Candidate transient sources (\\citealt{Waxman&Loeb09}) include classical Gamma-Ray Bursts (GRBs; \\citealt{Waxman95}; \\citealt{Vietri95}; \\citealt{Milgrom&Usov95}; \\citealt{Waxman04,Waxman06}; \\citealt{Dermer10}), low luminosity GRBs (e.g.~\\citealt{Murase+06}), AGN flares (e.g.~\\citealt{Farrar&Gruzinov09}), and relativistic (`engine-driven') supernovae (e.g.~\\citealt{Chakraborti+10}). The arrival directions of UHECRs provide a potentially important probe of their origin. Measurements by the Pierre Auger Observatory (PAO) rule out isotropy for the highest energy cosmic rays at $\\sim98\\%$ confidence (\\citealt{Armengaud+08}), and PAO has furthermore discovered a correlation between the arrival directions of UHECRs with energies $E > 57$ EeV and nearby ($\\lesssim 75$ Mpc) AGN (\\citealt{Abraham+08}). This result does not, however, imply that UHECRs necessarily originate from AGN, because AGN trace local Galactic structure, such that the correlation is consistent with a variety of other sources (\\citealt{Kashti&Waxman08}; \\citealt{Ghisellini+08}; \\citealt{Takami+09}; \\citealt{Takami&Sato09}).\\footnote{In addition, the latest PAO results suggest that the significance of the AGN correlation is reduced from previous measurements (\\citealt{Abraham+09}).} At present the sources of UHECR cannot therefore be deduced from their arrival directions alone. The composition of UHECRs also provides important clues to their origin. Although the composition is measured directly at low energies ($\\lesssim 10^{14}$ eV), at ultra-high energies it must be inferred indirectly by measuring the shower depth at maximum elongation $X_{\\rm max}$. Recent measurements by PAO show that the average shower depth $\\langle X_{\\rm max} \\rangle$ and its RMS variation decrease systematically moving to the highest energies \\citep{Abraham+10}. This suggests that the UHECR composition transitions from being dominated by protons below the ankle to being dominated by heavier nuclei with average masses similar to Si or Fe at $\\sim 5\\times 10^{19}$ eV. We caution, however, that HiRes has not verified this finding \\citep{Abbasi+05}.\\footnote{One possible explanation for this discrepency is that PAO and HiRes observe different portions of the sky. We note that HiRes also does not find the correlation of UHECRs with AGN seen by PAO (e.g.~\\citealt{Sokolsky&Thomson07}).} The UHECR composition measured by Auger is puzzling. One possible explanation is that the accelerated material has an intrinsically `mixed' composition (with e.g.~solar abundances), such that protons are accelerated to a maximum energy $E = E_{\\rm p,max} \\sim 10^{18.5}$ eV, beyond which only heavier nuclei are accelerated. This seems plausible {\\it a priori} because accelerator size considerations show that the maximum achievable energy increases linearly with the nuclear charge $Z$ \\citep{Hillas84}. On the other hand, this explanation appears to require fine tuning because the maximum energy to which, for instance, Fe nuclei are accelerated $E_{\\rm Fe,max} \\sim Z\\times E_{\\rm p,max} \\sim 8\\times 10^{19}$ (Z/26) eV must (by coincidence) be close to the cut-off observed at $\\sim 6\\times 10^{19}$ eV and expected to occur independently from the GZK effect. A `mixed' composition with metal abundance ratios similar to the Sun or Galactic cosmic rays also appears inconsistent with modeling of the propagation of UHECRs through the EBL (\\citealt{Allard+08}), which suggest that the injected composition has a fairly narrow distribution in charge (e.g.~\\citealt{Hooper&Taylor10}). A second possibility is that the accelerated material is dominated by heavy nuclei. In this case the proton-dominated composition measured near the ankle may be explained as secondary particles produced by the interaction of the nuclei with the EBL (e.g.~\\citealt{Hooper&Taylor10}). A heavy-rich composition is unlikely in the case of AGN, galaxy clusters, and supernova shocks because the accelerated material originates from the interstellar medium. For a solar composition, the fraction of the total mass in Fe nuclei and heavier is just $X_{\\rm Fe} \\sim 10^{-3}$, such that only for extremely super-solar metallicity ($\\sim 10^{3} Z_{\\odot}$) could heavy nuclei dominate the total UHECR mass. In this paper we show that UHECRs from GRBs, unlike AGN, may indeed be composed of almost entirely very heavy nuclei. In particular, if the outflow from the central engine is strongly magnetized we find that Fe-group nuclei and possibly heavier elements (A $\\gtrsim$ 90) are synthesized during its expansion. Although it is well established that long duration GRBs originate from the core collapse of massive stars \\citep{Woosley&Bloom06}, it remains debated whether the central engine is a hyper-accreting black hole \\citep{Woosley93} or a rapidly spinning, strongly magnetized neutron star (a `proto-magnetar'; e.g.~\\citealt{Usov92}). We focus here on the proto-magnetar model, which recent work has shown can explain many of the observed properties of GRBs (\\citealt{Thompson+04}; \\citealt{Metzger+07}; \\citealt{Bucciantini+07}; \\citealt{Metzger+10}). However, similar considerations may apply to accretion-powered models, provided that the jet is magnetically-dominated rather than a thermally-driven fireball ($\\S\\ref{sec:BH}$). ", "conclusions": "\\label{sec:discussion} The spectrum and composition of UHECRs measured by Auger are consistent with an accelerated composition dominated by heavy nuclei. However, until now there has been no astrophysical motivation for considering such a source. In this paper, we have shown that magnetically-dominated GRB outflows may synthesize heavy $A \\approx 40-200$ nuclei as they expand away from the central engine (Fig.~\\ref{fig:Xh}). Focusing on the millisecond proto-magnetar model, we have shown that the regions of magnetic reconnection or shocks responsible for the GRB emission also allow heavy nuclei to be accelerated to ultra-high energies $E_{\\rm max} \\gtrsim 10^{20}$ eV while not being disintegrated by GRB photons (Fig.~\\ref{fig:emax}). Models that invoke GRBs as the source of UHECRs have been criticized on the grounds that the required energetics may be insufficient if the rate and observed gamma-ray fluences from GRBs are a proxy for the UHECR flux (e.g.~\\citealt{Farrar&Gruzinov09}; \\citealt{Eichler+10}). However, this argument depends on the (uncertain) local rate of GRBs (e.g.~\\citealt{Le&Dermer07}), the fraction of the jet energy used to accelerate baryons (versus electrons; e.g.~\\citealt{Sironi&Spitkovsky11}), and the slope of the injected UHECR spectrum. It is thus interesting to note that a shallow injected energy spectrum both alleviates GRB energetic constraints ($\\S\\ref{sec:mag}$) and may be necessary to fit the UHECR spectrum and composition measured by Auger if the injected composition is indeed dominated by heavy nuclei (e.g.~\\citealt{Hooper&Taylor10}). \\citet{Metzger+10} found that observed GRB spectra were best understood in the proto-magnetar model if magnetic reconnection ($\\S\\ref{sec:reconnection}$) was responsible for powering the prompt gamma-ray emission rather than shocks ($\\S\\ref{sec:shocks}$). It is thus also important to note that numerical studies of magnetic reconnection indeed tend to predict flat accelerated spectra (e.g.~\\citealt{Romanova&Lovelace92}). Because the electron fraction in GRB outflows is uncertain, we cannot determine whether the heavy nuclei synthesized in GRB outflows are dominated by Fe group elements ($A \\sim 40-60$) or whether the distribution extends to even heavier nuclei $(A \\gtrsim 90$). If the latter are present in at least a modest subset of events, a unique prediction of our model is that the UHECR composition may continue to increase to nuclei heavier than Fe at yet higher energies $\\gtrsim 10^{20}$ eV. Making a measurement of the composition that is sufficiently accurate to test this prediction will, however, require both better statistics and a better understanding of the hadronic physics used to interpret the air showers. Another consequence of an ultra-heavy composition is that the accessible distance of sources may be appreciably larger than for Fe or protons (see $\\S\\ref{sec:ebl}$ and eq.~[\\ref{eq:chi75}]). This alleviates energetic constraints on the GRB model for UHECRs, provided that the intergalactic magnetic field is not too strong ($\\S\\ref{sec:mag}$). Although heavy nuclei probably compose a substantial fraction of the mass $X_{\\rm h} \\sim 1$ in magnetized GRB outflows, the remainder is locked into $^{4}$He. Helium is easily disintegrated into protons by the EBL and hence could contribute to the proton flux near the ankle. Alternatively, the proton-rich composition measured near the ankle could be secondary particles produced by the disintegration of heavier nuclei by the EBL, or they could represent an entirely different source of UHECRs (Galactic or extra-galactic). If heavy nuclei from GRBs are indeed an important source of UHECRs, this would have several important consequences for GRB physics. For one, it would imply that GRBs outflows are magnetically-dominated, rather than thermally-driven fireballs (at least at their base). There is in fact growing evidence from {\\it Fermi} observations that GRB jets may be magnetically-dominated (e.g.~\\citealt{Zhang&Pe'er09}). If heavy nuclei are accelerated in GRB jets, this would also disfavor models in which GRBs are powered by heating from neutron-proton collisions (\\citealt{Beloborodov10}). The high densities in magnetically-driven outflows make in unlikely that free nuclei will avoid being captured into heavy nuclei." }, "1101/1101.3077_arXiv.txt": { "abstract": "I review some recent results about the molecular content of galaxies, obtained essentially from the CO lines, but also dense tracers, or the dust continuum emission. New results have been obtained on molecular cloud physics, and their efficiency to form stars, shedding light on the Kennicutt-Schmidt law as a function of surface density and galaxy type. Large progress has been made on galaxy at moderate and high redshifts, allowing to interprete the star formation history and star formation efficiency as a function of gas content, or galaxy evolution. In massive galaxies, the gas fraction was higher in the past, and galaxy disks were more unstable and more turbulent. ALMA observations will allow the study of more normal galaxies at high z with higher spatial resolution and sensitivity. ", "introduction": "A new survey (HERACLES) has been completed with the IRAM receiver array in the CO(2-1) line, allowing extended maps of nearby galaxies, at 12'' resolution (Leroy et al 2009). The survey contains an atlas of 18 nearby galaxies, observed at multi-wavelengths, and in particular in the HI line, and in the mid infrared by Spitzer. Among the results, it is interesting to note a very good correlation between CO and HI kinematics. The excitation of the molecular gas, as traced by the first two rotational lines of CO, is usually low in the disk (R = CO(2-1)/CO(1-0) = 0.6) while it is higher in nuclei (R=1), indicating denser gas. The CO emission is compatible with optically thick clouds at a kinetic temperature of T=10K. A more refined view of the star formation law in galaxies has been obtained by Bigiel et al (2008). The Schmidt-Kennicutt law relating star formation and gas density has a different slope $n$, according to the gas surface density. At high surface density, when the gas is molecular, the gas forms stars at a constant efficiency ($n$=1), and the time-scale for star formation is 2 10$^9$ yrs. While, at low surface density, when the gas is atomic, the slope is much higher $n$= 2 or more. At sub-kpc scale, the star formation rate is not strongly correlated with HI surface density. The transition between HI and H$_2$ occurs when the surface density is $>$ 9 M$_\\odot$pc$^{-2}$. In order to better determine the change across the spiral arms of the molecular gas physical properties, Schinnerer et al (2010) have mapped at interferometric resolution several lines of CO and its isotopes, together with dense gas tracers, such as HCN and HCO$^+$, in two selected regions across M51 spiral arms. They find no change across the arms, and the GMC population in the spiral arms of M51 is similar to those of the Milky Way, even if the star formation rate is much higher. On the contrary, the low surface brightness dwarf spiral M33 reveals different conditions than in the Milky Way. In the center, the lower metallicity is the cause of a higher conversion factor, as will be described by J.Braine in this conference (see Gratier et al 2010), while surprisingly the outer molecular complexes show exceptionally bright CO emission, and therefore a lower conversion factor (Bigiel et al 2010). The excitation of the molecular gas has also been investigated with the first 3 CO rotational lines, at high resolution with the PdB and SMA interferometers in two NUGA-sample galaxies (Boone et al 2010). At about 100pc-scale resolution, it is possible to distinguish warmer and less dense gas towards the center, may-be heated by the AGN, and colder and denser components in a circumnuclear arc. According to LTE analysis, more then 50\\% of the gas is optically thick in both galaxies. The CO excitation has recently ben estimated even more completely by determining the spectral line energy distribution or SLED in bright starbursts and quasars, with the Herschel Spire FTS instrument. It is possible to obtain the full spectrum at once, although with low spectral and spatial resolution, up to the CO(13-12) line. In Messier 82, the SLED obtained by Panuzzo et al (2010) reveals a significant part of the molecular gas at high temperature T=500K, where the H$_2$ lines are the main coolant. The peak of the CO line intensity occurs at the level $J$=7. Extremely high CO line excitation is a clue to the presence of an AGN and its strong X-ray heating (XDR). The star formation region are characterized by PDR, where dust is heated efficiently. Also there is a richer chemistry in XDR (H$_2$O, H$_2$O+, OH+..). Figure \\ref{fig1} reveals the CO SLED of Mrk231, a typical quasar, where the excitation is dominated by an XDR. \\begin{figure}[b] \\begin{center} \\includegraphics[width=9.2cm]{combes-f1.ps} \\caption{ Energy distribution in the various CO lines from Mrk231: the high frequency measurements (filled symbols) have been done with SPIRE on Herschel, while the low frequency ones (open symbols) are measured from the ground. The data can be reproduced by the combination of three models: two model PDR components (red and green lines) and an XDR component (blue line). The sum of these three components is the black line, made to fit the CO measurements. The PDR alone are not sufficient. From van der Werf et al (2010).} \\label{fig1} \\end{center} \\end{figure} For more common lower-energy AGN, the XDR is visible only very close to the nucleus. The ALMA spatial resolution is then required to resolve these regions. A typical example is the Seyfert 2 galaxy NGC 1068, where the XDR dominates the starburst regions only at r $<$ 70pc. Garcia-Burillo et al (2010) have mapped SiO, CN with the IRAM interferometer, which are tracers of shocks as well as CH$_3$OH, HNCO. The starbursts in barred galaxies are frequently found in nuclear rings, at Lindblad resonances, such as in NGC1097. The recent observation with Herschel of the dust morphology in this ring by Sandstrom et al (2010) has revealed a high uniformisation of the dust heating in the ring, suggesting some kind of smoothing, or that the ISRF is a significant source of heating. This smooth structure of the dust contrasts with the clumpy nature of the line emission, in the infrared lines: [OI] 63$\\mu$, [OIII] 88 $\\mu$, [NII] 122 $\\mu$, [CII] 158 $\\mu$ and [NII] 205 $\\mu$ (Beirao et al 2010). Galaxy formation models required strong feedback to limit star formation. This feedback is observed as molecular gas outflow out of starbursting region, such as the prototypical dwarf galaxy M82. Recently, gas outflows have been observed also generated by AGN feedback, in the nearby quasar Mrk 231, with OH lines and outflow velocities up to 1400 km/s (Fischer et al 2010). In the same quasar, the observation of the CO(1-0) line with the IRAM interferometer revealed broad wings, with velocities up to 750km/s, dragging the molecular gas in strong outflows (Feruglio et al 2010). According to this interpretation, most of the gas around the nucleus could be depleted in a time-scale of tens of Myrs. In Mrk 231, Gonzalez-Alfonso et al (2010) have reported about emission and absorption lines of H$_2$O, which is abundant due to shocks, or XDR chemistry, and evaporation of ice from grains. A clear example of outflows has also been seen in 4C 31.04 (Garcia-Burillo et al 2007). The HCO$^+$ profile is very wide, broader than 1000km/s. There is both emission and deep absorption in the blue-side, and an interferometer map has been able to map both, according to its spatial coincidence with the resolved jet in radio continuum. Another striking example, is the radio source 3C293, observed in HI absorption with a width of 1400km/s (Morganti et al 2003). The absorption is blue-shifted, indicating an outflow. Garcia-Burillo et al (2010) have mapped the source in CO and HCO$^+$, where a deep absorption is also detected. The strongest evidence of AGN feedback until now has been seen in the center of cool core clusters. Cold molecular gas has been detected in the CO lines, associated to the cooling flow in Perseus (Salome et al 2006). Recently, the cooling lines OI and CII have been mapped with PACS and SPIRE in some cooling flows (Edge et al 2010, Mittal et al, in prep). They were seen with the same morphology, and when possible with the same kinematics than for the CO lines. They appear to come from the same gas, cooling through different phases, showing no rotation, but in(-out)flows. A strong result from the first Science with Herschel is the evidence of a cold dust component in dwarfs and the outer parts of spiral galaxies (Grossi et al 2010). This has been confirmed with Planck, and also ground based bolometer, such as LABOCA. In dwarfs, the CII/CO ratio is very high, due to low-metallicity effects (Cormier et al 2010): the high UV environment, due to the lack of dust, provides large-scale photodissociation of the molecular gas. However, the SED of dust emission reveals and overabundance of cold dust, or a flat opacity law $\\beta <$1.5 in low-Z systems (Boselli et al 2010). This is the case in the LMC (Meixner et al 2010), and also M33 (Quintana-Lacaci et al. in prep). A 2-component grey body fit with $\\beta$=2 indicates a very cold component at 5.7K, 15 times more massive than the 21K component. This excess is clearly seen in the low-Z dwarf NGC 1705 (O'Halloran et al 2010). ", "conclusions": "" }, "1101/1101.5651_arXiv.txt": { "abstract": "The inclination of galaxies induces both reddening and extinction to their observed spectral energy distribution, which in turn impact the derived properties of the galaxies. Here we report a significant dependence of the error in photometric redshift (\\photoz) on the inclination of disk galaxies from the Sloan Digital Sky Survey. The bias in the \\photoz \\ based on the \\sdssphotozA \\ approach increases from \\biassdssonefaceon \\ in face-on to \\biassdssoneedgeon\\ in edge-on galaxies. A Principal Component Analysis on the full sample of photometry reveals the inclination of the galaxies to be represented by the \\photopcamodeinclination \\ mode. The corresponding eigenspectrum resembles an extinction curve. The isolation of the inclination effect in a low-order mode demonstrates the significant reddening induced on the observed colors, leading to the over-estimated \\photoz \\ in galaxies of high inclinations. We present approaches to correct the \\photoz \\ and the other properties of the disk galaxies against the inclination effect. ", "introduction": "The inclination of galaxies has been used as a tool to infer the opacity of disk galaxies \\citep[e.g., ][]{1958MeLu2.136....1H,1989MNRAS.239..939D,1990Natur.346..153V,1992MNRAS.254..677H,1993MNRAS.260..491D,1994AJ....107.2036G,1994A&A...283...12B,1995osd..conf.....D}. The effect of the inclination on disk galaxies are twofold: the reddening and the extinction on its spectral energy distribution, supported by many of the recent studies based on large samples of galaxies \\citep[e.g.,][]{2007MNRAS.379.1022D,2007ApJ...659.1159S,2008ApJ...681..225B,2008ApJ...687..976U,2008MNRAS.388.1321P,2009ApJ...691..394M,2010ApJ...709..780Y,2010ApJ...718..184C}. If these effects are not corrected for, one would expect an impact on the derived properties of the galaxies. One such property is the photometric redshift (\\photoz) of a galaxy, because it relies on the observed colors and magnitudes \\citep[e.g.,][]{1985AJ.....90..418K,1995AJ....110.2655C} of the galaxy. Many panoramic sky surveys will measure primarily broadband photometry of galaxies. Considering how the distance to a galaxy bares its influence from the inferred properties of the galaxy to the large scale structures in the universe, the correct estimation of the \\photoz \\ of galaxies is of utmost importance. Studies in cosmology are also impacted by the accuracy in the redshift of galaxies of various inclinations. Notably, \\citet{2010Natur.468..539M} have recently constrained dark energy content with statistics of the inclination of galaxies in pairs where the redshifts are known. We therefore explore and quantify in this \\paper \\ the dependence of the error in the \\photoz \\ on the inclination of disk galaxies. Among all of the Hubble morphological types, the geometry of disk galaxies deviates substantially from the spherical symmetry. One would expect a relatively large amplitude in any inclination-dependent effect. We present the sample of disk galaxies in \\S\\ref{section:sample}. We quantify the \\photoz \\ error as a function of the inclination of the galaxies in \\S\\ref{section:photozerror}. We present approaches to correct the \\photoz \\ and the other properties of the disk galaxies against the inclination effect in \\S\\ref{section:correction}. ", "conclusions": "The reddening in the spectral energy distribution of a disk galaxy caused by its inclination, if not taken into account, impacts the accuracy of the derived \\photoz. We present several approaches to correct the respective property of disk galaxies against the inclination effect. The considered properties are the restframe magnitudes, the flux densities of an arbitrary stellar population model for the disk galaxies, and the \\photoz. We evaluate the performance of the inclination-dependent color corrections by using the accuracy of \\photoz \\ as a diagnostics, and find that the corrections give statistically correct face-on colors of the disk galaxies. We identify the inclination of the disk galaxies to be represented by a low order PCA mode of the sample, namely the \\photopcamodeinclination \\ mode. The inclination therefore modulates significantly the variance in the photometric sample. By considering the first two eigenspectra, the variance is revealed to be related to the reddening effect on the spectral energy distribution. The reddening effect leads to the aforementioned large \\photoz \\ error." }, "1101/1101.0903_arXiv.txt": { "abstract": "We present a study of the stellar populations of ring and/or arm-like structures in a sample of S0 galaxies using {\\it GALEX} far- and near-ultraviolet imaging and SDSS optical data. Such structures are prominent in the UV and reveal recent star formation. We quantitatively characterize these rejuvenation events, estimating the average age and stellar mass of the ring structures, as well as of the entire galaxy. The mass fraction of the UV$-$bright rings is a few percent of the total galaxy mass, although the UV ring luminosity reaches 70\\% of the galaxy luminosity. The integrated colors of these S0s locates them in the red sequence (NGC 2962) and in the so-called green valley. We suggest that the star formation episodes may be induced by different triggering mechanisms, such as the inner secular evolution driven by bars, and interaction episodes. ", "introduction": "S0 galaxies have been introduced by \\cite{Hubble36} in his ``tuning fork\" galaxy classification as a more or less hypothetical transition class between ellipticals and spirals. The bulge and the disk are the defining structures of S0s. Later classification schemes by \\citet[][RC3 hereafter]{RC3} and \\citet[][RSA hereafter]{RSA} take into account the presence of the several sub-structures detected within the S0s class: bars, inner and outer rings, as well as lenses (from here the widely used term of lenticulars as synonym of S0s) and ovals are often found. \\begin{figure}% \\includegraphics[width=\\columnwidth]{Fig1Proc.eps} \\vspace{-0.2cm} \\caption{Composite {\\it GALEX} images (FUV=blue and NUV=yellow; left panels) and SDSS ($g$=blue, $r$=green, $i$=red; right panels) color composite images of the S0 galaxies NGC 1533, NGC 2962, NGC 2974, NGC 4245 and NGC 5636 and of the elliptical NGC 5638. The outer ring structures appear in the UV less smooth than in optical (in spite of the SDSS higher resolution) and bluer than the nucleus, suggesting the presence of young stellar populations} \\label{fig1} \\end{figure} S0s are, as a rule, admitted in the vast class of early-type galaxies (ETGs hereafter), together with Ellipticals, with which they share typically passively evolving stellar populations. From an evolutionary point of view it is widely believed that S0s were initially spirals which lost their interstellar medium (ISM hereafter) during collisions \\citep{Spitzer51} by ``harassment'' \\citep{Moore96} or by ``ram-pressure''. Simulations show that such events could occur both in clusters \\citep{Gunn72} and in groups \\citep[see e.g.][]{Bekki09}. From the kinematical point of view, genuine S0s are quite distinct from giant ellipticals since they are fast rotating, like late-type galaxies. At the same time, if S0s were originally spirals, they should have formed their mass over a significant fraction of the Hubble time, following a star formation history more similar to that of late-type galaxies rather than that of giant elliptical galaxies. In this context, our multi-wavelength study of nearby ETGs \\citep{Marino10a} shows that S0s, characterized through their luminosity profiles and by the low values of the $n$ index of the Sersic law \\citep{Sersic68}, have the lowest values of the $\\alpha-$enhancements ($[\\alpha/Fe]$) in the sample. The low $[\\alpha/Fe]$ values suggest a ``more prolonged'' star formation in the S0s galaxies with respect to ellipticals in the sample \\citep[][see also Rampazzo et al. these proceedings]{Annibali07,Annibali10}. Several mechanisms, both internal and external to the galaxy, may be envisaged to produce the signature of a ``prolonged'' star formation. The removal of the ISM from a possible spiral precursor may quench the ongoing star formation transforming the debris in an S0 with signatures of prolonged star formation, as well as an external ``wet'' accretion and/or the inner secular evolution, e.g. driven by a tumbling bar \\citep[see e.g.][]{vandenbosch98}. The {\\it Galaxy Evolution Explorer} satellite ({\\it GALEX} hereafter) has widely contributed to explore the above mechanisms. ``Wet'' accretions, and the rejuvenation of the stellar population, has been evidenced in ETGs \\citep[see e.g.][and Rampazzo et al. in these proceedings]{Marino09}. Recently, \\citet{Thilker10} found in NGC~404, a nearby well known S0, an external ring-like structure with signature of recent star formation. 70\\% of the FUV comes from an HI ring forming stars at a rate of 2.5$\\times$10$^{-3} M_{\\odot}$ yr$^{-1}$. The structure has been likely produced by a ``wet accretion/merger''. Along this line, using {\\it GALEX} we analyze five S0s, showing outer rings and/or arm-like structures, from the sample of \\cite{Marino10a, Marino10c} aiming at understanding the nature of the ring and to map possible star formation episodes. \\begin{figure*}[!t] \\centering \\includegraphics[width=12.2cm]{Fig2Proc.eps} \\caption{In the top right panel the optical and UV surface brightness profiles of NGC 5636 are plotted. Blue filled circles represent FUV, empty red circles NUV, and magenta empty squares SDSS$ r$. For comparison, the smooth FUV, NUV and $r$ surface brightness profiles of the elliptical NGC 5638 are shown in the top left panel. In the bottom panels the full dots represent the FUV, NUV, $u$, $g$, $r$, $i$, $z$ spectral energy distribution of these galaxies. Best fit models, assuming only foreground extinction, are indicated with continuous lines (see text). Ages, derived from SED fitting, are reported in Myrs} \\label{fig2} \\end{figure*} ", "conclusions": "" }, "1101/1101.0132_arXiv.txt": { "abstract": "The 3 mm wavelength spectra of 10 galaxies have been obtained at the Five College Radio Astronomy Observatory using a new, very broadband receiver and spectrometer, called the Redshift Search Receiver (RSR). The RSR has an instantaneous bandwidth of 37 GHz covering frequencies from 74 to 111 GHz, and has a spectral resolution of 31 MHz ($\\sim$100 km s$^{-1}$). During tests of the RSR on the FCRAO 14 m telescope the complete 3 mm spectra of the central regions of NGC 253, Maffei 2, NGC1068, IC 342, M82, NGC 3079, NGC 3690, NGC 4258, Arp 220 and NGC 6240 were obtained. Within the wavelength band covered by the RSR, 20 spectral lines from 14 different atomic and molecular species were detected. Based on simultaneous fits to the spectrum of each galaxy, a number of key molecular line ratios are derived. A simple model which assumes the emission arises from an ensemble of Milky Way-like Giant Nolecular Cloud cores can adequately fit the observed line ratios using molecular abundances based on Galactic molecular cloud cores. Variations seen in some line ratios, such as \\thco/HCN and HCO$^+$/HCN, can be explained if the mean density of the molecular gas varies from galaxy to galaxy. However, NGC 3690, NGC 4258 and NGC 6240 show very large HCO$^+$/HCN ratios and require significant abundance enhancement of HCO$^+$ over HCN, possible due to the proximity to active galactic nucleus activity. Finally, the mass of dense molecular gas is estimated and we infer that 25-85\\% of the total molecular gas in the central regions of these galaxies must have densities greater than 10$^4$ cm$^{-3}$. ", "introduction": "The molecular emission lines observed in galaxies are powerful probes of the physical and chemical properties of the gas most directly connected to the star formation process. Over the past several decades there have been many studies of the molecular line emission from galaxies, and these studies have provided invaluable information on the molecular gas content, star formation efficiency and molecular abundances \\citep{you91, omo07}. Although CO emission has been the primary means of deriving the molecular gas content in galaxies, emission from other less abundant molecular species, is much better in determining the physical and chemical properties of the gas. As the chemistry of the molecular gas is affected by the local radiation field, these trace molecular constituents may also provide information on local environments of this gas. More frequently observations of molecular species such as HCN, HCO$^+$, HNC, and CS are being obtained \\citep{omo07} and in a few cases mapped in nearby galaxies \\citep{meie05}. Since these molecules have much larger permanent electric dipole moments than CO, they require considerable higher densities to collisionally excite, and thus, in general, trace denser molecular gas than the gas probed by CO. Emissions from these dense gas tracers, such as HCN, have been argued to be a much better probe of the star-forming molecular gas than emission from CO \\citep{gao04}. The utility of emission from these `high dipole moment molecules' to probe the properties of the nuclear regions of galaxies is well established, and there have been numerous papers on this subject in recent years (see review by \\citet{omo07}). In addition, several emission line ratios, such as those for the low-lying transitions of HCN/CO, HCO$^+$/HCN, and HNC/HCN, have been suggested as good diagnostics of the properties of the dense star-forming gas in galaxies \\citep{aal08}. Nearly all studies of the molecular emission in galaxies have been limited to just a few targeted molecular lines. However, \\citet{mar06} recently presented the first spectral scan of a galaxy providing an inventory of the molecular lines in NGC 253 within the 2 mm wavelength band. Such spectral scans can provide a much more complete description of the chemical complexities of the molecular gas in galaxies. The 3 mm wavelength band is equally well suited, and past spectral surveys have provided important information on the chemical and physical properties of Giant Molecular Clouds (GMCs) in the Milky Way \\citep{joh84, joh85, cum86, tur89}. One major disadvantage of most spectral line surveys is that the data are assembled from many observations with varying pointing accuracy and with potentially systematic calibration problems. The observations described here have been obtained {\\it simultaneously} over the full spectral band for each galaxy, and hence many systematic problems are eliminated. In this paper we present the first 3 mm spectral scans of the central regions of 10 galaxies. ", "conclusions": "We have presented the 3 mm wavelength spectra of the central regions of ten relatively nearby galaxies. The spectra of each galaxy were obtained simultaneously with the RSR. The important advantages of obtaining the 3 mm wavelength spectra simultaneously are the perfect pointing registration in all spectral lines and the very good relative calibration of the spectral line intensities. The only uncertainty in relative line intensities comes from the small variations of beam efficiency and beam size with frequency. Most of the lines that we detected arise from high-dipole moment molecules, which allows us to examine the physical and chemical properties of the dense gas in these galaxies. Although the galaxies in our survey are far from being a homogeneous sample, their differences can be exploited to study the extent to which various line ratios are found to vary and to what extent they require differences in the physical or chemical properties of the molecular gas. The spectra were fit with a template with 33 spectral features to determine the line intensities. We detected 20 of the 33 spectral features in our template, and these features arose from 14 different atomic and molecular species. Based on the fitted line intensities, a number of line intensity ratios were examined. The \\thco/\\ceio\\ and HNC/HCN line ratios are insensitive to density. The \\thco/\\ceio\\ ratio was found to have a very narrow range of values, from 3.3 to 7.8, supporting the idea that both $^{13}$C and $^{18}$O are secondary nucleosynthesis products. The HNC/HCN ratio also was relatively constant from galaxy to galaxy, varying between 0.4 - 0.7 and consistent with the emission arising in relatively warm gas. The emission from these galaxies was modeled with an ensemble of molecular cloud cores, each with a temperature of 35 K and a molecular hydrogen column density per unit line width of 1$\\times10^{22}$ cm$^{-2}$ (km s$^{-1}$)$^{-1}$. The molecular abundances were taken from well-studied Galactic clouds, although these abundances are consistent with those found in NGC 253 \\citep{mar06}. The sensitivity of various line ratios to density and temperature was examined and it was found that density has the greatest affect. The HCN/\\thco\\ ratio is particularly density sensitive, and we believe that the large variation observed for this ratio may be due to differences in the mean gas density in these galaxies. Varying the density from 10$^4$ to 10$^6$ cm$^{-3}$ can explain the range of HCN/\\thco\\ ratios observed and provide consistent results with the other density sensitive line intensity ratios. Varying the density can also help explain much of the observed variations in the HCO$^+$/HCN ratio. Thus, a simple model with fixed molecular abundances and varying density, fits well the 3 mm spectra of most of these galaxies. One of the most notable exceptions is the anomalously large HCO$^+$/HCN ratios detected in NGC 3690, NGC 4258 and NGC 6240, which cannot be fit without increasing the abundance of HCO$^+$ relative to HCN relative. Based on the XDR model of \\citet{mei07}, a plausible explanation is that the AGN activity in these galaxies is responsible altering the chemistry in these galaxies and producing the large HCO$^+$/HCN ratio. The cloud core model was used to estimate the filling factor of dense, high column density gas in our beam and to estimate the mass of dense molecular gas. The derived beam filling factors are strongly correlated with distance, as our beam varies in linear size from about 0.8 Kpc in the nearest galaxies to over 17 Kpc in the most distant galaxies. In the nearby galaxies, beam filling factors as great as 0.26 are found, suggesting that the central regions of these galaxies have a high filling factor of dense, high column density gas. We compared our mass of dense gas with estimates based on CO and an assumed {\\it X}$_{CO}$-factor and find that the dense gas represents approximately 25-85 \\% of the total molecular gas in the central regions of these galaxies." }, "1101/1101.2906_arXiv.txt": { "abstract": "In \\cite{wis_etal10}, paper I, we analyzed 15 years of spectroscopic and spectropolarimetric data from the Ritter and Pine Bluff Observatories of 2 Be stars, 60 Cygni and $\\pi$ Aquarii, when a transition from Be to B star occurred. Here we anaylize the intrinsic polarization, where we observe loop-like structures caused by the rise and fall of the polarization Balmer Jump and continuum V-band polarization being mismatched temporaly with polarimetric outbursts. We also see polarization angle deviations from the mean, reported in paper I, which may be indicative of warps in the disk, blobs injected at an inclined orbit, or spiral density waves. We show our ongoing efforts to model time dependent behavior of the disk to constrain the phenomena, using 3D Monte Carlo radiative transfer codes. ", "introduction": "The time evolution of the intrinsic continuum V-band polarization (V-pol) of $\\pi$ Aqr is shown in Figure 1. We find evidence of clockwise loop-like structures (Figure 2) when comparing the evolution of the polarization across the Balmer Jump (BJ) vs V-pol, particularly during polarimetric outburst events (red in Figure 1). 60 Cyg also displays this behavior. \\begin{figure}[h]% \\centering \\parbox{0.32\\textwidth}{\\includegraphics[width=0.32\\textwidth]{s4-05_ZacharyHDraper_fig1.eps}\\caption{}}% \\begin{minipage}{0.32\\textwidth}% \\includegraphics[width=1.0\\textwidth]{s4-05_ZacharyHDraper_fig2.eps} \\caption{} \\end{minipage}% \\begin{minipage}{0.30\\textwidth}% \\includegraphics[width=1.0\\textwidth]{s4-05_ZacharyHDraper_fig3.eps} \\caption{} \\end{minipage}% \\label{fig:1}% \\end{figure} ", "conclusions": "" }, "1101/1101.4307_arXiv.txt": { "abstract": "With the help of 3D MHD simulations we investigate the collapse and fragmentation of rotating turbulent prestellar core embedded into turbulent medium. The numerical code is based on a high resolution Godunov-type finite-difference scheme. Initial turbulence is represented by the ensemble of Alfven waves with power law spectrum. Our computations show that under realistic parameters two bound fragments can appear when the density increases at $100-10^3$ times. The distance between the fragments is about $0.1$ of the initial core radius and their orbital period is comparable to the initial free fall time of the core. These results can explain the origin of binary stars with separation $0.001-0.01$ pc in the Galaxy field. ", "introduction": "Prestellar cores in molecular clouds are gravitationally bound condensations with stellar masses. The collapse and fragmentation of such cores lead to formation of single or multiple stars. We investigate the influence of magnetic field, rotation and turbulence on a prestellar core collapse and early fragmentation using 3D magnetohydrodynamical (MHD) simulations. Similar problem without turbulence was studied by \\cite[Machida et al. (2008)]{Machida08}, \\cite[Hennebelle \\& Teyssier (2008)]{Hennebelle08} and \\cite[Duffin \\& Pudritz (2009)]{Duffin09}. The role of turbulence was investigated by \\cite[Price \\& Bate (2008)]{PriceBate08} and \\cite[Wang et al. (2010)]{Wang10} for massive clouds without rotation. We carry out low resolution simulations (with cell number 128$^3$) and focus the attention on initial conditions and results analysis. The computational domain contains both the core and surrounding medium to describe a turbulent energy redistribution via MHD waves propagation (see \\cite[Heitsch et al. 2001]{Heitsch01}). The initial pulsations amplitudes depend on density (see \\cite[McKee \\& Zweibel 1995]{MZ95}) because a turbulence in non-uniform medium is inhomogeneous one. A bound fragment is found with the help of internal motion separation on inward, outward and tangential components relative to fragment mass center. ", "conclusions": "We carried out 3D MHD simulations of the collapse and fragmentation of rotating turbulent prestellar core embedded into turbulent medium. We conclude that: \\begin{itemize} \\item under realistic parameters only two bound fragments can appear when the core density increases at $100-1000$ times, \\item the distance between such fragments is about $0.1$ of the initial core radius, \\item their orbital period is comparable to the initial free fall time of the core, \\item these results can explain the origin of binary stars in the Galaxy field with semi-major axis in the range $0.001-0.01$ pc. \\end{itemize} At the next stage of the project we will use the AMR technology realized in the code Megalion. That will improve the resolution and lead to more interesting results. It is also desirable to develop the turbulent initial conditions without MHD discontinuities at the core surface." }, "1101/1101.4900_arXiv.txt": { "abstract": "The integrated galactic initial mass function (IGIMF) is computed from the combination of the stellar initial mass function (IMF) and the embedded cluster mass function, described by a power law with index $\\beta$. The result of the combination is a time-varying IMF which depends on the star formation rate. We applied the IGIMF formalism to a chemical evolution model for the solar neighbourhood and compared the results obtained by assuming three possible values for $\\beta$ with the ones obtained by means of a standard, well-tested, constant IMF. In general, a lower absolute value of $\\beta$ implies a flatter IGIMF, hence a larger number of massive stars, higher Type Ia and II supernova rates, higher mass ejection rates and higher [$\\alpha$/Fe] values at a given metallicity. Our suggested fiducial value for $\\beta$ is 2, since with this value we can account for most of the local observables. We discuss our results in a broader perspective, with some implications regarding the possible universality of the IMF and the importance of the star formation threshold. ", "introduction": "The initial stellar mass function (IMF) is of primary importance in galactic chemical evolution models. The IMF regulates the relative fractions of stars of different masses, hence their relative contribution to the chemical enrichment of the interstellar medium (ISM) is tightly related to this quantity. For this reason, the analysis of abundance ratios in galaxies may allow one to put robust constraints on both the normalization and the slope of the IMF (Recchi et al. 2009; Calura et al. 2010). \\\\ The Solar Neighbourhood (S. N. hereinafter) can be considered the most valuable environment to achieve constraints on the main parameters regulationg chemical evolution models, since it is definitely the best studied Galactic environment and many observational investigations devoted to its study provide us with a large set of observables against which models can be tested. These observables include diagrams of abundance ratios versus metallicity, which are particularly useful when they involve two elements synthesised by stars on different timescales. An example is the [$\\alpha$/Fe] vs [Fe/H] diagram, since $\\alpha$ elements are produced mostly by massive stars ($m>8 M_{\\odot}$) on very short ($<0.03$ Gyr) timescales, while type Ia supernovae (SNe) produce most of the Fe on timescales ranging from 0.03 Gyr up to one Hubble time (Matteucci 2001). This diagnostic is a strong function of the IMF, but depends also on the assumed star formation (SF) history (Matteucci 2001; Calura et al. 2009). Another fundamental constraint is the stellar metallicity distribution (SMD), which depends mainly on the IMF and on the infall history (hence on the star formation history) of the studied system. Another example of a useful diagnostic test for the IMF is the present-day mass function, i.e. the mass function of living stars observed now in the Solar Vicinity (Elmegreen \\& Scalo 2006). \\\\ The integrated galactic initial mass function (IGIMF) originates from the combination of the stellar IMF within each star cluster and of the embedded cluster mass function (CMF). It relies on the observational evidence that small clusters are more numerous in galaxies and that the most massive stars tend to form preferentially in massive clusters (Weidner \\& Kroupa 2006). The IGIMF is star-formation dependent, hence it is time-dependent and its evolution with time is sensitive to the star formation history of the environment. \\\\ In this paper, we use all the local observables to study the IGIMF and its effects on the chemical evolution of the solar neighbourhood. The results obtained with the IGIMF are compared to those obtained with a non-star-formation dependent (hence constant in time), fiducial IMF. The aim is to derive some contraints on the main unknown parameter of the IGIMF, i.e. the index $\\beta$ of the power law expressing the embedded CMF.\\\\ This paper is organized as follows. In Section 2 we present a description of the theoretical scenario behind the IGIMF. In Section 3 we describe the chemical evolution model for the Solar Neighbourhood. In Sect. 4 we present our results and finally in Sect. 5 some open problems regarding the IMF are discussed and some conclusions are drawn. \\begin{figure} \\plottwo{calura_f_1.ps}{calura_f_2.ps} \\caption{\\emph{Left panel:} IGIMFs for different cluster mass functions. Upper panel: $\\beta=1$; central panel: $\\beta=2$; lower panel: $\\beta=2.35$. In each panel we have considered 7 possible values of SFRs, ranging from $10^{-\u22124}$ $M_{\\odot} \\, yr^{-1}$ (lowermost solid lines) to 100 $M_{\\odot} \\, yr^{-1}$ (uppermost solid lines), equally spaced in logarithm. \\emph{Right panel:} the open circles are the ``fitness'' as a function of $\\beta$ for various models with different SF efficiencies (indicated by the numbers beside each open circle in units of Gyr$^{-1}$.). The horizontal line indicates the fitness value computed for the standard model. } \\label{fig1} \\end{figure} \\begin{table*} \\vspace{0cm} \\begin{flushleft} \\caption[]{Solar neighbourhood observables, parameters to which they are most sensititve and references.} \\begin{tabular}{l|l|l} \\noalign{\\smallskip} \\hline \\hline \\noalign{\\smallskip} Observable & Parameter & Reference \\\\ \\hline \\hline SFR Surface density & SF efficency & Rana (1991)\\\\ type Ia SNR & Integrated SF history, IMF & Cappellaro (1996)\\\\ type II SNR & SF efficiency, IMF & Cappellaro (1996)\\\\ Gas surface density & SF efficiency, IMF & Kulkarni \\& Heiles (1987) \\\\ & & Olling \\& Merrifield (2001) \\\\ Stellar surface density & SF history, IMF & Weber \\& de Boer (2009) \\\\ Stellar abundance ratios & SF history, IMF & various authors \\\\ Stellar Metallicity distribution & SF history, IMF & Jorgensen (2000) \\\\ Present-day mass function & SF history, IMF & Miller \\& Scalo (1979) \\\\ \\hline \\hline \\end{tabular} \\label{tab1} \\end{flushleft} \\end{table*} ", "conclusions": "In this paper, We have modelled the physical properties of the S. N. within the IGIMF theory. In this scenario, the IGIMF can be calculated by combining the cluster mass function with the stellar IMF, which represents the mass function of stars born within clusters and which can be described by a double-slope power law. An important feature of the IGIMF is that it depends on the star formation rate, which in turn evolves with time. The parameter $\\beta$ regulating the cluster mass function may have an important impact on the predicted properties of the Solar Neighbourhood. In general alower value for $\\beta$ corresponds to a flatter IGIMF. In terms of chemical evolution, a flatter IGIMF translates into higher mass ejection rates from dying stars, hence globally a lower mass fraction incorporated into stellar remnants and higher gas mass densities. This implies that the evolution of all the models computed assuming $\\beta=1$ and $\\beta=2$ are not sensitive to the star formation threshold and the star formation histories do not exhibit the ``gasping'' features typical of the standard model, which in turn is dominated by threshold effects at evolutionary times greater than 10 Gyr. Moreover, the lower the value of $\\beta$, the higher the SN rate, the higher the metallicity and the larger the $\\alpha$-enhancement visible in the abundance pattern. The statistical test used to compare model results and the obervables indicates that the model which best reproduces the local observables is carachterized by $\\beta=2$ as the index of the CMF. The results of the best model are very similar to those obtained with the standard case. A possible diagnostic which could help us disentangling between the two is represented by the present-day mass function (PDMF). The PDMF represents the mass function of living stars as observed in the solar neighbourhood. This quantity is an important diagnostic since it provides pieces of information complementary to the ones from the previously discussed observables. In the left panel of Fig.~\\ref{fig3}, we show the PDMF observed in the S. N. and predicted by means of our models. The PDMF computed with the standard IMF agrees with the observations in the range 0.4 $ M_{\\odot}$ - 2 $M_{\\odot}$. At very low stellar masses, the standard IMF seems too steep, whereas the distribution of stars with masses $>30 M_{\\odot}$ is underestimated. Once again, this is due to the SF threshold, which has strong effects on the SF history of the solar neighbourhood at late times, inhibiting recent SF and hence causing the underabundance or absence of very massive stars. In contrast, the models calculated with the IGIMF provide all similarly a very good fit to the observed PDMF. \\\\ The analysis of Fig.~\\ref{fig3} seems to suggest that the SF threshold should not play a dominant role in the late evolution of the S. N. Within the IGIMF theory, the existence of a SF threshold may be an observational selection effect, naturally explained in this context as shown by Pflamm-Altenburg et al. (these proceedings). However, as shown by chemical evolution results, the SF threshold is fundamental in reproducing the metallicity gradients observed in the MW and in local galaxies, unless a variable star formation efficiency through the disc is assumed (Colavitti et al. 2009). The study of the abundance gradients within the IGIMF theory may be of crucial help in sheding light on this issue and will be considered in future work. \\\\ Another important issue concerns the time evolution of the IGIMF. In the right panel of Fig.~\\ref{fig3}, we show how the IGIMF varies as a function of time in the case of the three best models computed with different values of $\\beta$. The best model ($\\beta=2$) shows very small variation of the IGIMF with cosmic time. Strong variations are predicted by assuming $\\beta=2.35$, since in this case the star formation history is very much influenced by the effects of the SF threshold. It may be interesting to test the effects of the IGIMF in spiral, Milky Way-like galaxies in a cosmological context. Cosmological semianalytical models predic strong variations in the star formation histories of spiral galaxies (Calura \\& Menci 2009), which present a large number of spikes due to merging events and which should manifest into strong variations of the IGIMF with redshift. \\\\ The universality of the IGIMF is another issue that deserves particular attention in the future. A chemical evolution study of elliptical galaxies within the IGIMF theory shows that the best value for $\\beta$ is 2.35, allowing to reproduce best the integrated $\\alpha/Fe$ ratios observed in the local early-type galaxies. This value is in contrast with the best value suggested by the analysis of the S. N. features. A further study of the IGIMF in local dwarf irregular galaxies and dwarf spheroidals could certainly be of some help in this regard. \\\\ Currently, the slope of the IGIMF in extreme SF conditions is another largely debated topic. Various indirect indications in external galaxies (Dabringhausen, these proceedings) and in our Galaxy (Stolte, these proceedings) seem to suggest that in strongly star forming systems, the slope of the stellar IMF should be flatter than the Salpeter one. Moreover, the assumption of a slightly top-heavy IMF in starbursts helps alleviating the discrepancy between cosmological models and observations regarding the $\\alpha/Fe$-$\\sigma$ relation observed in local ellipticals (Calura \\& Menci 2009). Addressing this subject within the IGIMF theory will be of primary importance in the nearest future." }, "1101/1101.4077_arXiv.txt": { "abstract": "Estimation of the black hole mass in bright X-ray sources of nearby galaxies is crucial to the understanding of these systems and their formation. However, the present allowed black hole mass range spans five order of magnitude ($10M_\\odot < M < 10^5 M_\\odot$) with the upper limit obtained from dynamical friction arguments. We show that the absence of a detectable optical counterpart for some of these sources, can provide a much more stringent upper limit. The argument is based only on the assumption that the outer regions of their accretion disks is a standard one. Moreover, such optically dark X-ray sources cannot be foreground stars or background active galactic nuclei, and hence must be accreting systems residing within their host galaxies. As a demonstration we search for candidates among the point-like X-ray sources detected with {\\it Chandra} in thirteen nearby elliptical galaxies. We use a novel technique to search for faint optical counterparts in the {\\it HST} images whereby we subtract the bright galaxy light based on isophotal modeling of the surface brightness. We show that for six sources with no detectable optical emission at the 3$-$sigma level, their black hole masses $M_{BH} < 5000M_\\odot$. In particular, an ultra-luminous X-ray source (ULX) in NGC~4486 has $M_{BH}< 1244 M_\\odot$. We discuss the potential of this method to provide stringent constraints on the black hole masses, and the implications on the physical nature of these sources. ", "introduction": "Compact, off-nuclear X-ray point sources in nearby galaxies, with luminosities $10^{39}-10^{41}\\rm~ergs~s^{-1}$ are referred to as Ultra-Luminous X-ray sources (ULXs). Detected in the early 1980's, with the {\\it Einstein} X-ray satellite \\citep{Fab89}, these objects were further studied with {\\it ROSAT} \\citep{Col99} and {\\it ASCA} \\citep{Mak00}. The {\\it XMM-Newton} and {\\it Chandra} X-ray observatories with their significantly higher angular resolution, dramatically confirmed the presence of ULXs \\citep{Kaa01}, and have enabled their spectral and temporal properties to be studied in detail \\citep[see ][for reviews]{Mil04,Mus04,Mus06,Rob07}. The observed luminosities of ULXs exceed the Eddington limit for a $10 M_{\\odot}$ black hole. Since ULX are off-nuclear sources, their masses must be $< 10^{5} M_{\\odot}$ from dynamical friction arguments \\citep{Kaa01}. Thus, ULX may represent a class of Intermediate Mass Black Holes (IMBHs) whose mass range ($10 M_{\\odot} < M < 10^{5}M_{\\odot}$) lies between that of stellar mass black holes and super-massive black holes observed in galaxy centers \\citep{Mak00}. Alternatively, ULX may be stellar mass black hole systems exhibiting super-Eddington accretions with their radiation geometrically beamed \\citep{Sha73,Kin08}. X-ray spectroscopy has provided supporting evidence in favor of IMBHs of $\\sim1000M_{\\odot}$ in ULXs \\citep{Mil03a,Mil03b,Cro04,Dew04,Rob05,Dev08}. Moreover, X-ray timing characteristics, i.e. presence of low frequency QPOs and/or breaks in the power density spectra, also suggests that ULX may harbor $\\sim 100$-$1000M_{\\odot}$ black holes \\citep{Str03,Dew06,Muc06,Str09}. While indicative these results are not conclusive, since there are also several arguments against IMBHs in ULXs \\citep[see e.g.,][]{Mus04,Rob07} and further investigations are required to reveal the true nature of these sources. Study of the host galaxy properties of ULXs reveals that their number and total X-ray luminosity is related to recent star formation activity, suggesting that they originate in young short-lived systems \\citep{Swa04,Swa09}. While the number of ULXs per galaxy is roughly the same for both spirals and ellipticals, the ones in the spirals have significantly higher luminosities \\citep{Swa04}. Optical counterparts have been reported for some ULX \\citep{Liu04,Kun05,Ram06,Ter06}. While some of the counterparts have been identified as O stars \\citep{Liu02,Liu07}, for most ULXs, the optical counterparts are stellar clusters \\citep{Goa02,Pta06}. However, for many ULX, the optical counterparts reveal that they are either background AGN \\citep{Gut06, Bon09} or foreground stars. ULXs found in elliptical galaxies may have contamination from background sources at $\\sim 44$\\% level \\citep{Swa04}. Detailed studies of X-ray sources in general and their connection with globular clusters have been undertaken \\citep{Kim06, Kim09} who note that the X-ray properties of the the sources in the field (i.e. without optical counterparts) are not different from those in globular clusters. The allowed black hole mass range for X-ray sources in nearby galaxies span five orders of magnitude ($10 M_{\\odot} < M < 10^{5}M_{\\odot}$) and it is important to obtain tighter constrains. Here, we show that the absence of a detectable optical emission allows us to impose an upper limit on the black hole mass for these accreting systems based on some standard assumptions. Moreoever, We argue that these {\\it optically dark} X-ray sources cannot be foreground stars or background AGN and hence are a true sample of sources located within the host galaxy. To demonstrate the technique, we search for candiatates among X-ray bright sources detected by {\\it Chandra} in archival {\\it HST} ACS, and WFPC2 images. \\begin{deluxetable} {lccc} \\tablewidth{0pt} \\tablecaption{Sample Galaxy Properties} \\tablehead{ \\colhead{Galaxy} & \\colhead{Distance (Mpc)} &\\colhead{$N_{x}$}&\\colhead{$N_{d}$}} \\startdata NGC 1399 & $18.3$ & $26$ & $4$ \\\\ NGC 4649 & $16.6$ & $12$ & $5$ \\\\ NGC 4697& $11.8$ & $11$ & $3$ \\\\ NGC 1291 & $8.9$ & $5$ & $1$ \\\\ NGC 4365 & $20.9$ & $4$ & $3$ \\\\ NGC 1316 & $17.0$ & $7$ & $3$ \\\\ NGC 4125 & $24.2$ & $3$ & $3$ \\\\ NGC 3379 & $11.1$ & $3$ & $1$ \\\\ NGC 4374 & $17.4$ & $2$ & $1$ \\\\ NGC 4486 & $15.8$ & $5$ & $2$ \\\\ NGC 4472 & $15.9$ & $1$ & $0$ \\\\ NGC 1407 & $17.6$ & $2$ & $2$ \\\\ NGC 4552 & $15.9$ & $3$ & $0$ \\\\ \\enddata \\tablecomments{(1) Host galaxy name; (2) Distance to the galaxy ; (3) Number of X-ray sources within HST field of view; (4) Number of X-ray sources without optical counterparts} \\label{Sample} \\end{deluxetable} \\begin{figure*} \\begin{center} \\subfloat[NGC 4486]{\\label{fig:NGC 4486}\\includegraphics[width=0.25\\textwidth]{f1a.eps}} \\subfloat[NGC 4697]{\\label{fig:NGC 4697}\\includegraphics[width=0.25\\textwidth]{f1b.eps}} \\subfloat[NGC 4649]{\\label{fig:NGC 4649}\\includegraphics[width=0.25\\textwidth]{f1c.eps}}\\\\ \\subfloat[NGC 4374]{\\label{fig:NGC 4374}\\includegraphics[width=0.25\\textwidth]{f1d.eps}} \\subfloat[NGC 1399]{\\label{fig:NGC 1399}\\includegraphics[width=0.25\\textwidth]{f1e.eps}} \\subfloat[NGC 1316]{\\label{fig:NGC 1316}\\includegraphics[width=0.25\\textwidth]{f1f.eps}} \\caption{The HST images of the first six X-ray sources (marked X) listed in Table \\ref{BHmass}. Overlaid are $1\\arcsec$ and $3\\arcsec$ circles centered on the shifted {\\it Chandra} positions. The plus signs mark the original {\\it Chandra} positions. Note that for four of the images, there are no optical sources even within $3\\arcsec$ of the X-ray position. } \\label{optimg} \\end{center} \\end{figure*} ", "conclusions": "Optically dark X-ray sources cannot be foreground stars or background AGN, otherwise their optical emission would be significantly higher than what is detected. Hence, these are a clean sample of sources within the host galaxies, which are probably accreting black hole systems. The optical emission from a standard accretion disk scales as mass of the black hole $M^{2/3}$ and hence the non-detection of optical emission imposes an upper limit on the black holes mass $M_U$. For ten of the sources $M_U < 10000 M_\\odot$. For a source in NGC 4486 with an X-ray luminosity clearly exceeding $10^{39}$ ergs/s (and therefore a bona-fide ULX by definition), the estimated black hole mass is smaller than $1244 M_\\odot$. This is two orders of magnitude smaller than the constraint obtained from dynamical friction, which is $ 10^5 M_\\odot$. These sources with black hole mass, $M_U < 5000 M_\\odot$ cannot be accreting systems with massive black holes residing in star clusters, or in the nuclei of merged satellite galaxies. For typical low-luminosity dwarf galaxies \\citep[$M_{B} \\sim -8.0$;][]{Mat98}, such an optical counterpart would be easily detected, given that our 3$-$sigma limits on the $HST$ images are much fainter. Even a compact nucleus of a merged dwarf galaxy \\citep[$M_{B} \\sim -7.0$;][]{Lot04} would have been easily identified. If they are binary systems, their companion cannot be a massive O star as such a star would have been detected in the optical image. Assuming an O star, with $M_{B} \\sim -5.5$, we find that the possibility of such a companion can be ruled out in all cases for which $M_U < 5000 M_\\odot$ (Table \\ref{BHmass}). In all of the above arguments, we have ignored the effect of dust obscuration in the host galaxies, because we are only considering elliptical galaxies in the present work. Although some ellipticals are known to have dust lanes, and nuclear rings in their centers, the X-ray sources we are considering here are distributed at fairly large radial distance from the center to be significantly affected by dust. Even for the best cases of optically dark X-ray sources presented in Table (\\ref{BHmass}), the range of black hole mass allowed is still large as the source could be a $\\simless 1000 M\\odot$ intermediate mass black hole, or a few solar mass object emitting at super-Eddington luminosities. The brightest X-ray sources in the sample have a luminosity of a few times $10^{39}$ ergs/s. This is unfortunate, since a dark source with luminosity $ > 10^{40}$ would have provided an order of magnitude better constraint on the black hole mass. Since the black hole mass upper limit $M_U \\propto D^3$ a bright X-ray source in a more nearby galaxy ($D \\sim 2 Mpc$) would have also provided significantly better constraints. A systematic search for such sources in very nearby galaxies may indeed prove fruitful. Another point to note is that bright X-ray sources in these galaxies are known to be variable in X-rays. Our analysis in this work, implicitly assumes that the X-ray luminosity observed through a single {\\it Chandra} observation, represents an average luminosity which is used to derive an average accretion rate $\\dot M = L_X/(\\eta c^2)$ which in turn is used to estimate the upper limit on the black hole mass (Eqn \\ref{M_U}). This assumption is required because the expected optical emission arises from the outer part of the disk and the local accretion rate there may be different than the one in the inner region which produces the X-rays. Any accretion rate fluctuation in the outer disk will be transfered along the disk on the viscous time-scale which could be significantly longer than a day. Thus, in principle one needs to ascertain the average X-ray luminosity of a source, and using a single very bright, but rare, X-ray observation of the source will not represent the average accretion rate. A systematic and comprehensive multi-wavelength study (using also other bands like infra-red and radio), along with X-ray variability studies, can shed further light on the nature of these sources." }, "1101/1101.4846_arXiv.txt": { "abstract": "Based on our recent MHD simulations, first, a concept of the successive merging of plasmoids and fragmentation in the current sheet in the standard flare model is presented. Then, using a 2.5-D electromagnetic particle-in-cell model with free boundary conditions, these processes were modelled on the kinetic level of plasma description. We recognized the plasmoids which mutually interacted and finally merged into one large plasmoid. Between interacting plasmoids further plasmoids and current sheets on smaller and smaller spatial scales were formed in agreement with the fragmentation found in MHD simulations. During interactions (merging - coalescences) of the plasmoids the electrons were very efficiently accelerated and heated. We found that after a series of such merging processes the electrons in some regions reached the energies relevant for the emission in the hard X-ray range. Considering these energetic electrons and assuming the plasma density 10$^{9}$--10$^{10}$ cm$^{-3}$ and the source volume as in the December 31, 2007 flare (Krucker at al. 2010), we computed the X-ray spectra as produced by the bremsstrahlung emission process. Comparing these spectra with observations, we think that these processes can explain the observed above-the-loop-top hard X-ray sources. Furthermore, we show that the process of a fragmentation between two merging plasmoids can generate the narrowband dm-spikes. Formulas for schematic fractal reconnection structures were derived. Finally, the results were discussed. ", "introduction": "It is commonly accepted that plasmoids play a very important role in the magnetic field reconnection in solar flares. Their importance, for the first time, was recognized by \\citet{OhyamaShibata1998}. In the 1992 October~5 flare, observed in soft X-rays by {\\it Yohkoh} satellite, they analyzed the plasmoid which was ejected during the impulsive phase upwards into the corona. Studying the same flare, \\citet{Kliemetal2000} showed that this plasmoid ejection was associated with the drifting pulsating structure (DPS) on radiowaves. They proposed the model of this radio emission, which was further developed in the papers by \\citet{Karlickyetal2002}, \\citet{Karlicky2004}, \\citet{KarlickyBarta2007}, \\citet{Bartaal2008a}, and \\citet{Karlickyetal2010}. In this model, in the current sheet, due to tearing and coalescence processes the plasmoids are formed. As shown by \\citet{Drakeetal2005, Drakeetal2006}, \\citet{Hoshino2005}, \\citet{Pritchett2006, Pritchett2008}, and \\citet{Karlicky2008} during these processes electrons are very efficiently accelerated. The electrons are then trapped in plasmoids, where they generate Langmuir waves, which through a wave transformation produce the electromagnetic waves recorded on the radio spectrum as DPSs. Due to limited range of plasma densities in the plasmoid the DPS is generated in the limited range of frequencies. In the vertical current sheet the plasmoids move upwards or downwards or even stay without any motion in dependance on a form of the surrounding magnetic field \\citep{Bartaal2008a, Bartaal2008b}. Due to a preference of divergent magnetic field lines in the upward direction, most of the plasmoids move upwards and corresponding DPSs drift towards lower frequencies. Nevertheless, in some cases the plasmoids move downwards and even interact with the underlying flare arcade as observed by \\citet{KolomanskiKarlicky2007} and \\citet{Milliganetal2010}. Recently, \\citet{Okaetal2010} studied the electron acceleration by multi-island coalescence processes in PIC model with periodic boundary conditions. They found that the most effective acceleration process is during the coalescence of plasmoids (\"anti-reconnection\"), see also \\citet{Pritchett2008} and \\citet{KarlickyBarta2007}. Furthermore, \\citet{ShibataTanuma2001} proposed that the current sheet, stretched by a rising magnetic rope, is fragmented to smaller and smaller plasmoids by the tearing mode instability in subsequently narrower and narrower current sheets (cascading reconnection). This suggestion has been recently further theoretically developed by \\citet{Loureiro+:2007} and \\citet{Uzdensky+:2010} into the theory of chain plasmoid instability. An advantage of this concept is that it explains how very narrow current sheets with high current densities (requested for the anomalous resistivity generation and fast reconnection) are generated. Moreover, many X-points in this model give sufficient volume for an acceleration of particles. Besides this fragmentation described by \\citet{ShibataTanuma2001}, \\citet{Bartaal2010b} found a new fragmentation in the region between two merging plasmoids using MHD simulations. This fragmentation is caused by the tearing mode instability in the current sheet generated between these interacting plasmoids and this process is repeated in smaller and smaller spatial scales. This fragmentation is driven by a merging process of the plasmoids. Considering all the above-mentioned processes, in the present paper, we focus our attention to two processes: (a) successive merging of plasmoids to large plasmoid and (b) fragmentation process between merging plasmoids. We selected these processes because we think that the successive merging of plasmoids can explain the above-the-loop-top hard X-ray source (as a large stationary plasmoid). On the other hand, the fragmentation can explain the narrowband dm-spikes. Because both these phenomena are generated by accelerated electrons, in following simulations we use the particle-in-cell (PIC) model instead of MHD models \\citep[e.g.][]{Bartaal2008a, Bartaal2008b}. The above-the-loop-top hard X-ray sources belong to the most discussed topics in recent years. The well-known example of such a hard X-ray source is that observed in the $\\sim$30--50 keV energy range by \\citet{Masudaetal1994}. However, such events are very rare \\citep{Tomczak2001, Petrosian2002, KruckerLin2008}. Another very interesting example was published just recently by \\citet{Kruckeretal2010}. They presented the hard X-ray source (with the energy up to $\\sim$80 keV) which was located 6 Mm above thermal flare loops. They derived the upper limit of the plasma density and source volume as n$_e \\sim$ 8$\\times$10$^{9}$ cm$^{-3}$ and $V \\sim$ 8$\\times$ 10$^{26}$ cm$^{3}$, respectively. Just a relatively low plasma density in such hard X-ray sources attracts attention of scientists. \\citet{Kruckeretal2010} concluded that these hard X-ray sources have to be close to the acceleration region and the distribution function of electrons emitting hard X-rays is strongly non-thermal or the plasma in the source is very hot (up to T$_e$ $\\sim$ 200 MK). Several ideas explaining these X-ray sources have been proposed, e.g. the magnetic or turbulent trapping, and dense (collisionally thick) coronal sources \\citep[see][]{Fletcher1995,Jakimiec1998,VeronigBrown2004,ParkFleishman2010} The narrowband dm-spikes (further spikes) belong to the most interesting radio bursts due to exceptionally high brightness temperatures (T$_{b}$~$\\approx$~10$^{15}$~K) and short durations \\citep[$\\leq$~0.1~s, see the review by][]{Benz1986}. Their observational characteristics were described in many papers \\citep[e.g.][]{Slottje1981,Karlicky1984,Fuetal1985,StahliMagun1986, Benzetal1982,ZlobecKarlicky1998,Meszarosova2003}. On the other hand, the theoretical models can be divided into two groups: a)~based on the plasma emission and acceleration processes \\citep{Kuijpersetal1981,Tajimaetal1990,Wentzel1991,BartaKarlicky2001}, and b)~based on the electron-cyclotron maser \\citep{Holmanetal1980,MelroseDulk1982,VlahosSharma1984,Wingleeetal1988, Aschwanden1990,FleishmanYastrebov1994}. To distinguish between these two types of models polarization and harmonic structures of the spikes have been also studied \\citep{Gudel1990,GudelZlobec1991,KruckerBenz1994}. Searching for a~characteristic bandwidth of individual spikes \\citet{Karlicky1996,Karlickyetal2000} found that the Fourier transform of the dynamic spectra of spikes have a~power-law form with power-law indices close to~$-5/3$. Based on these results \\citet{BartaKarlicky2001} proposed that the spikes are generated in turbulent reconnection outflows. This paper is organized as follows: First, we present our model scenario and explain the successive merging and fragmentation process. Then using a 2.5-D PIC model we simulate these processes. The results are then used in the interpretation of the above-the-loop-top hard X-ray sources and narrowband dm-spikes. \\begin{figure} \\begin{center} \\includegraphics[width=8.5cm]{fig1.eps} \\end{center} \\caption{Model scenario.} \\label{figure1} \\end{figure} ", "conclusions": "Important aspect of our model is that we used free boundary conditions which enabled successive merging of small plasmoids into a large final plasmoid. In simulations, we recognized also a fragmentation of current sheet between two merging plasmoids. We showed that these processes very efficiently accelerate electrons to energies relevant for the emission in the hard X-ray range. Based on this result we propose that the above-the-loop-top hard X-ray sources are produced by the successive merging of plasmoids in the region with the turbulent reconnection outflow, in the region just above the loop arcade. Computed X-ray spectra support this idea. To explain a difference between the slopes of the computed and observed X-ray spectra we propose that the observed power-law spectrum is a sum of emissions from many locations with different thermal and nonthermal distribution functions. In simulations we used the particle-in-cell model. Although, the studied processes are self-similar (i.e. they does not depend on scales, see e.g. the MHD and PIC simulations in the paper by \\citet{Tajimaetal1987}), the results need to be taken with a caution, because the spatial and time scales in the model and real plasmoids differ in several orders of magnitude. On the other hand, it is beyond a possibility of any present numerical model to take into account all these scales. Further aspect of the PIC modelling is that the Coulomb collisions were not considered in our model. Namely, the time interval of our computations is much shorter than the collision time in the-above-the-loop-top X-ray sources; e.g. $\\omega_{pe} t$ = 7800 for the plasma density n$_e$ = 10$^9$ cm$^{-3}$ corresponds to 4.5 $\\times$ 10$^{-6}$ s and the collision time to 0.115 T$_7^{3/2}$ s, where T$_7$ = T$_e$/10$^7$. On the other hand, the Coulomb collisions are essential for the bremsstrahlung X-ray emission. The hard X-ray spectra in the December 31, 2007 flare were detected as the mean ones over intervals of several seconds. Furthermore, the-above-the-loop-top hard X-ray source lasted of about 2 minutes. Thus comparing these times with the collision one, the collisions influence not only the observed spectra, but also a duration of the X-ray source. Therefore some re-acceleration of electrons is even needed. In the present model such a re-acceleration is very probable, because the plasmoids are generated in a broad range of spatial scales and electrons can travel several times though regions of interacting plasmoids. Although an inclusion of the Coulomb collisions and corresponding prolongation of computations are beyond the possibility of our present PIC modelling, we think that the observed spectra are given by a competition of the collisions with the re-acceleration of electrons. On the other hand, in PIC model the anomalous collisions (wave-particle interactions), which are much more effective than the Coulomb collisions, are present as confirmed by fast thermalization of nonthermal distribution functions. \\begin{figure} \\begin{center} \\includegraphics[width=8.5cm]{fig8.eps} \\end{center} \\caption{The radio spectrum observed during the 28 March 2001 flare by two Ond\\v{r}ejov radiospectrographs (0.8--2.0 and 2.0--4.5 GHz) supporting our model. It shows the drifting pulsating structures (DPSs) which drift towards narrowband dm-spikes. The 0.8--2.0 GHz spectrum was shortly interrupted at 12:08:34--12:08:36~UT.} \\label{figure8} \\end{figure} We made additional computations with different initial parameters and we found: a) the energy gain of accelerated electrons increases with the decrease of the plasma beta parameter, and b) the increase of the proton-electron mass ratio m$_p$/m$_e$ makes computations longer, but results are similar. We compared the present simulation also with that in the numerical model which size was two times smaller ($L_x \\times L_y$ = 600$\\Delta$ $\\times$ 2000$\\Delta$) and in which only 5 plasmoids were initiated (contrary to 10 plasmoids in the present simulation). In this case the final mean energy of accelerated electrons was 5.3 times greater than the initial one, compare with that of 10.7 times (from the initial temperature 10~MK to final temperature 107~MK) in the present case. Namely, each coalescence process increases the energy of accelerated electrons, therefore the number of successive coalescence processes is essential for their final energy. For calculations of the hard X-ray spectra, presented in Fig.~\\ref{figure5}b we used two methods. The obtained results are similar. Small differences are due to differences in these methods and deviations of the computed distribution functions from the thermal one. The plasmoids in 2-D are in reality 3-D magnetic ropes. While in 2-D the trapping of energetic electrons is a natural consequence of a close magnetic field structure of the plasmoid, in 3-D, this structure is only semi-closed. However, we consider the merging processes in the turbulent reconnection outflow therefore the magnetic trapping of electrons, similar to that proposed by \\citet{Jakimiec1998} is highly probable. Moreover, the coalescence fragmentation process, which generates the reverse electric currents (which in 3-D has to be closed in finite volume) will contribute to a full trapping of electrons. In agreement with the conclusions by \\citet{Kruckeretal2010}, in the model the acceleration region is very close to the hard X-ray source. It enables to re-accelerate energetic electrons, which loss their energy due to collisions. Acceleration regions are among interacting plasmoids and also between the plasmoids and the arcade of flaring loops. This model can explain not only the above-the-loop-top hard X-ray sources, but also the loop-top sources because the arcade of loops is, in principle, the 'plasmoid' fixed in its half height at the photosphere. Considering all aspects of the fragmentation process (power-law spatial scales of plasmoids, effective acceleration of electrons, trapping of electrons in plasmoids, location in the reconnection plasma outflow) we think that this process can explain a generation of the narrowband dm-spikes. We supported this idea by the radio spectrum observed during the 28 March 2001 showing DPSs drifting towards the narrowband dm-spikes. Furthermore, it is known that more than 70~$\\%$ of all groups of dm-spikes are observed during the GOES-rising-flare phases \\citep{Jirickaeta2001}. Although these arguments support the presented idea, further analysis of the narrowband dm-spikes and their modelling is necessary. In this first study, we considered only the neutral current sheet, i.e. $B_z=0$. For more realistic description, we plan to extend our study also to cases with non-zero guiding magnetic field. The presented model is a natural extension of our previous models explaining the plasmoid formation, its ejection, and corresponding DPS. The question arises why the above-the-loop-top hard X-ray sources are very rare comparing, e.g., with DPSs or dm-spikes observations. We think that the above-the-loop-top hard X-ray source is large and stationary plasmoid, which is sufficiently dense and in which there is sufficient amount of energetic electrons. Positional stationarity and location of this plasmoid is given by surrounding magnetic field and location where this plasmoid starts to be formed, see the paper by \\citet{Bartaal2008b}. On the other hand, the plasmoids generating DPSs or dm-spikes need not to be so dense and they need not to have such amount of energetic electrons. It is known that for generation of the radio emission a number of energetic electrons can be much smaller than that for the hard X-ray emission. Namely, the intensity of the radio emission depends on derivatives of the electron distribution function in the momentum space; not on the absolute amount of energetic electrons as in the case of the hard X-ray emission." }, "1101/1101.5592_arXiv.txt": { "abstract": "Loop quantum gravity and cosmology are reviewed with an emphasis on {\\em evaluating} the dynamics, rather than constructing it. The three crucial parts of such an analysis are (i) deriving effective equations, (ii) controlling the theory's microscopic degrees of freedom that lead to its spatial discreteness and refinement, and (iii) ensuring consistency and anomaly-freedom. All three issues are crucial for making the theory testable by conceptual and observational means, and they remain challenging. Throughout this review, the Hamiltonian nature of the theory will play a large role for properties of space-time structure within the framework discussed. ", "introduction": "In its different incarnations, quantum gravity must face a diverse set of fascinating problems and difficulties, a set of issues best seen as both challenges and opportunities. One of the main problems in canonical approaches, for instance, is the issue of anomalies in the gauge algebra underlying space-time covariance. Classically, the gauge generators, given by constraints, have weakly vanishing Poisson brackets with one another: they vanish when the constraints are satisfied. After quantization, the same behavior must be realized for commutators of quantum constraints (or for Poisson brackets of effective constraints), or else the theory becomes inconsistent due to gauge anomalies. If and how canonical quantum gravity can be obtained in an anomaly-free way is an important question, not yet convincingly addressed in full generality. Posing one of the main obstacles to a complete formulation of quantum gravity, this issue is hindering progress toward a detailed evaluation of quantum gravitational dynamics. A reliable phenomenological analysis must, after all, start with a consistent set of sufficiently general dynamical equations. But the strong and tough requirement of anomaly-freedom is also an opportunity, for it allows an analysis of quantum space-time and the changes in its structure possibly implied by quantum gravity. Addressing the anomaly problem is, moreover, crucial for an understanding of the dynamics of quantum gravity, both in the sense of {\\em constructing} consistent dynamical equations at the quantum level and in the sense of {\\em evaluating} equations and their solutions to bring out physical effects. Although the anomaly problem has not been addressed in full generality, several model systems have by now been analyzed in loop quantum cosmology, as reviewed by \\cite{LivRev}, with this question in mind. Loop quantum cosmology is a rather wide area within loop quantum gravity, analyzing several classes of model systems and perturbations around them. Loop quantum gravity, detailed by \\cite{Rov}, \\cite{ALRev} and \\cite{ThomasRev}, is a canonical quantization of general relativity based on holonomies (the eponymous loops) as elementary variables. The use of holonomies allows a background-independent formulation free of auxiliary metrics, and it implies several specific properties of the resulting dynamics. In all the systems used in loop quantum cosmology, quantization techniques close to those of a general loop quantization are used; they can thus be seen as capturing at least some of the crucial properties of full loop quantum gravity. To different degrees, most of these models make additional use of symmetry reduction as introduced by \\cite{SymmRed}, simplifying much of the quantum geometry and thereby providing rather direct access to the much less understood quantum dynamics. Thanks to these steps, implications of the quantum dynamics of loop quantum gravity, as generally defined based on \\cite{RS:Ham} and \\cite{QSDI}, have been evaluated quite explicitly for the first time. In general terms as well as for particular questions arising in loop quantum cosmology and loop quantum gravity, three key issues regarding quantum space-time arise, not unlike what one would expect for any fundamental gauge theory with microscopic degrees of freedom: \\begin{description} \\item[Effective dynamics:] In quantum gravity, geometry is described unsharply by whole states with all their fluctuations and correlations, in addition to the expectation values for an average geometry. Equations of motion for expectation values receive quantum corrections in their effective dynamics, as it may describe a quantum geometry. Along with this consequence of quantizing gravity come not only new mathematical space-time structures but also a vast enlargement of the number of degrees of freedom by quantum variables. The strongest control of such a high-dimensional dynamical quantum system is usually obtained for dynamical coherent states, defined as states saturating uncertainty relations at all times. If such states exist, they provide insights into the minimal deviations from classical behavior expected for a quantum system. The form and behavior of dynamical coherent states in loop quantum cosmology can be highlighted in several models, bringing out the role of space-time fluctuations and correlations as degrees of freedom beyond the classical ones. Exact dynamical coherent states exist only in special models and for specific initial values. Nevertheless, they allow interesting views on the generic quantum dynamics as it arises in quantum gravity. Before all corrections are derived for a large class of models, a clear analysis provided by dynamical coherent states, when they exist, unambiguously shows the first deviations from classical behavior. More generally, when exact dynamical coherent states do not exist, additional quantum corrections will result. They can often be computed perturbatively, analogously to loop corrections in interacting quantum field theories. \\item[Discrete dynamics:] In addition to those generic effects due to non-classical state parameters, underlying space-time structures typically change even for the expectation values of a quantum-gravity state. Most importantly, discrete geometry, at least in purely spatial terms which is by now well understood in loop quantum gravity following \\cite{AreaVol,Area,Vol2}, shows several detailed properties of importance for the dynamics and thus for space-time geometry. A spatial slice in space-time is equipped with a discrete quantum geometry, roughly seen as making space built from atomic patches of certain discrete sizes. One of the main problems of quantum dynamics is to show how these spatial atoms along slices fit together to form a quantum space-time --- or, more dynamically, how the spatial atoms change, merge, subdivide and interact as time is let loose. For an expanding universe, one would expect the discrete spatial structure to be refined as the volume increases; otherwise discrete sizes would be enlarged by huge factors, especially during inflation, making them macroscopic. The full dynamics of loop quantum gravity has indeed provided several hints that the number of discrete building blocks must change from slice to slice, once the dynamics is implemented consistently. This lattice refinement can be modelled even in the simplest, most highly symmetric situations of loop quantum cosmology, laid out by \\cite{InhomLattice,CosConst}. And it has shown several specific implications by which its precise form can already be constrained. \\item[Consistent dynamics:] Dynamics unfolds in time, but time is relative. Making sure that descriptions using different notions of time, corresponding to measurements by different observers, can agree about their physical insights requires the consistency conditions of general covariance. Since general covariance in gravity is implemented by gauge transformations, any quantization or even just a modification of the theory must, for consistency, respect this principle and be anomaly free. While the previous two points manifest themselves already in homogeneous models, where they can most easily be studied, the anomaly problem arises only in inhomogeneous situations. Within homogeneity, all but one of the gauge transformations underlying covariance are fixed. Anomalies can only arise if there are at least two independent gauge transformations; after quantization, they would be anomalous if their composition is no longer a gauge transformation. One could, of course, make it a gauge transformation by definition, by declaring the whole group generated by the independent gauge transformations as the gauge group. But if this group is too large, it would identify variables which are to be considered physically distinct, removing observables and degrees of freedom and in many cases leaving no non-trivial solutions. Changing the gauge transformations of a classical theory by quantum effects requires much care; only so-called consistent deformations of the classical gauge generators can provide well-defined quantizations. \\end{description} Addressing these questions is crucial, not only for a complete formulation of quantum gravity but also for reliable cosmological applications based on the resulting set of equations (such as singularity removal or structure formation). In what follows, we will describe the current status based on the models of loop quantum cosmology. ", "conclusions": "" }, "1101/1101.2467_arXiv.txt": { "abstract": "{} { We derive the evolution of the infrared luminosity function (LF) over the last 4/5ths of cosmic time, using deep 24 and 70~$\\mu$m imaging of the GOODS North and South fields. } { We use an extraction technique based on prior source positions at shorter wavelengths to build the 24 and 70 $\\mu$m source catalogs. The majority ($93\\%$) of the sources have a spectroscopic ($39\\%$) or a photometric redshift ($54\\%$) and, in our redshift range of interest (i.e., $1.31.3$, the SFR history of the Universe has been derived by several studies \\citep{perez_2005,caputi_2007} using deep 24 $\\mu$m imaging and infrared bolometric correction estimated from local spectral energy distribution (SED) libraries \\citep[][]{chary_2001,lagache_2003,dale_2002}. All of these studies concluded that the relative contribution of ULIRGs to the SFR density of the Universe increases with redshift, and may even be the dominant component at $z\\thicksim2$. However, these conclusions still need to be confirmed since there are large uncertainties at high-redshift in transforming observed 24~$\\mu$m flux densities to far-infrared luminosities \\citep{papovich_2007,daddi_2007}. To study the high-redshift evolution of the SFR density, one has to combine deep mid- and far-infrared observations in order to infer robust bolometric corrections and to clearly constrain the location of the break of the infrared luminosity function (LF). \\\\ \\\\} \\indent{ At $z\\thicksim2$, the observed 70 $\\mu$m emission corresponds approximately to the rest-frame 24 $\\mu$m luminosity, which was proven to be a good SFR estimator in the local Universe \\citep{calzetti_2007}. The reliability of this SFR estimator seems to hold at high-redshift since the SFR of $z\\thicksim2$ galaxies estimated from their observed 70~$\\mu$m flux densities and their radio emissions are in good agreement \\citep{daddi_2007a}. Thus, to get robust estimates of the SFR of distant galaxies, we decided to use deep 70~$\\mu$m observations obtained by \\textit{Spitzer}. \\\\} \\indent{ The main difficulty of using 70 $\\mu$m observations to study star-formation at $z\\thicksim2$ resides in the limited depth of the existing \\textit{Spitzer} data, even from the very deepest observations such as those in the Great Observatories Origins Deep Survey (GOODS). In this study, we overcome this limitation by using a stacking analysis. As shown in \\citet{papovich_2007} this analysis allows characterization of the 24 vs 70~$\\mu$m correlation and thus constrains the bolometric corrections to be applied to the 24~$\\mu$m flux densities. Using deep 24 and 70 $\\mu$m images of the GOODS-North and South fields we find that the 24 vs 70~$\\mu$m correlation observed at high redshift is significantly different from predictions by standard SED libraries. This deviation can be interpreted as a possible signature of an obscured active galactic nuclei (AGN) \\citep{daddi_2007} or simply as a SED evolution characterized by stronger polycyclic aromatic hydrocarbon (PAH) emission \\citep{papovich_2007}. Both interpretations are discussed and two different bolometric corrections are inferred. Based on these bolometric corrections we derive the infrared LF in two redshift bins (i.e., $1.31$ reinforces the idea that at this redshift we observe a change of the properties of star-forming galaxies as already suggested by the reversal of the star formation-density relation at $z\\thicksim1$ \\citep{elbaz_2007}. \\\\} \\indent{ The evolution of the SFR density of the Universe provides a strong constraint on the main mechanism which triggers the SFR in galaxies. At $z<1$ the decrease of the SFR density might be driven by a gradual gas exhaustion as suggested by the continuous decrease of the SFR vs M$_{\\star}$ relation in this redshift range \\citep{noeske_2007b,elbaz_2007}. Between $z \\sim 2$ and $z \\sim 1$, the relatively constant SFR density still needs to be understood in the framework of large-scale structure formation, merging, and/or AGN activity. \\\\} \\indent{ The main limitation of our study is the uncertainty on the influence of obscured AGN on the infrared bolometric correction to be applied to bright 24 $\\mu$m sources. This influence will soon be assessed using new far-infrared observations from \\textit{Herschel}. In particular, the GOODS-Herschel Open Time Key Programme (PI: David Elbaz) will reach the faintest flux limits at 100~$\\mu$m in an ultradeep field within GOODS-S, expected to provide individual measurements for most $z \\sim 2$ galaxies detected at 24~$\\mu$m by \\textit{Spitzer}, where here we could only derive average values based on 70~$\\mu$m stacking. This should help disentangle the contributions of AGN and star formation for sources over a broad swath of the high-redshift infrared luminosity function. }" }, "1101/1101.0181_arXiv.txt": { "abstract": "The 2010 outburst of the transient neutron star low-mass X-ray binary IGR J17480--2446 has exhibited a series of unique X-ray bursts, as well as millihertz (mHz) quasi-periodic oscillations (QPOs) related to these bursts. It has been recently proposed that these are type-II bursts, powered by the gravitational energy. This implies that the current nuclear-burning based model of mHz QPOs is not correct, and this timing feature cannot be used as a tool to measure the neutron star parameters. We report the analysis of the {\\it Rossi X-ray Timing Explorer} data of IGR J17480--2446 to show that the burst properties of this source are quite different from the properties of the type-II bursts observed from the rapid burster and GRO J1744--28. For example, the inferred ratio ($\\sim 50-90$) of the non-burst fluence to burst fluence is consistent with the thermonuclear origin of IGR J17480--2446 bursts, and is significantly different from this ratio ($\\lsim 4$) for type-II bursts. Our results suggest that the bursts and the mHz QPOs from IGR J17480--2446 are powered by the nuclear energy. ", "introduction": "The first observed outburst \\citep{Bordasetal2010} of the globular cluster Terzan 5 source IGR J17480--2446 in Oct/Nov, 2010 revealed that it is a neutron star low-mass X-ray binary (LMXB) and an 11 Hz pulsar \\citep{Chenevezetal2010, StrohmayerMarkwardt2010}. An X-ray burst was observed from this source with {\\it Rossi X-ray Timing Explorer} ({\\it RXTE}) on Oct 13, 2010 \\citep{StrohmayerMarkwardt2010}. While this is not surprising even for a pulsar, what made this source very interesting is the recurrent bursts observed for next few days \\citep{Altamiranoetal2010b, Papittoetal2011}. As the source intensity increased, the bursts became more frequent, and eventually disappeared, and millihertz (mHz) quasi-periodic oscillations (QPOs) appeared \\citep{Altamiranoetal2010b, Linaresetal2010}. Finally, the bursts reappeared, and as the source intensity decreased, they became less frequent \\citep{ChakrabortyBhattacharyya2010}. The burst recurrence frequencies observed between October 13 and 16 converge asymptotically towards the mHz QPO frequency, which suggests that these QPO is related to bursts \\citep{Linaresetal2010}. If these bursts are type-I or thermonuclear \\citep{StrohmayerBildsten2006}, then such a relation is expected from a widely accepted mHz QPO model, which interprets this timing feature as a signature of marginally stable nuclear burning \\citep{Hegeretal2007}. However, \\citet{GallowayZand2010} reported that, except the Oct 13 burst, no other burst up to Oct 26 showed a clear cooling trend during burst decay. This motivated them to suggest that, except the first burst, others are type-II bursts. The type-II bursts are believed to be caused by the accretion instability, and hence powered by the gravitational potential energy. This would mean that the mHz QPOs are related to the accretion dynamics, rather than nuclear burning, and hence plausibly could not be used as a tool to measure the neutron star parameters \\citep{Hegeretal2007, Bhattacharyya2010}. Type-II bursts have so far been observed from two sources: rapid burster and GRO J1744--28 \\citep{Tanetal1991, Lewinetal1995, Lewinetal1996}. \\citet{GallowayZand2010} have also indicated that IGR J17480--2446 might be a GRO J1744--28 analogue (since both have relatively wide orbit and slow pulsations), further strengthening the gravitational origin argument for the IGR J17480--2446 bursts. It is therefore extremely important to investigate whether the X-ray bursts from IGR J17480--2446 are powered by the nuclear energy or by the gravitational energy (1) in order to understand the nature of the unique source IGR J17480--2446, and its unique bursts, and (2) to find out the origin of mHz QPOs and their usefulness in understanding thermonuclear bursts and in measuring neutron star parameters. We have performed spectral and timing analysis of the entire RXTE Proportional Counter Array (PCA) data from IGR J17480--2446, the details of which will be reported elsewhere. In this Letter, we report and discuss the results required to address the single very important question mentioned above. ", "conclusions": "The gradual (as opposed to abrupt) change of all the burst properties throughout the outburst (Table~\\ref{Properties} and Fig.~\\ref{lc}) strongly indicate that all the bursts from IGR J17480--2446, including the Oct 13 burst, originate from the same physical mechanism. The points 3, 4 and 5 of \\S~\\ref{ResultsandDiscussion} suggest that these bursts are of thermonuclear origin, and mHz QPOs, plausibly caused by the quasi-stable nuclear burning, can be used as a tool to measure the neutron star parameters. The point 2 of \\S~\\ref{ResultsandDiscussion} indicates that a clear cooling may not always be present during a thermonuclear burst decay. The points 1, 3, 4, 5, 6 and 7 of \\S~\\ref{ResultsandDiscussion} show that IGR J17480--2446 is not a GRO J1744--28 analogue (\\S~\\ref{Introduction}). Finally, we note that the IGR J17480--2446 bursts are similar to the GS 1826-238 bursts in their clock-like recurrence times and in the shortening of the recurrence time with the non-burst flux increase \\citep{StrohmayerBildsten2006}.\\footnotemark[1]" }, "1101/1101.5137_arXiv.txt": { "abstract": "Two types of dust disks around white dwarfs (WDs) have been reported: small dust disks around cool metal-rich WDs consisting of tidally disrupted asteroids, and a large dust disk around the hot central WD of the Helix planetary nebula (PN) possibly produced by collisions among Kuiper Belt-like objects. To search for more dust disks of the latter type, we have conducted a {\\it Spitzer} MIPS 24 \\um\\ survey of 71 hot WDs or pre-WDs, among which 35 are central stars of PNe (CSPNs). Nine of these evolved stars are detected and their 24 \\um\\ flux densities are at least two orders of magnitude higher than their expected photospheric emission. Considering the bias against detection of distant objects, the 24 \\um\\ detection rate for the sample is $\\gtrsim$15\\%. It is striking that seven, or $\\sim$20\\%, of the WD and pre-WDs in known PNe exhibit 24 \\um\\ excesses, while two, or 5--6\\%, of the WDs not in PNe show 24 \\um\\ excesses and they have the lowest 24 \\um\\ flux densities. We have obtained follow-up {\\it Spitzer} IRS spectra for five objects. Four show clear continuum emission at 24 \\um, and one is overwhelmed by a bright neighboring star but still show a hint of continuum emission. In the cases of WD\\,0950+139 and CSPN K\\,1-22, a late-type companion is present, making it difficult to determine whether the excess 24 \\um\\ emission is associated with the WD or its red companion. High-resolution images in the mid-IR are needed to establish unambiguously the stars responsible for the 24 \\um\\ excesses. ", "introduction": "\\label{sec:intro} The {\\it Spitzer Space Telescope} \\citep{Wetal04}, with its superb sensitivity and resolution at infrared (IR) wavelengths, provides an excellent opportunity to study planetary debris disks around stars \\citep{Suetal06,Tetal08,Caetal09}. For instance, a comprehensive {\\it Spitzer} 24 \\um\\ survey of main-sequence A-type stars has shown that up to $\\sim$50\\% of young ($\\lesssim$30 Myr) stars have little or no 24 \\um\\ excess emission from debris disks, large debris-disk excesses decrease significantly at ages of $\\sim$150 Myr, and much of the dust detected may be generated episodically by collisions of large planetesimals \\citep{Rietal05}. The dust in these debris disks would have dissipated long before the stars evolve off the main sequence. Dust can be replenished during late evolutionary stages. As a low- or intermediate-mass star loses a significant fraction of its initial mass to become a white dwarf (WD), its planetary system expands. Sub-planetary objects, such as asteroids and comets, can be injected to very small radii and be tidally pulverized by the WD, while orbital resonances with giant planets can raise the collision rates among sub-planetary objects and generate dust \\citep{DS02}. This freshly produced dust can be detected through IR excesses and allows us to peer into the late evolution of planetary systems. The first two WDs reported to possess dust disks were G29-38 and GD\\,362, both exhibiting near-IR excesses that were confirmed spectroscopically to be dust continuum emission \\citep{ZB87, Betal05, Ketal05,Reetal05}. A subsequent {\\it Spitzer} survey of 124 WDs at 4.5 and 8.0 $\\mu$m found one additional dust disk, around WD\\,2115$-$560 \\citep{Metal07, vHetal07}. As G29-38, GD\\,362, and WD\\,2115$-$560 are WDs with photospheric absorption lines of heavy elements, searches for dust disks have been conducted for DAZ and DBZ WDs, and indeed many more dust disks were discovered. For example, a near-IR spectroscopic survey of 20 DAZ WDs found a dust disk around GD\\,56 \\citep{Ketal06}, the presence of a dust disk around the DAZ WD\\,1150$-$153 was diagnosed by $K$-band excess and confirmed spectroscopically \\citep{KR07}, and {\\it Spitzer} photometric observations of 9 DAZ/DBZ WDs revealed dust disks around GD\\,40, GD\\,133, and PG\\,1015+161 \\citep{Jetal07}. One common characteristic of these dust disks around DAZ/DBZ WDs is that they are all small, with outer radii $\\ll$0.01 AU. As the dust disks are completely within the Roche limits of the WDs and the dust mass is estimated to be only $\\sim$10$^{18}$ g, it is suggested that tidally disrupted asteroids produce the dust disks and that the accretion of this dust enriches refractory metals, such as Ca, Mg, Fe, and Ti, in the WD atmospheres \\citep{Jura03,Jetal07, Zetal07}. To date, $\\sim$20 dust disks around DAZ/DBZ WDs have been reported, all consistent with this suggested origin of tidal disruption of asteroids \\citep{Fetal10}. More dust disks of this type are being found from the {\\it Wide-Field Infrared Survey Explorer} observations of WDs \\citep[e.g.,][]{Detal11}. An entirely different kind of dust disk has been discovered around the central WD of the Helix planetary nebula (PN), WD\\,2226$-$210 \\citep{Suetal07}. {\\it Spitzer} observations of the Helix Nebula show a bright compact source coincident with the central WD in the 24 and 70 \\um\\ bands, and follow-up Infrared Spectrograph \\citep[IRS;][]{Hetal04} observations have verified that the mid-IR emission originates from a dust continuum. The spectral energy distribution (SED) of this IR-emitter indicates a blackbody temperature of 90-130 K, and its luminosity, 5-11$\\times$10$^{31}$ ergs s$^{-1}$, requires an emitting area of 3.8-38 AU$^2$. These properties can only be explained by the presence of a dust cloud; furthermore, little extinction exists toward the WD, so the dust cloud must be flattened with a disk geometry. Adopting a stellar effective temperature of 110,000 K for WD\\,2226$-$210 \\citep{Nap99} and assuming astronomical silicates with a power-law size distribution and a maximum grain radius of 1000 \\um\\ for the dust grains, models of the SED indicate that the dust disk extends between 35 and 150 AU from the WD and has a mass of $\\sim$0.13 $M_\\oplus$. Since this dust must have been generated recently and since the radial location of the dust disk corresponds to that of the Kuiper Belt in the Solar System, \\citet{Suetal07} suggest that the dust disk around WD\\,2226$-$210 was produced by collisions of Kuiper Belt-like objects (KBOs) or the break-up of comets. To simulate the dust disk of WD\\,2226$-$210, the dynamic evolution of a debris disk around a 3 $M_\\odot$ star has been modeled from the main sequence (corresponding to $\\sim$A0\\,V) to the WD stage, and it is found that collisions among KBOs may produce the amount of dust observed \\citep{BW10,Detal10}. If the dusk disk around the Helix central star is indeed produced by collisions of KBOs, similar dust disks should be found around other WDs and a survey would allow us to assess their frequency of occurrence and physical properties. These results can then be compared with models of debris disks evolution \\citep{BW10,Detal10} for implications on their planetary systems. To search for dust disks similar to that around the central WD of the Helix Nebula, we have conducted three surveys using {\\it Spitzer} observations: (1) 24 \\um\\ survey of hot WDs and pre-WDs (this paper), (2) archival survey of IR excesses of WDs \\citep{Cetal11,Retal11}, and (3) archival survey of IR excesses of central stars of PNe \\citep[CSPNs;][]{Betal11a}. The combined results from these three surveys will provide a comprehensive picture of post-main sequence dust production and dynamic evolution of debris disks. This paper reports the results of the {\\it Spitzer} 24 \\um\\ survey of hot WDs as well as follow-up spectroscopic and imaging observations for a subset of hot WDs with 24 \\um\\ excesses. Section 2 describes the target selection and observations, Section 3 reports the results, and Section 4 discusses the implications. A summary is given in Section 5. ", "conclusions": "\\label{sec:discussion} \\subsection{Statistical Properties} Among our sample of 71 hot WDs, nine show 24 \\um\\ excesses, corresponding to almost 13\\%. Figure 17 shows the distribution of the sample in $V$ and $J$. The number of detections in each magnitude bin is too small to provide meaningful statistics. If the sample is divided into a brighter group that has photometric measurements and a fainter group that has no photometric measurements, it can be seen that 24 \\um\\ excesses are detected in 15-16\\% of the brighter hot WDs, and only $\\sim$8\\% among the fainter hot WDs. As all of the WDs in our sample have high temperatures and a small range of radii, their brightnesses are indicative of distances, with the fainter ones being at larger distances. The different percentages of brighter and fainter WDs exhibiting 24 \\um\\ excesses most likely reflect the fact that the limiting sensitivity of our MIPS 24 \\um\\ survey precludes the detection of distant objects. The true percentage of hot WDs exhibiting 24 \\um\\ excesses is likely greater than 15-16\\%. For the nine hot WDs with 24 \\um\\ excesses, their 24 \\um\\ flux densities are plotted against their $J$-band flux densities in Figure 18. No correlations are seen in this plot. This is expected, as the $J$-band flux density is a rough indicator of distance and the 24 \\um\\ excess should not be dependent on distance. Another reason for the lack of correlation is the diverse physical conditions of the 24 \\um\\ emitters, as discussed later in Section 4.2. Seven of the nine hot WDs with 24 \\um\\ excesses are still surrounded by PNe: CSPN K\\,1-22, CSPN NGC\\,2438, WD\\,0103+732 in EGB\\,1, WD\\,0127+581 in Sh\\,2-188, WD\\,0439+466 in Sh\\,2-216, WD\\,0726+133 in Abell\\,21 (YM\\,29), and WD\\,0950+139 in EGB\\,6; while two are not in PNe: WD\\,0109+111 and WD\\,1342+443. There is a striking difference in the frequency of occurrence of 24 \\um\\ excesses between the WDs and pre-WDs in PNe and those without PNe: 20\\% and 5--6\\%, respectively. The two WDs not in known PNe are also the faintest in 24 \\um\\ among the nine. WDs without PNe are more evolved than those that are still surrounded by PNe. The smaller 24 \\um\\ excesses of the WDs without PNe appear to indicate an evolutionary trend of diminishing excess; however, this trend is not obvious among the 24 \\um\\ excesses of the WDs with PNe, if the nebular sizes \\citep{Cetal09} are indicative of their evolutionary status. Considering the diversity in the progenitors' masses and evolutionary paths of these WDs, the current sample of hot WDs with 24 \\um\\ excesses is too small to allow us to distinguish between an evolutionary effect and other effects. \\subsection{Nature of the 24 \\um\\ Excesses} The IR excesses of the CSPN Helix are detected in the 8, 24, and 70 \\um\\ bands, but not at shorter wavelengths \\citep{Suetal07}. High spatial resolution {\\it HST} observation has ruled out any resolved companion earlier than M5 \\citep{Cetal99}. Furthermore, the photometric accuracies in the IRAC bands (1-$\\sigma$ of 20, 24, 26 and 17 $\\mu$Jy at the 3.6, 4.5, 5.4, 8.0 $\\mu$m bands, respectively) can further constrain the mass of a possible companion. At an age of 1 Gyr, the 2MASS and IRAC photometry can also rule out, at the 3-$\\sigma$ level, any companion with mass greater than 20 Jupiter masses, i.e., early T dwarfs with $T_{\\rm eff} \\lesssim$650 K. Three of the nine new detections of mid-IR excesses of hot WDs show similar SEDs: WD\\,0103+732 (CSPN EGB\\,1) shows excess emission at 8 and 24 \\um, while WD\\,0439+466 (CSPN Sh\\,2-216) and WD\\,0726+133 (CSPN Abell 21) show excess emission at only 24 \\um. Their lack of near-IR excesses does not support the presence of late-type companions. Four of the hot WDs with 24 \\um\\ excesses also exhibit excess emission in the IRAC bands: CSPN K\\,1-22, CSPN NGC\\,2438, WD\\,0127+581 (CSPN Sh\\,2-188), and WD\\,0950+139 (CSPN EGB 6). IRS observations indicate that this IR excess is continuum in nature, and crude fits to the SEDs suggest blackbody emitter temperatures of 500--1200 K. These temperatures are in the ranges for brown dwarfs, but the emitting areas, $\\sim 1\\times10^{23}$ to $3\\times10^{25}$ cm$^2$, are too large for brown dwarfs. It is most likely that these IR excesses in the IRAC bands originate from hot dust emission. The two cases with cooler emitter temperatures (500-700 K), CSPN K\\,1-22 and WD\\,0950+139, are known to have late-type companions \\citep{Cetal99,Bond09}. The relationship between these companions and the hot dust components is uncertain. Future searches for companions of CSPN NGC\\,2438 and WD\\,0127+581 may help us understand the roles played by stellar companions in producing the dust component responsible for the IRAC excesses. Finally, two hot WDs with 24 \\um\\ excesses, WD\\,0109+111 and WD\\,1342+443, have no IRAC observations to determine whether excess emission is present in the IRAC bands; their 2MASS measurements do show any excess emission. In summary, among the ten hot WDs and pre-WDs that show excess emission at 24 \\um, 40\\% also show near-IR excesses associated with an additional warmer emission component that might be related to the presence of a companion. Whether a 24 \\um\\ excess is accompanied by near-IR excess or not, the shape of the SED suggests that the 24 \\um\\ emission originates from a component distinct from the WD's photospheric emission or another near-IR emitter. The 24 \\um\\ emission must originate from a source cooler than 300 K. Assuming that this source is heated solely by stellar radiation, we can determine the covering factor of the emitter from the luminosity ratio $L_{\\rm IR}/L_*$, where $L_{\\rm IR}$ is the luminosity of the excess emitter and $L_*$ is the luminosity of the WD. We adopt the stellar effective temperature, assume a blackbody model, and use the distance and extinction corrected photometry to determine the stellar luminosity $L_*$. (Useful physical parameters are listed in Table 6.) The luminosity of the excess emitter is calculated by assuming a 150 K blackbody model normalized at the observed 24 \\um\\ flux density. In the case of WD\\,0103+732 and WD\\,2226$-$210, blackbody temperatures are determined from model fits to the 8, 24, and 70 \\um\\ flux densities. For the four WDs exhibiting excess emission in the IRAC bands, we add another blackbody component to fit these measurements. These IR emitter models, illustrated in Figure 2, are used to calculate their approximate luminosities. The $L_{\\rm IR}/L_*$ ratios of the nine new cases and the Helix CSPN are listed in Table 6. The covering factors range from $4.9\\times10^{-6}$ to $4.7\\times10^{-4}$. A perfect absorbing body heated by a 100,000 K WD to 150 K would be at a distance of $\\sim$10 AU, and 100 K at 20 AU. Even at a distance of 10 AU, a covering factor of $4\\times10^{-6}$ corresponds to $5\\times10^{-3}$ AU$^2$, or $3\\times10^6$ $R_\\odot^2$, too large to be a star or a planet. The most likely origin of this mid-IR emitter is a dust disk, as proposed for the CSPN Helix Nebula \\citep{Suetal07}. The presence of near-IR excesses indicates the existence of a binary companion or a dust disk at higher temperatures, or both. Detailed modeling and high-resolution images are needed to decipher the true nature of the IR excesses. We defer the modeling of the SEDs and IRS spectra of hot WDs with 24 \\um\\ excesses to another paper (Bilikova et al.\\ 2011b, in preparation). The connection among CSPNs, circumstellar dust disks, and binarity has been alluded to by various observations and theoretical models. For example, Keplerian circumstellar dust disks have been observed to be associated with binary post-AGB stars \\citep{DRetal06}. Furthermore, the close binary stellar evolution of CSPNs has been suggested to play a very important role in the formation and shaping of PNe \\citep{Soker98,NB06,dM09}, although such binary companions are difficult to detect. Our detection of a warmer emitter (500--1200 K) in addition to the colder dust component (at 100--200 K) in four CSPNs, two of which have known companions, appear to further the connection. However, the dust disks responsible for 24 \\um\\ excesses are much larger than the circumstellar disks ejected from common-envelope binaries \\citep[a few AU at most;][]{TR10}, and have very different geometry from the Keplerian circumstellar dust disks around binary post-AGB stars, especially in the covering factor ($L_{\\rm IR}/L_*$). The Keplerian circumstellar dust disks around binary post-AGB stars typically have $L_{\\rm IR}/L_*$ $\\sim$ 0.2--0.5 \\citep{DRetal06}, several orders of magnitude higher than those responsible for 24 \\um\\ excesses of hot WDs and pre-WDs. This large discrepancy in covering factors suggests that these dust disks have different origins, and the small covering factors are more consistent with debris disks observed in main sequence stars \\citep{Suetal06,Tetal08,wyatt08}. It is thus likely that the origin of the 24 \\um\\ excesses of hot WDs and pre-WDs is dynamically rejuvenated debris disks as suggested by \\citet{Suetal07}." }, "1101/1101.1115_arXiv.txt": { "abstract": "We present a comprehensive study of the emission spectra from accreting sources. We use our new reflection code to compute the reflected spectra from an accretion disk illuminated by X-rays. This set of models covers different values of ionization parameter, solar iron abundance and photon index for the illuminating spectrum. These models also include the most complete and recent atomic data for the inner-shell of the iron and oxygen isonuclear sequences. We concentrate our analysis to the $2-10$~keV energy region, and in particular to the iron K-shell emission lines. We show the dependency of the equivalent width (EW) of the Fe K$\\alpha$ with the ionization parameter. The maximum value of the EW is $\\sim 800$~eV for models with log~$\\xi\\sim 1.5$, and decreases monotonically as $\\xi$ increases. For lower values of $\\xi$ the Fe K$\\alpha$ EW decreases to a minimum near log~$\\xi\\sim 0.8$. We produce simulated CCD observations based on our reflection models. For low ionized, reflection dominated cases, the $2-10$~keV energy region shows a very broad, curving continuum that cannot be represented by a simple power-law. We show that in addition to the Fe K-shell emission, there are other prominent features such as the Si and S L$\\alpha$ lines, a blend of Ar~{\\sc viii-xi} lines, and the Ca~{\\sc x} K$\\alpha$ line. In some cases the S~{\\sc xv} blends with the He-like Si RRC producing a broad feature that cannot be reproduced by a simple Gaussian profile. This could be used as a signature of reflection. ", "introduction": "Accreting systems are observed to emit copious radiation in the X-ray energy range which suggests emission from the innermost regions of an accretion disk. Analysis of the X-ray spectra is crucial to study the complex mixture of emitting and absorbing components in the circumnuclear regions of these systems. However, there are only few observables that provide indication of the existence of an accretion disk. These need to be understood in order to correctly interpret the physics of these systems. In a general sense, the observed radiation from accreting sources can be divided into: a thermal component, in the form of a black body emitted at the surface of the disk with typical temperatures of $\\sim 0.01-10$~keV (for the mass range $\\sim 10^8-10~M_{\\sun}$); a coronal component in the form of a power law covering energies up to $\\sim 100$~keV, believed to arise from inverse Compton scattering in a hot gas that lies above the disk; and a reflected component, resulting from the interaction of some of the coronal X-ray photons and the optically thick material of the disk. In the reflected component, the most prominent feature is the iron K$\\alpha$ emission line at $\\sim 6.4$~keV, produced by transitions of electrons between the $1s$ and $2p$ atomic orbitals. These are ubiquitous in the spectra of accreting sources \\citep{pou90, nan94,mil07}. Other reflection signatures are the so called Compton shoulder (next to the Fe K-line), and the Compton hump (above $\\sim 10$ keV), produced by the down-scattering of high energy photons by cold electrons. Much theoretical effort has gone into studies of X-ray illuminated disks over the past few decades. Most models assume that the gas density is constant with depth \\citep{don92,ros93,mat93,cze94,kro94,mag95,ros96,mat96,pou96,bla99}. Although constant density models may be appropriate for radiation-pressure dominated disks, other studies have shown significant differences when the gas density is properly solved via hydrostatic equilibrium \\citep{roz96,nay00, nay01, bal01, dum02, ros07}. Recently we have developed a new model for the reflected spectra from illuminated accretion disks called {\\sc xillver} \\citep{gar10}. Although our code is similar in its principal assumptions to previous models, {\\sc xillver} includes the most recent and complete atomic data for K-shell of all relevant ions \\citep{kal04,gar05,pal08a,gar09a,pal10}. This has a dramatic impact on the predicted spectra, in particular the K$\\alpha$ emission from iron. With this model we can study the effects of incident X-rays on the surface of the accretion disk by solving simultaneously the equations of radiative transfer and ionization equilibrium over a large range of column densities. Plane-parallel geometry and azimuthal symmetry are assumed, such that each calculation corresponds to an annular ring at a given distance from the source of X-rays. The redistribution of photons due to Compton down-scattering is included by using a Gaussian approximation for the Compton kernel. With {\\sc xillver} we are able to solve the reflection problem with great detail, i.e., with very high energy, spatial and angular resolution. In this paper we present a systematic analysis of our models for reflected spectra from X-ray illuminated accretion disks. We show how the most relevant atomic features in the spectra depend on the assumed properties of the irradiated gas. We pay particular attention to the Fe K-shell emission lines, and we quantify its strength in terms of the equivalent widths predicted by our models. These models are also used to produce faked CCD spectra in order to simulate the effects of the instrumental response and limited spectral resolution. These results will be helpful diagnostic tools in the interpretation of accreting sources observations. In the next Section we describe briefly the numerical methods used in our reflection code. In Section~\\ref{secres} we present the results of our analysis on the simulated spectra, as well as comparisons with observations from Seyfert galaxies and X-ray binaries. The main conclusions are presented in Section~\\ref{seccon} ", "conclusions": "\\label{seccon} In this paper we have presented a comprehensive study of the emission spectra from accreting sources. Using our new reflection modeling code \\citep{gar10}, we have computed the reflected spectra from an X-ray illuminated accretion disk. We have concentrated our analysis to the iron K-shell emission lines, although other spectral features are also discussed. Our models predict equivalent widths for the Fe K$\\alpha$ emission of $\\sim 10$~eV for high ionization parameters (log~$\\xi\\sim 4$), and maximum values of $\\sim 800$~eV for models with log~$\\xi\\sim 1.5$. For lower values of the ionization parameter the equivalent widths decrease to a minimum near log~$\\xi\\sim 1$, contrary to what other models predict. These differences are due to the atomic data used by each simulation. We have shown that the behavior of the Fe K$\\alpha$ equivalent widths with respect to the ionization parameter of the gas is consistent with the line emissivities shown by \\cite{kal04}, where the same atomic data was used. Additionally, these equivalent widths display a linear dependency with the iron abundance normalized to its solar values. This seems to be true for all the models within the range of ionization parameter values considered in this paper. Simple analysis of simulated CCD spectra reveals that for low ionized, reflection dominated cases the 2-10~keV continuum cannot be represented by a simple power-law. Instead, a broad Gaussian profile is also required. This type of continuum can be used as a strong reflection signature. These simulations also indicate that in addition to the iron K-shell lines, the S~{\\sc xv} K$\\alpha$ emission line at 2.45~keV is one of the most prominent features in the spectra. For cases with log~$\\xi\\sim 2$, this line blends with the He-like silicon RRC providing a broad feature that cannot be represented by simple Gaussian profiles. This could also be used as another reflection signature while analyzing observations. The iron K$\\alpha$ equivalent widths and line centroid energies are in good agreement with values reported in the literature for both AGN Seyfert galaxies, and galactic sources such as LMXB. In particular, both model and observations show that large equivalent widths ($> 200$~eV) can be achieved in many situations. According to our models, the largest values of the equivalent widths correspond to the lower $\\beta$ parameters, i.e., reflection dominated cases. Only observational data from AGN Sy 1.8-2 galaxies, or those with log~$N_{\\mathrm H} \\ge 23$~cm$^{-2}$ coincide with these large equivalent widths, but only for low values of the ionization parameter (log~$\\xi < 2$). The majority of the Sy 1-1.2, or those sources with log~$N_{\\mathrm H} < 21$~cm$^{-2}$ show equivalent widths below 200~eV. In order to reproduce those values with our models, we need to increase the contribution of the direct power-law up to ten times the contribution of the reflected component. These results agree with the general idea that type 1 and 2 Seyfert galaxies only differ on their orientation with respect to the observer. In the case of LMXB, both the model and the observed data coincide in equivalent widths of $\\sim 10-200$~eV. Since all these sources show hotter Fe K-lines (with energies around $6.7-6.9$~keV), they correspond to models with values of the ionization parameter larger than those for AGN ($2 < \\mathrm{log}~\\xi < 4$)." }, "1101/1101.1862_arXiv.txt": { "abstract": "We discuss and compare two alternative models for the two-point angular correlation function of galaxies detected through the sub-millimetre emission using the \\emph{Herschel} Space Observatory. The first, now-standard \\emph{Halo Model}, which represents the angular correlations as arising from one-halo and two-halo contributions, is flexible but complex and rather unwieldy. The second model is based on a much simpler approach: we incorporate a fitting function method to estimate the \\emph{matter} correlation function with approximate model of the bias inferred from the estimated redshift distribution to find the \\emph{galaxy} angular correlation function. We find that both models give a good account of the shape of the correlation functions obtained from published preliminary studies of the HerMES and H-ATLAS surveys performed using \\emph{Herschel}, and yield consistent estimates of the minimum halo mass within which the sub-millimetre galaxies must reside. We note also that both models predict an inflection in the correlation function at intermediate angular scales, so the presence of the feature in the measured correlation function does not unambiguously indicate the presence of intra-halo correlations. The primary barrier to more detailed interpretation of these clustering measurements lies in the substantial uncertainty surrounding the redshift distribution of the sources. ", "introduction": "\\label{SECintro} Current theoretical models predict that galaxies form and evolve in cold dark matter (CDM) halos. Galaxies consequently tend to trace the distribution of mass, although the manner in which they do this may be {\\em biased} \\citep{Kaiser1984,BBKS,Coles1993,Mo1996}. In principle, therefore, once the bias is allowed for, it is possible to use measurements of the clustering of galaxies to determine the clustering properties of the dark matter, especially if measurements can be made as a function of redshift. Because clustering evolution is sensitive to the parameters underlying the background cosmological model, such observations can provide another (independent) test of the concordance cosmological model; see, e.g. \\cite{Coles2005}. In addition, clustering observations can be used to constrain properties of the galaxies themselves, such as the minimum halo mass within which they reside, which may yield clues about the processes of galaxy formation and evolution. A steadily increasing number of surveys of large-scale galaxy clustering are now becoming available. In the optical wavebands there are projects such as the UKIDSS Ultra Deep Survey \\citep{Hartley2010} and the SDSS Redshift Survey \\citep{Connolly2010, Ross2009} which are being used to extract information on clustering as a function of redshift. The \\emph{Herschel} Space Observatory was launched in 2009 and is the only space observatory to cover a spectral range from the far infrared to sub-millimetre and therefore provides a new and unique window through which to study high-redshift galaxy clustering. Two surveys of particular interest to this article, HerMES \\citep{Oliver2010} and H-ATLAS \\citep{Eales2010}, have already released angular correlation results \\citep{Cooray2010,Maddox2010} from their Science Demonstration Phase (SDP) which we will discuss further later, so it is timely to raise the issue of modelling sub-millimetre galaxy clustering. The so-called Halo Model has been used extensively over the past few years in modelling galaxy clustering in a variety of contexts. The Halo Model combines approximations of the dark matter profile within individual halos, the mass function, and bias models to estimate the correlation function for given cosmological parameters and characteristic halo masses. It has been shown to provide accurate and reliable predictions of clustering measurements, at the price of a certain degree of complexity and modelling freedom. The main focus of this paper is a comparison of the Halo Model and a fitting function method which was initially introduced by \\cite{Hamilton1991}, and subsequently developed by \\cite{Peacock1994}, \\cite{JMW95} and especially \\cite{Peacock1996}, to estimate the matter correlation function. Work by \\cite{Matarrese97} followed by \\cite{Moscardini1998} and \\cite{Coles1998}, showed how to incorporate this idea into a technique for providing detailed predictions of angular correlation in high-redshift galaxy surveys. In this paper we compare the predictions of this much simpler approach with results from the Halo Model. The paper is organized as follows. In the next section, we outline the methodology for the fitting function method. Following that in Section \\ref{SECpoint2}, the models are compared to recent \\emph{Herschel} data and the best-fitting values of the free parameters are extracted and compared. In Section \\ref{SECcon}, we conclude with a summary. ", "conclusions": "\\label{SECcon} In this paper, we have compared models of the angular correlation function to data from the science demonstration phase of \\emph{Herschel}. We highlight a Fitting Function method which provides an improved fit to the data than a power law, and similar to that of the Halo model. It has just one free parameter, the minimum halo mass, compared to the two and three for a simple power-law and the halo model respectively. The halo mass is more meaningful physically as a parameter than those involved in the power law fit, so the Fitting Function is a much better method than the power-law for a quick-and-simple analysis. Although neither as sophisticated nor as flexible as the Halo model, it remains a useful tool that is perfectly adequate for modelling currently available data. For example, the minimum halo mass found using our Fitting Function model is consistent with that found using the Halo Model. Another point of interest is that, in fact, the linear angular correlation function also provides a reasonable fit to the data. Owing to the limited resolution of the \\emph{Herschel} telescope it is difficult to identify pairs of galaxies sufficiently close together on the sky to probe anything but the mildly non-linear regime. The currently available data provide some evidence of a transition between the linear and non-linear regimes, but they provide no unambiguous detection of the change of shape in $\\omega_\\textsuperscript{obs}$ that the Halo Model predicts. This does not mean the Halo Model is incorrect, of course, but what it does mean is that, at least for the time being, the paraphernalia involved in modelling intra-halo correlations is rather superfluous for these objects; simpler models can yield perfectly adequate results. Finally, we stress that the data sets to which we have applied these models are preliminary. The biggest stumbling-block to a more complete analysis relates to the considerable uncertainties in the redshift distribution $N(z)$ of the sources involved. The choices we adopted for this analysis are for illustration only so the results should not be regarded as definitive. Further data, especially ancillary data providing spectroscopic redshifts, will be needed before the precise nature of these galaxies can be determined." }, "1101/1101.3783_arXiv.txt": { "abstract": "Geometrically thick accretion flows may be present in black hole X-ray binaries observed in the low/hard state and in low-luminosity active galactic nuclei. Unlike in geometrically thin disks, the angular momentum axis in these sources is not expected to align with the black hole spin axis. We compute images from three-dimensional general relativistic magnetohydrodynamic simulations of misaligned (tilted) accretion flows using relativistic radiative transfer, and compare the estimated locations of the radiation edge with expectations from their aligned (untilted) counterparts. The radiation edge in the tilted simulations is independent of black hole spin for a tilt of $15^\\circ$, in stark contrast to the results for untilted simulations, which agree with the monotonic dependence on spin expected from thin accretion disk theory. Synthetic emission line profiles from the tilted simulations depend strongly on the observer's azimuth, and exhibit unique features such as broad ``blue wings.\" Coupled with precession, the azimuthal variation could generate time fluctuations in observed emission lines, which would be a clear ``signature'' of a tilted accretion flow. Finally, we evaluate the possibility that the observed low- and high-frequency quasi-periodic oscillations (QPOs) from black hole binaries could be produced by misaligned accretion flows. Although low-frequency QPOs from precessing, tilted disks remains a viable option, we find little evidence for significant power in our light curves in the frequency range of high-frequency QPOs. ", "introduction": "In standard thin disk accretion theory \\citep{shaksun1973,novthorne}, the angular momentum axis of the accretion flow is assumed to be aligned with the black hole spin axis. \\citet{bardeenpetterson1975} found that even if the initial angular momentum axis of the accretion flow is misaligned from the black hole spin axis, the inner part of the disk will still align on the viscous timescale. However, this so-called ``viscous\" regime only operates when $H/R \\lesssim \\alpha$, where $H/R$ is the scale height of the accretion disk, and $\\alpha$ is the parameterized viscosity \\citep{papaloizoulin1995}. This is applicable in active galactic nuclei (AGN) and the high/soft or thermal state of black hole X-ray binaries. On the other hand, advection-dominated accretion flows (ADAFs) are expected in the low/hard state of black hole X-ray binaries \\citep{narayanyi1995adaf,esinetal1997} and in low-luminosity AGN. ADAFs are unable to cool through efficient radiation, and are geometrically thick. It is likely that the accretion flow in many of these sources is misaligned, or ``tilted.'' Contemporary general relativistic MHD simulations \\citep[GRMHD, ][]{devilliers2003,gammie2003} currently provide the most physically realistic description of the inner portion of accretion flows around spinning black holes. Radiation can be calculated from these simulations in post-processing by assuming that it is dynamically and thermodynamically negligible. This method has been used to look for high frequency quasi-periodic oscillations (HFQPOs) in simulated data \\citep{schnittman2006} and to create radiative models of Sagittarius A* \\citep{noble2007,moscibrodzka2009,dexter2009,dexteretal2010}. All of this work assumed alignment between the angular momentum axis of the accretion flow and the black hole spin axis. \\citet{fragile2007,fragileetal2009,fragiletilt2009} were the first to do GRMHD simulations of disks with a tilt between these two axes. These new simulations yielded a number of unexpected features. First, the main body of the disk remained tilted with respect to the symmetry plane of the black hole; thus there was no indication of a Bardeen-Petterson effect in the disk at large. The torque of the black hole instead principally caused a global precession of the main disk body \\citep{fragile2005,fragile2007}. The time-steady structure of the disk was also warped, with latitude-dependent radial epicyclic motion driven by pressure gradients attributable to the warp \\citep{fragile2008}. The tilted disks also truncated at a larger radius than expected for an untilted disk. In fact, based on dynamical measures, the inner edge of these tilted disks was found to be independent of black hole spin \\citep{fragiletilt2009}, in sharp contrast to the expectation that accretion flows truncate at the marginally stable orbit of the black hole. Finally, \\citet{henisey2009} found evidence for trapped inertial waves in a simulation with a black spin $a=0.9$, producing excess power at a frequency $118 (M/10 M_\\sun)^{-1}$ Hz. In this work we use relativistic ray tracing to produce images and light curves of some of these numerically simulated tilted and untilted black-hole accretion disks. Our goal in this paper is to discuss observable differences between the two types of accretion flows, and to identify observational signatures of tilted black hole accretion disks. \\begin{figure} \\epsscale{1.2} \\plotone{f1.eps} \\caption{\\label{imgs}Sample images of the thermal emission model for the 90h (left) and 915h (right) simulations at $60^\\circ$ inclination. The observed photon energy is $E_0=10$keV for a $10 M_\\sun$ black hole, and each panel is $54 M$ across. The color scale is linear, increasing from blue to red to yellow to white.} \\end{figure} ", "conclusions": "Tilted accretion flows will inevitably be present in a significant fraction of black hole sources with $L/L_{edd} \\lesssim 0.05$ and possibly $L/L_{edd} \\gtrsim 0.3$ (thick or slim disks). Using relativistic ray tracing and a set of simple emissivities, we have compared the radiation edge, emission line profiles and power spectra of simulated black hole accretion flows with a tilt of $15^\\circ$ to their untilted counterparts. We find the radiation edge is independent of black hole spin, while the untilted simulations agreed with the expected qualitative trend of decreasing inner radius with increasing spin. These results for the radiation edge confirm the work of \\citet{fragiletilt2009}, who used dynamical measures to locate the inner edge. Due to the independence of inner edge on spin, the red wing of tilted accretion flow emission line profiles is also fairly independent of spin. This introduces a possible complication for attempts to measure black hole spin from sources which may be geometrically thick. In general, measurements of small spin (large inner radius) may be unreliable unless the disk is known to be untilted. A reliable estimate of a large black hole spin (small inner radius), in contrast, could rule out the presence of a tilted disk. The blue wing can be much broader for tilted accretion flows, and the tilted-disk line profiles depend strongly on the observer azimuth as well as inclination. Since a tilted disk is expected to precess \\citep{fragile2007}, highly variable emission line profiles could signify the presence of a tilted accretion flow, as pointed out by \\citet{hartnollblackman2000} for warped thin disks. Since many LLAGN and X-ray binaries in the low/hard state should be tilted, time-variable emission lines should be quite common, and this effect is unlikely to significantly depend on accurate reflection spectrum modeling. Although the simulations can only be run for a short time compared to the precession time scale, precession is a possible source of low frequency quasi-periodic oscillations when the accretion flow is optically thin due to the modulation of Doppler shifts as the velocities in the accretion flow align and misalign with the observer's line of sight \\citep[see][for more discussion of QPOs from precessing tilted disks]{ingram2009}. Finally, we have studied power spectra for our simple models. We find broken power law spectra with break frequencies around $100 M_{10}^{-1}$Hz and power law indices in the range 0-2 (3-4) pre- (post-) break for both tilted and untilted simulations. Previous studies \\citep{armitagereynolds2003,noblekrolik2009} found single power laws with index $\\sim$2. \\citet{armitagereynolds2003} found that power spectra from individual annuli are well described by broken power laws where the break frequency is close to the local orbital frequency -- the averaging of many annuli with an emissivity that falls with radius smooths the power spectrum into a single power law. We see the same behavior in our simulations; the break frequencies from power spectra of individual radial shells agree with the local orbital frequency for both simulations 90h and 915h. A break frequency $100 M_{10}^{-1}$Hz then implies a radius of $r \\approx 16$M. Our broken power law spectra are therefore likely due to the fact that our emissivity peaks relatively near the outer radius used for the ray tracing, $r=25$M. A larger radial domain would likely shift the break to smaller frequencies. Observed break frequencies in the low/hard state are typically $\\nu_b \\sim 0.1-1$Hz, which may be caused by the transition from a thin disk to a thicker, ADAF flow \\citep{esinetal1997}. That would imply a transition radius $r_t \\simeq 200-1000 M_{10}^{2/3} M$. Our results for pre- and post-break slopes from both tilted and untilted simulations agree with those found in Cygnus X-1 \\citep{revnivtsevetal2000} for an inclination $i=30^\\circ$. In GRO J1655-40 \\citep{remillardetal1999} our pre-break slopes agree for all inclinations. However, the PSD for that source is well described by a single power law. There is no clear evidence in our work for high frequency QPOs due to the trapped inertial waves identified by \\citet{henisey2009}, although there are more features in power spectra from the 915h simulation at higher significance than in 90h. Even when computing PSDs for sets of spherical shells from the simulations, there are no clear features in the tilted power spectra that are not also present in the untilted case. It is possible that this result could depend on the chosen emissivity. Alternatively, the excess power in trapped inertial waves could be insufficient to rise above the red noise continuum. The independence of the inner radius of the tilted simulations on black hole spin is attributable to the extra angular momentum transport provided by the asymmetric standing shocks. These shocks are only present in the tilted simulations. Their strength scales with black hole spin, which is a necessary condition for countering the greater centrifugal support at higher spins. The standing shocks, in turn, appear to be attributable to epicyclic motion within the disk driven by pressure gradients associated with the warped structure. Again, this effect scales with the spin of the black hole, which contributes to the stronger shocks. For small tilt angles, the orbital eccentricity scales as $e \\sim \\beta$. This suggests that significant deviations between the spin-dependence of the radiation edge and the marginally stable orbit should be present even at modest tilt angles $\\beta \\gtrsim 5^\\circ$. At larger tilts, it is unclear if the increasing eccentricity will lead to an inner edge that increases with spin. This is both due to the uncertainty in the radial tilt and twist profiles $\\beta(r)$ and $\\gamma(r)$ at larger tilts, and to the lack of a quantitative connection between inner disk edge and eccentricity. The dynamical measures from \\citet{fragiletilt2009} place the location of the inner edge in a simulation with $a=0.9M$ and $\\beta=10^\\circ$ closer to the location of 915h than 90h. This data point supports the idea that a noticeable departure between $r_{\\mathrm{edge}}$ and $r_ {\\mathrm{ms}}$ should exist between tilted and untilted disks even for $\\beta \\gtrsim 5^\\circ$. It also suggests that at larger tilt angles, $r_{\\mathrm{edge}}$ is likely to increase with spin unless the effect saturates at $\\beta \\approx 15^\\circ$. Simulations with larger tilt angles will be able to address this question with certainty." }, "1101/1101.1265_arXiv.txt": { "abstract": "The recent detection of $\\gamma$-ray emission from four radio-loud narrow-line Seyfert 1 galaxies suggests that the engine driving the AGN activity of these objects share some similarities with that of blazars, namely the presence of a $\\gamma$-ray emitting, variable, jet of plasma closely aligned to the line of sight. In this work we analyze the $\\gamma$-ray light curves of the four radio-loud narrow-line Seyfert 1 galaxies for which high-energy gamma-ray emission has been discovered by {\\it Fermi}/LAT, in order to study their variability. We find significant flux variability in all the sources. This allows us to exclude a starburst origin of the $\\gamma$-ray photons and confirms the presence of a relativistic jet. Furthermore we estimate the minimum {\\it e}-folding variability timescale (3 -- 30 days) and infer an upper limit for the size of the emitting region (0.2 -- 2 pc, assuming a relativistic Doppler factor $\\delta=10$ and a jet aperture of $\\theta=0.1$ rad). ", "introduction": "Narrow-line Seyfert 1 (NLS1) galaxies are a class of AGN characterized by a rather narrow width of the H$\\beta$ emission line (FWHM $\\la$ 2000 km/s), and by the presence of a strong FeII bump, a soft X-ray excess and flux ratio [OIII]/H$\\beta <$3 \\citep{1985-Osterbrock-spectra_of_nls1, 2000-Pogge-review_nls1}. These sources are usually radio-quiet, although, in a few cases, they show a radio-loudness parameter R (ratio between the 5 GHz and optical B flux densities, \\citealt{1989-Kellermann-def_radio_loudness}) greater than 1000 \\citep{2006-Komossa-rlnls1_quasar, 2008-Yuan-population_rlnls1_with_blazar_prop}. Such sources, dubbed radio-loud narrow-line Seyfert 1 galaxies (RL-NLS1), were thought to be inactive in $\\gamma$-rays, although several authors speculated the occurrence of similarities with blazars \\citep{2003-Zhou-pmnj0948, 2006-Komossa-rlnls1_quasar, 2007-Zhou-1h0323, 2008-Yuan-population_rlnls1_with_blazar_prop, 2009-Foschini-blazar_nuclei_in_rlnls1, 2010-Gu-CompactRadioNLS1}. The important discovery of $\\gamma$-ray emission from the RL-NLS1 source PMN J0948+0022 \\citep{2009-Foschini-discovery_pmnj0948, 2009-Abdo-discovery_pmnj0948} confirmed these similarities, i.e. the presence of a jet closely aligned to the line of sight as a source of Compton up-scattered $\\gamma$-ray photons. PMN J0948+0022 ($z$=0.585, \\citealt{refz-0948}) is one of the strongest radio source among the RL-NLS1; for its fast radio variability, source compactness, inverted radio spectrum and high brightness temperature, PMN J0948+0022 is one of the first RL-NLS1 for which the presence of a jet has been hypothesized \\citep{2003-Zhou-pmnj0948, 2006-Doi-VLBIObs0948}. The detection of $\\gamma$-ray emission definitively confirms the presence of a jet and allows to build a complete SED which closely resemble that of a typical blazar, with two broad peaks in the far IR and in the $\\gamma$-ray range respectively. By modeling the SED with the synchrotron and inverse-Compton model \\citep{2009-Ghisellini-Canonical_blazar} it is possible to roughly estimate important parameters such as the black hole mass, the Eddington ratio, and the power carried by the jet. The resulting values, although model-dependent, are usually compatible with estimates found independently in other studies. In the case of PMN J0948+0022 the black hole mass turns out to be $\\sim $1.5 $\\times$ 10$^8$ M$_{\\sun}$ (compatible with the range of values found by \\citealt{2003-Zhou-pmnj0948} using the empirical relation between the FWHM of emission lines and the continuum luminosity), and the Eddington ratio which is $\\sim$40\\%. A $\\gamma$-ray variability of a factor 2.2 rules out the possibility that the $\\gamma$-ray emission is due to a starburst contribution \\citep{2009-Abdo-mw_monitor_pmnj0948}. Shortly after this discovery, three more RL-NLS1 had been observed in $\\gamma$-rays \\citep{2009-Abdo-rlnls1_newclass}: 1H 0323+342, PKS 1502+036 and PKS 2004-447, thus confirming that at least some RL-NLS1 are $\\gamma$-ray emitting AGN. The SED of these sources has also been modeled with the synchrotron and inverse-Compton model, providing black hole mass estimates in agreement with previous studies \\citep[see][and references therein]{2009-Abdo-rlnls1_newclass}. 1H 0323+342 ($z$=0.06, \\citealt{refz-0323}) is considered to be a composite nucleus of a NLS1 (for its optical spectrum properties typical of NLS1) and a blazar (for its flat radio spectrum, core compactness and X-rays variability, \\citealt{2007-Zhou-1h0323}). It is the only $\\gamma$-ray detected RL-NLS1 for which we have an {\\it HST}/WFPC2 optical image of the host galaxy. The host shows a ring-like structure of 15.6 kpc in diameter and the entire galaxy looks like a one-armed spiral \\citep{2007-Zhou-1h0323}. Another study, based on NOT data, suggests that 1H 0323+342 may be the remnant of a galaxy merger \\citep{2008-Anton-colour1h0323}. The most striking feature of this source is the high luminosity of the accretion disc which, according to SED modeling, is approximately 90\\% of Eddington value \\citep{2009-Abdo-rlnls1_newclass}. PKS 1502+036 ($z$=0.41, \\citealt{refz-1502}) is the second most powerful source in terms of Eddington ratio (80 per cent). It carries a jet power, inferred from SED modeling, sligthly lower than PMN J0948+002 \\citep{2009-Abdo-rlnls1_newclass}. Finally, PKS 2004-447 ($z$=0.24, \\citealt{refz-2004}) is an unusual RL-NLS1. It is a strong radio emitter, with the highest radio-loudness parameter (R$>$6000) among the four $\\gamma$-ray detected RL-NLS1. PKS 2004-447 has an unusually weak Fe II complex compared to other NLS1 and a rather steep (although variable) radio spectral index compared to the other RL-NLS1. \\citet{2001-Oshlack-PKS2004} suggests that it may be well classified as a low-luminosity compact steep spectrum (CSS) radio quasar, also considering its low inferred black hole mass ($\\sim$10$^6$ M$_{\\sun}$) \\citep{2009-Abdo-rlnls1_newclass}. Spectral fit of the SED may require a thermal Comptonization component \\citep{2006-gallo-2004} or an external Compton component \\citep{2009-Abdo-rlnls1_newclass}. The detection of $\\gamma$-rays from the four mentioned RL-NLS1 allows to identify a new class of $\\gamma$-ray emitting AGN. The radio properties (namely temporal variability, flat spectrum and high brightness temperature) together with the $\\gamma$-ray detection suggests the presence of a relativistic jet closely aligned to the line of sight. Low power, mildly relativistic and poorly collimated radio jets have already been observed in a few spiral galaxies hosting Seyfert nuclei \\citep[e.g.][and references therein]{2006-Keel-0313-192}, but in the case of PMN J0948+0022 and PKS 1502+036 the power carried by the jet, as inferred from SED modeling, is in the range of quasars, while in 1H 0323+342 and PKS 2004-447 is in the range of BL Lac objects \\citep{2009-Abdo-rlnls1_newclass}. Such powerful jets are observed only in blazars hosted in elliptical galaxies \\citep{2009-Marscher} with black hole masses in the range 10$^8$ -\u2013 10$^9$ M$_{\\sun}$. By contrast black hole masses in $\\gamma$-ray detected RL-NLS1 (10$^6$ -\u2013 10$^8$ M$_{\\sun}$) are up to three orders of magnitude smaller than for blazars, thus suggesting that these sources emit at high Eddington ratios. Furthermore, it seems that the M$_{\\rm BH}$-$\\sigma_*$ scaling relation does not apply to NLS1, mainly due to their small FWHM of permitted lines. \\citet{2008-Decarli} showed that a reconciliation with the scaling relation is possible if the broad line region is assumed to have a disc-like geometry \\citep[but see][]{2008-Marconi}. While the $\\gamma$-ray variability of PMN J0948+0022 has already been studied in \\citet{2009-Abdo-mw_monitor_pmnj0948} and, with a higher level of significance, during the outburst occurred in July 2010 \\citep{2010-Foschini-outburst}, it has never been studied for the remaining three sources. Aim of this paper is to analyze the {\\it Fermi}/LAT light curves of the afore-mentioned RL-NLS1 galaxies in order to put the $\\gamma$-ray variability on a firm basis, and to find the minimum $\\gamma$-ray variability timescale. Throughout the paper, we assume a $\\Lambda$CDM cosmology with H$_0$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}$ = 0.27, $\\Omega_\\Lambda$ = 0.73. ", "conclusions": "In this work, we report the discovery of a statistically significant variability in the $\\gamma$-ray light curves of the four RL-NLS1 detected with {\\it Fermi}/LAT, with minimum variability timescales in the range 3 -- 30 days. This excludes a potential starburst origin of the $\\gamma$-ray emission, and supports the hypothesis of the presence of a jet closely aligned to the line of sight. A hint for photon index variations on timescales $\\sim$tens of days is also found in the data. Variability appears to be a feature common to the four first $\\gamma$-ray detected RL-NLS1. The minimum timescales found in the {\\it Fermi}/LAT energy range, appropriately scaled with the flux, are comparable to those found in the most luminous blazars. Thus, it is not possible to exclude variability as fast as that observed in blazars. This study goes in the same direction of the finding by \\citealt{2010-Foschini-outburst} who reported compelling evidence of similarities in the SED shape of PMN J0948+0022 (also analyzed here) and the archetypal blazar 3C 273. We are confident that in the near future more RL-NLS1 will be identified as $\\gamma$-ray emitters among the many unidentified $\\gamma$-ray sources observed by {\\it Fermi}/LAT." }, "1101/1101.5393_arXiv.txt": { "abstract": "We have searched for compact stellar structures within 17 tidal tails in 13 different interacting galaxies using \\textit{F606W}- and \\textit{F814W}- band images from the Wide Field Planetary Camera 2 (WFPC2) on the \\textit{Hubble Space Telescope} (\\textit{HST}). The sample of tidal tails includes a diverse population of optical properties, merging galaxy mass ratios, \\HI\\ content, and ages. Combining our tail sample with Knierman et al.\\ (2003), we find evidence of star clusters formed \\textit{in situ} with M$_V <$ -8.5 and $V$-$I$ $<$ 2.0 in 10 of 23 tidal tails; we are able to identify cluster candidates to M$_V$ = -6.5 in the closest tails. Three tails offer clear examples of ``beads on a string\" star formation morphology in $V$-$I$ color maps. Two tails present both tidal dwarf galaxy (TDG) candidates and cluster candidates. Statistical diagnostics indicate that clusters in tidal tails may be drawn from the same power-law luminosity functions (with logarithmic slopes $\\approx$ -2 -- -2.5) found in quiescent spiral galaxies and the interiors of interacting systems. We find that the tail regions with the largest number of observable clusters are relatively young ($\\lesssim$ 250 Myr old) and bright ($V \\lesssim$ 24 mag arcsec$^{-2}$), probably attributed to the strong bursts of star formation in interacting systems soon after periapse. Otherwise, we find no statistical difference between cluster-rich and cluster-poor tails in terms of many observable characteristics, though this analysis suffers from complex, unresolved gas dynamics and projection effects. ", "introduction": "Galaxy interactions and mergers are commonly observed phenomena, from intermediate (z $\\sim$ 1) to low redshifts. The complex gravitational potential of interacting galaxies drastically affects their morphologies, producing tidal tails and other disturbed features (\\citealp{toomre}; \\citealp{schweizer78}). Encounters between galaxies are also agents of photometric evolution, often localized a few kiloparsecs from the nucleus \\citep{schombert} or strewn across tidally distorted features (\\citealp{hibvan}; \\citealp{duc}; \\citealp{K07}; \\citealp{schombert}; \\citealp{hibbard05}). The intensity of these events may range from relatively rare global starbursts (e.g.\\ NGC 6240 and Arp 220), to more frequent local concentrations of star-forming behavior that have little impact on the integrated optical colors of the host galaxies \\citep{bergvall03}. Tidal tails are physically interesting environments for star formation. They are prevalent for $\\sim$ 500 Myr--1 Gyr from their inception at the galaxies' initial encounter, before dispersing into the intergalactic medium \\citep{binney08}. They can host alternating sequences of tidally compressive and extensive regions \\citep{renaud09} that may shape any observed characteristics of emerging cluster populations, and are bereft of the familiar periodic density waves of spiral arms that can trigger new generations of clusters. Moreover, the low stellar and gaseous masses and densities of tidal tails (e.g. \\citealp{elmegreen98}) compared to their progenitor galaxies ensures their star formation histories are easier to disentangle than the often heavily extincted, star-forming engines of interacting galaxy interiors. Tail internal extinctions are low (A$_V$ $\\lesssim$ 0.5; e.g.\\ \\citealp{temporin05}), simplifying the photometric analysis often made arduous and convoluted by intricate extinction maps and inconvenient galaxy orientations. The resulting low stellar densities present challenges in studying the sparse stellar populations within tidal debris. One way to circumvent this issue is to focus on the clustered stellar component of tidal tails. With the observational resources of \\textit{HST}, star clusters are luminous tracers of the tidal tail star formation history (\\citealp{iraklis09}; \\citealp{RdG09}) out to distances $\\approx$ 10--70 Mpc. Thus they allude to the overall star formation capacity of their tidal homes and the interplay of physical properties that allow such clusters to form and survive. Star clusters have been observed in tidal debris. For instance, populations of 40--50 compact clusters are found in the tidal tails of the Tadpole and Mice mergers (\\citealp{RdG}; \\citealp{tran}). \\citet{bastian05} also find clusters in the tails of NGC 6872, which appear to have luminosity and mass functions similar to those of clusters within the main bodies of mergers. It is difficult to assess the prevalence of star clusters across a variety of tidal tails, however, because studies have different resolution and completeness limits, and there are strongly age-dependent effects of cluster evolution (i.e.\\ fading). Furthermore, the connection between star formation in clusters and many observed features of galaxy mergers -- e.g.\\ age, multiwavelength brightness and colors, and dynamical properties -- is not straightforward. For instance, \\citet{K03}, hereafter K03, do not find a clear relationship between the number of star clusters they detect in tidal tails and various properties of a galaxy merger. They identify dozens of compact clusters brighter than M$_V$ = -8.5, with $V$-$I$ colors ranging from 0.2 to 0.9 in one system (NGC 3256), but do not find such an excess in three other galaxies (NGC 4038, NGC 3921, and NGC 7252). Their sample of interactions is limited to relatively old ($\\approx$ 400 -- 730 Myr), optically bright major mergers, however. It is uncertain what this implies for merging galaxies in the full spectrum of interaction ages, mass ratios, and other observable characteristics. Tidal tails do not appear to consistently have star clusters, unlike in the disks of gas-rich mergers (\\citealp{miller97}; \\citealp{whitmore99}; \\citealp{zepf99}). The single tidal tail with a confirmed cluster excess from K03 (NGC 3256W) has 0.11 $\\pm$ 0.03 cluster candidates kpc$^{-2}$ for sources with $V$-$I$ $<$ 0.7 and M$_{V} <$ -8.5. Additionally, clusters found in tidal debris often have ages similar to young clusters found within the interacting galaxies. In the eastern tail of NGC 3256, K03 infer young ages from their broadband colors, and \\citet{trancho} spectroscopically verify that several of these clusters are $\\lesssim$ 200 Myr old. These ages are also younger than the 400 Myr age of the tail they inhabit, strongly suggesting \\textit{in situ} formation. \\citet{peterson09} identify clusters in the tidal bridge of the NGC 7714/15 system, whose ages are similar to clusters they find in the interacting galaxies. Unlike these interior clusters, however, they do not find many bridge objects with masses above their age-averaged mass completeness limit of 10$^{4.8}$ \\msun. The dependence of clustered star formation on \\HI\\ properties is also not apparent. \\citet{aparna07} find that a threshold of \\HI\\ column density N$_{{\\mbox{\\scriptsize{\\HI}}}} \\sim$ 10$^{20.6}$ cm$^{-2}$ is a \\textit{necessary} but not \\textit{sufficient} condition for generating clusters at the M$_V <$ -8.5 level. It is true that projection effects may confuse the translation from column density to \\HI\\ volume density and pressure, but this questions how the complex interplay of \\HI\\ content, tail densities, pressures, etc.\\ shapes the star-forming environments of tidal tails. K03 also find that some systems devoid of star clusters seem to prefer harboring larger star- or cluster-forming complexes whose luminosities and sizes are consistent with tidal dwarf galaxies (TDGs). TDGs are potentially self-gravitating dwarf galaxy-sized accumulations of tidal debris (\\citealp{schweizer78}; \\citealp{bh92}; \\citealp{duc04}). NGC 7252 and NGC 3921 were presented in this context; they both show 7--8 clusters associated with their tidal dwarfs, and no statistical in-tail cluster excess anywhere else. The canonical example of a TDG is in NGC 4038/9 (\\citealp{saviane04}; \\citealp{hibbard01}), which shows evidence for star formation 2 Myr ago \\citep{mirabel92}, or hundreds of Myr after the main tail formed. Through numerical simulations, \\citet{hibbard95} and \\citet{bournaud06} assert that many TDGs may live long enough to become independent from their host tails. Through H$\\alpha$ Fabry-Perot observations and modelling, \\citet{bournaud04} find that the TDG candidate of NGC 7252 is a real condensation of matter and not a chance projection effect. \\citet{boquien10} explore the SEDs and star-forming properties of this TDG (and others) in depth. It is possible that the apparent mutual exclusivity of compact clusters vs.\\ TDGs identified in the K03 tails stems from small sample statistics, but it is also plausible that the outcome is affected by certain dynamical properties. For instance, the tail velocity dispersion and the resulting ambient pressure may conceivably encourage or discourage the formation of TDGs (\\citealp{elm93}; \\citealp{elm97}) or other bound structures. K03 posit that physical conditions in the tail may spur different modes of star formation, e.g.\\ either TDG or tail cluster formation, but not both. Modeling work by \\citet{duc04} suggests that variations in dark matter halo structure and other interaction parameters are key in deciding tail lengths and kinematic properties encouraging either star cluster formation along the tail or concentrated in kinematically distinct TDGs at the tail tip. This implies a top-down scenario where several gross tail and merger properties have the potential to help determine the cluster population in tidal debris. Observationally, this hypothesis is limited to a small interaction sample (K03), and requires follow-up with additional interactions. Fundamentally, there may be environmental factors that dictate how effectively stars and star clusters form in interacting pair of galaxies. Evidence now suggests that the majority of star formation occurs in clusters (\\citealp{lada}; \\citealp{fall}; although see \\citealp{bressert10}). While the formation efficiency of bound clusters appears to scale directly with the total star and cluster formation rate of its environment \\citep{silvavilla10}, star clusters (observed and inferred) typically constitute $\\lesssim$ 3--10\\% of the stellar complement of evolved galaxies. Thus, most clusters must fall prey to a variety of disruption mechanisms and contribute to the dominant field star component. It is unknown how processes determining star cluster formation and survival, and the luminous subset of these clusters we observe, relate to the physical conditions within tidal tails and the large-scale observable properties (luminosity, color, mass, TDGs, etc.) they prescribe. Cumulatively, these recent insights into the mechanics of clustered star formation and galaxy interactions merit a re-examination of clusters in tidal debris, specifically in terms of the characterization of these clusters and their relationships to their tail environments. In this paper, we extend the sample of merging galaxies of K03 with \\textit{HST} WFPC2 \\textit{F606W}- and \\textit{F814W}- band observations of twelve additional galaxies. The seventeen tidal tails of these otherwise isolated systems were selected to test environmental extremes for star formation across a broad range of interactions. An overview of the observations and data reduction is provided in \\S 2, along with object detection and photometry. In \\S 3, the properties of star cluster candidates of the tails are discussed, and \\S 4 follows with a discussion of star cluster populations in the context of the observed tidal tail environment. \\S 5 offers conclusions based on current data. Lastly, the appendices highlight observational details of the individual tidal tails and our efforts in optimizing the number of clusters detected at faint magnitudes across the sample. ", "conclusions": "We have presented a survey of the star cluster candidates in 17 tidal tails of 12 interacting galaxies observed with {\\it HST}/WFPC2 in \\Vhst\\ and \\Ihst. We combined this sample with the six tails studied by \\citet{K03}. The 23 tails span the vast parameter space of interaction characteristics like ages, mass ratios, average \\HI\\ surface densities, \\HI\\ richness, tail lengths and surface brightnesses, and host galaxy star formation rates. Star cluster candidate (SCC) densities in the tail regions were computed by subtracting a background estimated from out-of-tail sources. A significant excess ($>$ 2.5$\\sigma$) was detected at M$_V <-8.5$ and $V$-$I$ $<$ 2.0 for 10 of 23 tails. This color and magnitude range was selected to ensure completeness, exclude single stellar contaminants, and allow for stochastic and systematic effects of cluster evolution. We have excluded tails whose SCCs were largely constrained to tidal dwarf galaxy (TDG) candidates. In some cases cluster formation seems confined to these structures, as was seen in \\citet{K03}. We also identify several cases with definite in-tail SCC excesses and TDG candidates, but cannot extrapolate the significance of this with the paucity of multiwavelength and kinematic data for several of these systems. We also find three cases of ``beads on a string\" morphology (with separations of $\\approx$ 3 kpc) for clumps of sources/SCCs, characteristic of large-scale star formation by gravitational instabilities also seen in spiral galaxies and tidal tails of other studies. For some nearer tails, fainter sources could be detected, and in the NGC 520 TDG candidate a significant cluster excess was found for M$_V <-7.5$ and M$_V<-6.5$, where none was detected at M$_V < -8.5$. Given the limitations of our exposure times and broadband coverage for this project, it cannot be conclusively determined if the in-tail cluster candidate excesses at fainter magnitudes are significant for other tails, or if more distant tails have faint clusters that are currently undetectable. To the extent that we can determine them with small numbers, the cluster luminosity function in the tail environment may follow the same power-law slope as in quiescent and starburst galaxies. Combined with the emergent consensus on the universality of clustered star formation, we therefore contend that cluster populations are likely to exist to some level in most, if not all tails. This level may not always be detectable by \\textit{HST}. A complex combination of various factors, including environmental differences in star formation, gas supply, tail ages, and projection effects appear to influence the populations we observe. Of the global properties we considered (interaction age, progenitor mass ratios and star formation rates, and total \\HI\\ tail masses), no parameter was individually responsible in shaping the cluster populations. However, we do find statistical differences in the populations or tails with and without cluster excesses in terms of $V$-band surface brightness and mean \\HI\\ density. The tails that have the highest excesses of star cluster candidates at M$_V <$ -8.5 tend to be the youngest ($\\lesssim$ 250 Myr), probably tracing the strongest episodes of star formation that are typically triggered close to periapse. Young tails also tend to have larger mean observed \\HI\\ surface densities than older tails. It is important to note, however, that there are exceptions to all of these trends. We conclude that there is no single explanation for a tail to have a cluster excess, though there are several important influential factors like local gas density and local star formation activity at the time of observation. Quantities like magnitude of the brightest observed SCC and the overall tail surface brightness appear to loosely trace SCC excess. Merger dynamics not addressed here (e.g.\\ encounter speeds and orientations; dark matter halo structure) may also be important in imprinting the resulting tidal field with a dynamical setup conducive to star and cluster formation. To disentangle the full assortment of interaction variables would be aided by a larger, deeper sample, more accurate cluster age-dating, and by a higher resolution, local study of the relationship between gas and star cluster formation." }, "1101/1101.0864_arXiv.txt": { "abstract": " ", "introduction": "Since the pioneering work by Paczynski \\cite{Paczynski} and the subsequent detections of MACHOs \\cite{Alcock,Aubourg}, microlensing has been one of the most vital areas in astrophysics. Now, it plays an important role also in searching extra-solar planets \\cite{Beaulieu,Gaudi}. Light rays are influenced by a local curvature of a spacetime \\cite{SEF}. It seems natural that the bending angle of light rays in modified gravity theories is different from that in general relativity. Does such a modified bending angle really makes a change in microlensing? We shall examine this problem in this paper. Conventional studies of microlensing are usually based on the lens equation that maps the source direction into the position on the lens plane, namely the direction of each image due to the lens effect \\cite{SEF}. In particular, the deflection angle is written as $4GM/bc^2$, where $G$ is the gravitational constant, $M$ is the lens mass, $b$ is the impact parameter of the light ray, and $c$ is the speed of light. For the derivation of the deflection angle and consequently the lens equation, the post-Minkowskian approximation $O(G)$ or the linear-order approximation $O(M)$ of the Schwarzschild metric is employed. Do second-order relativistic corrections affect microlensing through the amplification factor? This has been already answered \\cite{EOAK}. The total amplification remains unchanged at $O(G^2 M^2)$, whereas the magnification factor for each image is corrected at this order. The main purpose of this paper is to generalize their relativistic result. In particular, theoretical models beyond the theory of general relativity have attracted a lot of interests last decades, mostly motivated by the dark energy and dark matter problems. In this letter, therefore, we shall reexamine the amplification of lensed images by taking account of corrections in a rather general form, where we assume static, spherically symmetric spacetimes. We use the units of $c=1$ but keep $G$ in order to make iterative calculations clearer. ", "conclusions": "We investigated corrections to the conventional lens equation in terms of differentiable functions, so that they can express not only the second-order effects of $GM$ in general relativity but also modified gravity theories. It was shown that, provided that the spacetime is static, spherically symmetric and asymptotically flat, the total amplification by microlensing remains unchanged at the linear order of the correction to the deflection angle, if and only if the correction takes a particular form as the inverse square of the impact parameter, whereas the magnification factor for each image is corrected. It is concluded that the light curve shape by microlensing is inevitably changed and will thus allow us to probe modified gravity, unless a modification to the deflection angle takes the particular form. It is left as a future work to use the present formulation to probe modified gravity models by microlensing observations. Microlensing in our galaxy is sensitive to gravity at short scale around a few AU, whereas cosmological microlensing is more useful for the large distance physics \\cite{SEF}." }, "1101/1101.0325_arXiv.txt": { "abstract": "{{\\bfseries\\scshape Abstract} \\\\ \\par We consider a phenomenological extension of the minimal supersymmetric standard model which incorporates non-minimal chaotic inflation, driven by a quartic potential associated with the lightest right-handed sneutrino. Inflation is followed by a Peccei-Quinn phase transition based on renormalizable superpotential terms, which resolves the strong CP and $\\mu$ problems of the minimal supersymmetric standard model provided that one related parameter of the superpotential is somewhat small. Baryogenesis occurs via non-thermal leptogenesis, which is realized by the inflaton decay. Confronting our scenario with the current observational data on the inflationary observables, the baryon assymetry of the universe, the gravitino limit on the reheating temperature and the upper bound on the light neutrino masses, we constrain the effective Yukawa coupling involved in the decay of the inflaton to relatively small values and the inflaton mass to values lower than $10^{12}~\\GeV$. } \\\\ \\\\ {\\small \\sc Keywords}: {\\small Cosmology, Supersymmetric models};\\\\ {\\small \\sc PACS codes:} {\\small 98.80.Cq, 12.60.Jv}\\\\\\\\ \\hspace*{-1.43cm} \\publishedin{{\\sl J. Cosmol. Astropart. Phys. } {\\bf 02}, 019 (2011)} ", "introduction": "Recently \\emph{non-minimal inflation} (non-MI) \\cite{wmap3}, i.e. inflation arising in the presence of a non-minimal coupling between the inflaton field and the Ricci scalar curvature, $\\rcc$, has gained a fair amount of attention \\cite{sm1, love, unitarizing, nmi, ld, linde1, linde2}. In particular, it is shown that non-minimal chaotic inflation based on a quartic potential \\cite{nmchaotic} with a quadratic non-minimal coupling to gravity can be realized in both a non-supersymmetric \\cite{sm1, love, ld} and a \\emph{sypersymmetric} (SUSY) framework \\cite{linde1, linde2}, provided that the inflaton couples strongly enough to $\\rcc$. In the latter case, the recently developed \\cite{linde2} superconformal approach to \\emph{supergravity} (SUGRA) greatly facilitates the relevant model building. In most of the models proposed, the inflaton is identified with the Higgs field(s) of the \\emph{Standard Model} (SM) or the next-to-MSSM (\\emph{Minimal SUSY SM}) \\cite{linde1,linde2} -- see also \\cref{SusyHiggs}. Motivated by the various attractive features of the MSSM \\cite{eg} -- such as the resolution of the hierarchy problem, the achievement of gauge coupling unification and the candidature of the lightest SUSY particle as cold dark matter -- we consider it as the starting point of our investigation. Despite its successes, however, the MSSM fails to address a number of important issues. For instance, the strong CP and $\\mu$ problems, the generation of the observed \\emph{baryon asymmetry of the universe} (BAU) and the existence of tiny but non-zero neutrino masses are some fundamental issues which remain open within the MSSM. For the resolution of these, it seems imperative to supplement the MSSM with additional superfields, which in the simplest cases are singlets under the SM gauge group, ${G_{\\rm SM}}= SU(3)_{\\rm c}\\times SU(2)_{\\rm L}\\times U(1)_{Y}$, so that gauge coupling unification is not disrupted. Consequently, new candidates (besides the Higgs boson) for driving non-MI arise. In \\cref{suzuki, tony, rsym} a resolution to the aforementioned problems of the MSSM was proposed within a framework that implements a \\emph{Peccei-Quinn symmetry breaking phase transition} (PQPT). In those models non-renormalizable superpotential terms are added, involving some singlets that develop \\emph{vacuum expectation values} (VEVs) of the order of the PQ symmetry breaking scale. As a consequence, the $\\mu$ and the strong CP problems \\cite{pq} of the MSSM can be simultaneously solved, and in addition a new intermediate scale arises which generates Majorana masses for three \\emph{right-handed} (RH) neutrinos, $\\sni$. The inclusion of $\\sni$ is necessary so that the smallness of neutrino masses is explained through the well-known see-saw mechanism \\cite{seesaw}. These same superfields can play an important role in the generation of the BAU via non-thermal leptogenesis \\cite{inlept, chaotic1, baryo}. This latter attractive possibility is invalidated, though, in the cases studied in \\cref{suzuki, tony}, where the PQPT follows a period of thermal inflation \\cite{thermalI} that leads to a very low reheating temperature. An enormous entropy production occurs, diluting any preexisting, non-thermally created, lepton asymmetry. This dilution can be avoided if we adopt the scheme of \\cref{rsym} but then, the PQ field cannot be zero during inflation -- see below. On the other hand, non-thermal leptogenesis can be enhanced if the scalar component, $\\sn$, of the lightest $\\sni$ is the inflaton itself as firstly proposed in \\cref{murayama}. In this case the branching ratio of the inflaton decay (which now triggers leptogenesis) into a lepton plus a Higgs boson is \\cite{baryo} maximized. However, $\\sn$-inflation, in its simplest realization \\cite{tony, murayama, sneutrino1}, is of the chaotic type -- for other scenarios see \\cref{sneutrinoF, sneutrinoD} -- and therefore, trans-Planckian inflaton-field values are typically required to allow for a sufficiently long period of inflation. The implementation of inflation then necessitates the adoption of special types of \\Ka, as in \\cref{chaotic1}, so that SUGRA corrections are kept under control -- for other proposals related to chaotic inflation with a quadratic potential, see \\cref{sneutrino2}. Moreover, minimal chaotic inflation driven by a quartic potential seems \\cite{circ} to be ruled out by the fitting to the seven-year data of the \\emph{Wilkinson Microwave Anisotropy Probe Satellite} (WMAP7), \\emph{baryon-acoustic-oscillations} (BAO) and \\emph{Hubble constant} ($H_0$) data \\cite{wmap}. In this paper we construct a model of \\emph{non-minimal $\\sn$ inflation} (\\FHI) retaining the successful ingredients of the picture above. To this aim, $\\sn$ (the lightest RH sneutrino) is coupled to one of the PQ fields, which can be confined to zero during inflation -- see \\cref{linde2}. We then show that the model naturally leads to non-MI within SUGRA, provided that a particular parameter of the superpotential is sufficiently small. Sub-Planckian values of the inflaton field are allowed in a wide range of the parameter space, and the adopted type of \\Ka\\ is more or less well-motivated. Also the inflationary observables turn out to lie within the current data. The \\FHI is followed by a PQPT driven by renormalizable superpotential terms as in \\cref{goto, pqhi}, whereas the $\\mu$ parameter of the MSSM can be generated from the PQ scale via a non-renormalizable term as in \\cref{suzuki, tony, rsym}. The reheating temperature is determined exclusively by the decay of $\\sn$ and is high enough ($>100~\\GeV$) so that non-perturbative electroweak sphalerons are operative and, consequently, non-thermal \\cite{inlept} leptogenesis and the subsequent generation of the BAU can be realized. As usually in similar models -- cf. \\cref{baryo, Ndomination, sneutrino1, sneutrino2, sneutrinoF, sneutrinoD} -- consistency with the constraint on the gravitino ($\\Gr$) abundance \\cite{gravitino, brand, kohri} requires a relatively small effective Yukawa coupling constant $\\lf10^{-8}-10^{-3}\\rg$. The smallness of this coupling though may be explained through a broken flavor symmetry \\cite{baryo, Ndomination}. Below, we present the basic ingredients of our model (Sec.~\\ref{fhim}) and describe the inflationary (Sec.~\\ref{fhi}) and post-inflationary dynamics (Sec.~\\ref{pfhi}). We then restrict the parameters of our model (Sec.~\\ref{cont}) and summarize our conclusions (Sec.~\\ref{con}). Details concerning the formulation of non-minimally coupled scalar fields within SUGRA are presented in the Appendix. Throughout the text, we use natural units for Planck's constant, Boltzmann's constant and the speed of light ($\\hbar=c=k_{\\rm B}=1$); the subscript of type $,\\chi$ denotes derivation \\emph{with respect to} (w.r.t.) the field $\\chi$ (e.g., $_{,\\chi\\chi}=\\partial^2/\\partial\\chi^2$); charge conjugation is denoted by a star and $\\log~[\\ln]$ stands for logarithm with basis $10~[e]$. Finally, we follow the conventions of \\cref{kolb} for the quantities related to the gravitational sector of our model. ", "conclusions": "\\label{con} In this paper we attempted to embed within a realistic cosmological setting one of the recently formulated \\cite{linde2} SUSY models of chaotic inflation with non-minimal coupling to gravity. We concentrated on a moderate extension of the MSSM augmented by three RH neutrino superfields and three other singlet superfields, which lead to a PQPT tied to renormalizable superpotential terms. The coupling between the RH neutrinos and one of the fields associated with the PQPT plays a crucial role for the implementation of our scenario. We showed that the model non only supports non-MI driven by the lightest RH sneutrino, but it also resolves the strong CP and the $\\mu$ problems of the MSSM and, even more, it leads to the production of the required by the observations BAU via non-thermal leptogenesis, which accompanies the inflaton's decay. Moreover the $\\Gr$ abundance becomes observationally safe for $\\Gr$ masses even lower than $10~{\\rm TeV}$. An important prerequisite for all these is that the parameter of the superpotential related to the PQPT, $\\la$, is adequately small. Imposing a number of observational constraints arising from the data on the inflationary observables, the BAU, the concentration of the unstable $\\Gr$ at the onset of nucleosynthesis and the mass of the heaviest light neutrino, we restrict the effective Yukawa coupling, involved in the decay of the inflaton, to relatively small values, and the inflaton mass to values lower than $10^{12}~\\GeV$. \\begin{acknowledgement} \\paragraph{} We would like to cordially thank G. Lazarides for helpful discussions and J.~McDonald for an enlightening correspondence. \\end{acknowledgement} \\appendix \\setcounter{equation}{0} \\renewcommand{\\theequation}{A.\\arabic{equation}} \\renewcommand{\\thesubsection}{A.\\arabic{subsection}}" }, "1101/1101.5500_arXiv.txt": { "abstract": "{% The star $\\zeta$\\,Ophiuchi is one of the brightest massive stars in the northern hemisphere and was intensively studied in various wavelength domains. The currently available observational material suggests that certain observed phenomena are related to the presence of a magnetic field. We acquired spectropolarimetric observations of $\\zeta$\\,Oph with FORS\\,1 mounted on the 8-m Kueyen telescope of the VLT to investigate if a magnetic field is indeed present in this star. Using all available absorption lines, we detect a mean longitudinal magnetic field $\\left< B_z\\right>_{\\rm all}= 141\\pm45$\\,G, confirming the magnetic nature of this star. We review the X-ray properties of $\\zeta$\\,Oph with the aim to understand whether the X-ray emission of \\zoph\\ is dominated by magnetic or by wind instability processes. } ", "introduction": "\\label{sect:intro} During the last years a gradually increasing number of O, early B-type, and WR stars have been investigated for magnetic fields, and as a result, about a dozen magnetic O-type stars are presently known (e.g., Hubrig et al.\\ \\cite{Hubrig2008}; Martins et al.\\ \\cite{Martins2010}; Hubrig et al.\\ \\cite{Hubrig2011a}). The recent detections of magnetic fields in massive stars generate a strong motivation to study the correlations between evolutionary state, rotation velocity, and surface composition, and to understand the origin and the role of magnetic fields in massive stars. The star $\\zeta$\\,Ophiuchi (=HD\\,149757) of spectral type O9.5V is a well-known rapidly rotating runaway star with extremely interesting characteristics. It undergoes episodic mass loss seen as emission in H$\\alpha$, and it is possible that it rotates with almost break-up velocity with $v$\\,sin\\,$i=400$\\,km\\,s$^{-1}$ (Kambe et al.\\ \\cite{Kambe1993}). Various studies indicate different types of spectral and photometric variability. The UV resonance lines show multiple discrete absorption components (DAC) in the UV (e.g.\\ Howarth et al.\\ \\cite{Howarth1984}) and strong line profile variations in optical spectra reconciled with traveling sectorial modes of high degree (e.g.\\ Reid et al.\\ \\cite{Reid1993}). Highly precise {\\em MOST} (Microvariability and Oscillations of Stars) satellite photometry in 2004 has yielded at least a dozen significant oscillation frequencies between 1 and 10 cycles/day, hinting at a behaviour similar to $\\beta$~Cephei-type stars (Walker et al.\\ \\cite{Walker2005}). No unambiguous rotation period could be identified in spectroscopic and photometric observations, although Balona \\& Kambe (\\cite{BalonaKambe1999}) favored a period in the region of 1 cycle/day. $\\zeta$\\,Oph is also well-known for its variability in the X-ray band. Oskinova \\etal\\ (\\cite{Oskinova2001}) studied the {\\em ASCA} observations of \\zoph\\ that covered just more than the expected rotation period of the star. A clearly detected periodic X-ray flux variability with $\\sim$20\\%\\ amplitude in the {\\em ASCA} passband (0.5-10\\,keV) was reported. A period of 0.77\\,d was detected and a possible connection with the recurrence time (0.875\\,d$\\pm$0.167\\,d) of the DACs in UV spectra of the star was discussed. The DACs in the spectra of O stars are commonly explained by large-scale structures in the stellar wind, modulated by rotation and possibly related to a surface magnetic field (Cranmer \\& Owocki \\cite{CranmerOwocki1996}). Waldron \\etal\\ (in preparation, private communication) found that {\\em SUZAKU} data on \\zoph\\ suggest a period of $\\sim$0.98\\,d that is consistent but slightly larger than the X-ray periodicities found in {\\em ASCA} data (Oskinova \\etal\\ \\cite{Oskinova2001}) and in {\\em Chandra} HETGS data (Waldron \\cite{Waldron2005}). In addition, the HETGS data appear to indicate an additional cyclic period of $\\sim$0.33\\,d in the hard X-ray band ($>$1.2\\,keV). The results of our previous studies seem to indicate that the presence of a magnetic field is more frequently detected in candidate runaway stars than in stars belonging to clusters or associations (Hubrig et al.\\ \\cite{Hubrig2011b}; Hubrig et al.\\ \\cite{Hubrig2011a}). The currently best available astrometric, spectroscopic, and photometric data were used to calculate the kinematical status of magnetic O-type stars with previously unknown space velocities. The results suggest that most of the magnetic O-type stars can be considered as candidate runaway stars. The available observational material suggests that $\\zeta$\\,Oph is a main sequence single star in the field with runaway characteristics. Usually, to explain the origin of massive stars in the field, two mechanisms are discussed in the literature. In one scenario, close multibody interactions in a dense cluster cause one or more stars to be scattered out of the region (e.g.\\ Leonard \\& Duncan \\cite{LeonardDuncan1990}). For this mechanism, runaways are ejected in dynamical three- or four-body interactions. An alternative mechanism involves a supernova (SN) explosion within a close binary, ejecting the secondary due to the conservation of momentum (Zwicky \\cite{Zwicky1957}; Blaauw \\cite{Blaauw1961}). Blaauw (\\cite{Blaauw1952}) suggested the origin of $\\zeta$\\,Oph in the Scorpius OB2 association due to its proper motion vector, which points away from the association. More recently, Hoogerwerf et al.\\ (\\cite{Hoogerwerf2001}) suggested that the star gained it space velocity of $\\sim$30\\,km\\,s$^{-1}$ in a supernova explosion within a close binary in Upper Scorpius about 1--2\\,Myr ago. The authors identified PSR~B1929+10 as an associated pulsar with a characteristic age of $\\sim$3\\,Myr, consistent with the kinematic age of $\\zeta$\\,Oph within the uncertainties. Tetzlaff et al.\\ (\\cite{Tetzlaff2010}) reinvestigated the scenario of a binary SN in Upper Scorpius involving $\\zeta$\\,Oph and PSR~B1929+10 and concluded that it is very likely that both objects were ejected during the same supernova event. In their work, the considered association age range implies that the progenitor star of the produced neutron star had a spectral type between O6/O7 and O9 with a mass range from 18 to 37\\,\\Msun{}. The X-ray emission of the pulsar seems to be dominated by non-thermal radiation processes (e.g.\\ Becker et al.\\ \\cite{Becker2006}). An arc-like nebula surrounding PSR~B1929+10 and extending up to 10\\arcsec{} was identified in {\\em Chandra} data and interpreted as a bow-shock nebula (Hui \\& Becker \\cite{HuiBecker2008}). The estimated magnetic field strength in the shocked region accounts for $\\sim$75\\,$\\mu$G, while the typical magnetic field strength in the ISM is about 2--6\\,$\\mu$G. The presence of a bow-shock nebula has also been detected for $\\zeta$\\,Oph. Figure~\\ref{fig:bsh} shows an image based on archival {\\em Spitzer} IRAC maps (AOR 17774848). Recently, Kobulnicky \\etal\\ (\\cite{Kobulnicky2010}) analyzed a sample of bow shocks around massive stars in Cygnus-X. They used the analytical description of momentum-driven bow shocks and dust/polycyclic aromatic hydrocarbon emission models to estimate stellar mass loss rates from the observed properties of the bow shocks. It was found that mass-loss rates in the range between $10^{-7}$\\myr\\ and a few times $10^{-6}$\\myr\\ are required to generate the bow shocks around typical B2V - O5V type stars. \\begin{figure} \\centering \\includegraphics[width=0.45\\textwidth]{zoph-irac.2.ps} \\caption{(online colour at: www.an-journal.org) Combined IR {\\em Spitzer} IRAC (3.6\\,\\mim\\ blue, 4.5\\,\\mim\\ green, 8.0\\,\\mim\\ red) image of the bow shock around the runaway star \\zoph. Archival data have been used. Galactic coordinates are shown. The image size is $\\sim 36' \\times 31'$. } \\label{fig:bsh} \\end{figure} The mass-loss rate \\mdot\\ of this star was empirically obtained from different diagnostics by a number of authors. Repolust \\etal\\ (\\cite{Repolust2004}) fitted the H$\\alpha$ photospheric absorption line and derived the upper limit on the \\zoph\\ mass-loss rate as $1.8\\times 10^{-7}$\\myr. Fullerton \\etal\\ (\\cite{Fullerton2006}) determined the radio-based mass-loss rate of \\zoph\\ as $1.1\\times 10^{-7}$\\myr. The mass-loss rates determined from radio depend on the square of the density since the physical mechanism responsible for the radio emission is free-free emission. On the other hand, Fullerton \\etal\\ derive a much smaller mass-loss rate from fitting the UV P\\,{\\sc v} resonance doublet, the product of the mass-loss rate and the ion fraction of P$^{+4}$ being only $\\dot{M}q({\\rm P^{+4}}) \\lsim 1.3 \\times 10^{-10}$\\myr\\ with $q({\\rm P^{+4}})\\lsim 1$. The mass-loss rates derived from fitting the wind profiles of UV resonance lines depend linearly on the density. To resolve this discordance in mass-loss determinations based on $\\rho^2$- and $\\rho$-diagnostics, Fullerton \\etal\\ suggest that the winds are strongly clumped with a volume filling factor of $\\sim$10$^{-3}$--10$^{-5}$. Marcolino \\etal\\ (\\cite{Marcolino2009}) analyzed optical and UV spectra of \\zoph\\ among their sample of O-type dwarfs. They derive an upper limit on the mass-loss rate of \\zoph\\ as $1.6 \\times 10^{-9}$\\myr\\ if the wind was smooth. This value agrees with the $\\dot{M}q({\\rm P^{+4}})$ value obtained by Fullerton \\etal\\ (\\cite{Fullerton2006}). Using the example of the O-type supergiant $\\zeta$\\,Puppis, Oskinova \\etal\\ (\\cite{Oskinova2007}) demonstrated that the discordance of mass-loss rates found by Fullerton \\etal\\ can be overcome by accounting for stellar wind porosity (see also Sundqvist \\etal\\ \\cite{Sundqvist2010}). It was found for the O5Ia star $\\zeta$\\,Puppis that only a moderate reduction of the mass-loss rate by a factor of 2--3 (compared to the smooth wind models) is required to reproduce both H$\\alpha$ and P\\,{\\sc v} lines. If this result holds also for non-supergiant O type stars, the mass-loss rate of \\zoph\\ is only a few times lower compared to the radio-based mass-loss determined by Fullerton \\etal, i.e.\\ $\\sim$10$^{-7}$\\,\\myr. Importantly, this mass-loss rate is in agreement with values that are required to produce bow shocks around O stars (Kobulnicky \\etal\\ \\cite{Kobulnicky2010}). An additional aspect, which may hint at the presence of a magnetic field in runaway stars, is that a number of individual abundance studies indicate nitrogen enrichment in the atmospheres of runaway stars (e.g.\\ Boyajian et al.\\ \\cite{Boyajian2005}). Nitrogen enrichment was found in $\\zeta$\\,Oph by Villamariz \\& Herrero (\\cite{VillamarizHerrero2005}). Recent NLTE abundance analyses (e.g., Morel et al.\\ \\cite{Morel2008}; Hunter et al.\\ \\cite{Hunter2008}) suggest that slow rotators have peculiar chemical enrichment such as nitrogen excess or boron depletion, and these peculiarities are linked to the presence of a magnetic field. On the other hand, Hubrig et al.\\ (\\cite{Hubrig2011c}) showed that some magnetic massive stars previously assumed to be slow rotators, are in fact fast rotators, but are viewed close to their rotation poles. To test the magnetic nature of this particularly interesting rapidly rotating runaway star, we acquired spectropolarimetric observations with the low-resolution spectropolarimeter FORS\\,1 at the VLT. In this work we report the first detection of a magnetic field in this star. ", "conclusions": "\\label{sect:discussion} $\\zeta$\\,Oph has been extremely well studied in all wavelength ranges, from the X-ray by all major X-ray satellites (with the exception of \\xmm) to the infrared region with {\\em Spitzer}. In view of the detection of a magnetic field on \\zoph\\ reported in this work, we review its X-ray properties with the aim to understand whether the X-ray emission of \\zoph\\ is dominated by magnetic or wind instability processes. Babel \\& Montmerle (\\cite{BabelMontmerle1997}) studied the case of a rotating star with a dipole magnetic field sufficiently strong to confine stellar wind. The magnetic field locally dominates the bulk motion of stellar wind, when the ratio of magnetic to kinetic energy density, $B^2/\\mu_0\\rho v^2 > 1$, where $v$ is the supersonic flow speed. A collision between the wind streams from the two hemispheres in the closed magnetosphere leads to a strong shock and X-ray emission. MHD simulations in the framework of this magnetically confined wind shock (MCWS) model were performed by ud-Doula \\& Owocki (\\cite{udDoulaOwocki2002}) and Gagn{\\'e} \\etal\\ (\\cite{Gagne2005}). Using their notation, the wind is confined when \\mbox{$\\eta_\\ast\\equiv (R_\\ast^2B^2)(\\dot{M}\\vinf)^{-1} > 1$}. New observations are required to establish whether the magnetic field of \\zoph\\ is a dipole. However, for the purpose of this discussion, let us assume that the field has an average strength of 150\\,G. Using the stellar parameters of \\zoph\\ as inferred by Marcolino \\etal\\ (\\cite{Marcolino2009}), we estimate $\\eta_\\ast(\\zoph)\\sim 10^3$, i.e.\\ the magnetic field should dominate the wind motion up to the Alfv{\\'e}n radius that is located at $\\lsim$10\\,$\\Rstar$. In this case, the X-ray emission should mainly originate from the MCWS. The MCWS model predicts that the X-ray emitting plasma should be located at a few \\Rstar\\ from the photosphere; that the X-ray emission lines should be narrow; that the X-ray luminosity should be higher and the spectrum harder than in non-magnetic stars; that in oblique magnetic rotators the X-ray emission should be modulated periodically as a consequence of the occultation of the hot plasma by a cool torus of matter, or by the opaque stellar core. The lines of He-like ions observed in X-ray spectra are useful to derive the location of the line formation region in hot stars because forbidden line emission is depressed by ultraviolet pumping. The latter depends on the distance to the stellar photosphere (Gabriel \\& Jordan \\cite{GabrielJordan1969}; Blumenthal \\etal{} \\cite{Blumenthal1972}). The Si\\,{\\sc xiii} line observed in the \\cxo\\ HETGS/MEG spectrum is shown in Fig.~\\ref{fig:sixiii}. prominent forbidden line can easily be distinguished in this fugure, while normally forbidden lines are strongly suppressed in the spectra of OB-type stars. The presence of the forbidden line implies that the line formation region is located far from the photosphere, so that the radiative excitation does not lead to the depopulation of the corresponding metastable energy levels. Waldron \\& Cassinelli (\\cite{WaldronCassinelli2008}) found that the Si\\,{\\sc xiii} line is formed at $1.8 \\pm 0.7$\\,\\Rstar\\ in \\zoph\\ and that other He-like lines are formed even further out in the wind. Interestingly, the strong forbidden Si\\,{\\sc xiii} line is also observed in the {\\em Chandra} spectrum of the magnetic star $\\tau$\\,Sco. Cohen \\etal\\ (\\cite{Cohen2003}) derive a Si\\,{\\sc xiii} line formaiton radius for $\\tau$\\,Sco in the range between 1.1\\,$R_\\ast$ and 1.5\\,$R_\\ast$. These radii of line formation are smaller than those found in the prototypical MCWS model object $\\theta^1$\\,Ori\\,C, $3.4 \\pm 0.8$\\,\\Rstar\\ (Waldron \\& Cassinelli \\cite{WaldronCassinelli2008}). \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{SiXIII.ps} \\caption{The Si\\,{\\sc xiii} line observed in the spectrum of \\zoph{} (co-added MEG $\\pm$1). Vertical dashed lines indicate the rest-frame wavelength: $\\lambda_{\\rm R}$ -- resonant line, $\\lambda_{\\rm I}$ -- sum of intercombination lines, $\\lambda_{\\rm F}$ -- forbidden line. The rest-frame wavelengths are corrected for the radial velocity taken from Hoogerwerf \\etal\\ (\\cite{Hoogerwerf2001}).} \\label{fig:sixiii} \\end{figure} Oskinova \\etal\\ (\\cite{Oskinova2006}) studied the \\cxo\\ spectrum of \\zoph\\ among other O-type stars. They found that the X-ray emission lines in this star are narrow and that the signatures of wind absorption on line profiles are weak. Figure~\\ref{fig:vel} shows the Fe\\,{\\sc xvii} and Ne\\,{\\sc x} lines as measured by \\cxo\\ plotted over units of the wind terminal velocity, \\vinf=1550\\,km\\,s$^{-1}$. The lines are only slightly broadened, if at all. \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{twoline_vel.ps} \\caption{Fe\\,{\\sc xvii} (upper panel) and Ne\\,{\\sc x} (lower panel) lines observed in the spectrum of \\zoph\\ (co-added MEG $\\pm$1). Vertical dashed lines indicate the rest-frame wavelength, corrected for the radial velocity. } \\label{fig:vel} \\end{figure} The X-ray luminosity of \\zoph, $\\Lx=1.2\\times 10^{31}$\\,erg\\,s$^{-1}$ and the ratio $\\Lx/\\Lbol=4\\times 10^{-8}$ are quite usual among late type OV stars (Oskinova \\etal\\ \\cite{Oskinova2006}). Adopting the mass-loss rate from \\zoph\\ as $\\lsim 1.8\\times 10^{-7}$\\,\\myr{}, Oskinova \\etal{} (\\cite{Oskinova2006}) noticed that in \\zoph{} the ratio of X-ray to wind mechanical luminosity $L_{\\rm mech}$ ($\\mdot\\vinf^2/2$), $L_{\\rm X}/L_{\\rm mech} \\msim 8.5\\times 10^{-5}$, is a few times higher than in other single O-type stars. This may be related to the lower wind opacity in \\zoph, or it may hint at some additional mechanism of X-ray generation besides the intrinsic wind shocks. From their analysis of \\cxo\\ spectra, Zhekov\\& Palla (\\cite{ZhekovPalla2007}) derived the differential emission measure (DEM) for \\zoph\\ among other OB stars in their sample. They found that in \\zoph\\ the DEM sharply peaks at about 6\\,MK. While this is a significantly lower temperature than found for the DEM peak in case of $\\theta^1$\\,Ori\\,C (50\\,MK), it is higher than found for other stars of similar spectral types ($\\sim$3\\,MK). Thus, considering the X-ray temperature of \\zoph{}, it is not straightforward to attribute its X-ray emission to the MCWS. On the other hand, recent studies of O stars with detected magnetic fields (e.g., HD\\,191612, HD\\,108) show that their X-ray properties are diverse (Naz{\\'e} \\etal\\ \\cite{Naze2004}, \\cite{Naze2010b}) and may be difficult to fully reconcile with the predictions of the MCWS model. Clearly, new observations are needed to better understand the magnetic field of \\zoph\\ and its link with the X-ray emission from this star. {" }, "1101/1101.1503_arXiv.txt": { "abstract": "The brightness of gamma-ray burst (GRB) afterglows and their occurrence in young, blue galaxies make them excellent probes to study star forming regions in the distant Universe. We here elucidate dust extinction properties in the early Universe through the analysis of the afterglows of all known $z>6$ GRBs: GRB\\,090423, 080913 and 050904, at $z=8.2,~6.69$, and $6.295$, respectively. We gather all available optical and near-infrared photometry, spectroscopy and X-ray data to construct spectral energy distributions (SEDs) at multiple epochs. We then fit the SEDs at all epochs with a dust-attenuated power-law or broken power-law. We find no evidence for dust extinction in GRB\\,050904 and GRB\\,090423, with possible evidence for a low level of extinction in GRB\\,080913. We compare the high redshift GRBs to a sample of lower redshift GRB extinctions and find a lack of even moderately extinguished events ($A_V\\sim0.3$) above $z\\gtrsim4$. In spite of the biased selection and small number statistics, this result hints at a decrease in dust content in star-forming environments at high redshifts. ", "introduction": "} Dust formation in the early universe is a hotly debated topic. The properties and quantity of dust at these times are essential not only to our understanding of how dust forms and evolves, but also to models of star-formation and the appearance of the first generations of stars and galaxies. Currently we extrapolate dust properties in the local group to very different environments that existed less than a billion years after the Big Bang. Long-duration gamma-ray bursts (GRBs), associated with the deaths of massive stars \\citep{woosley, galama, hjorth, stanek, malesani,campana06}, have extremely bright and spectrally simple afterglows $(F_{\\nu}\\propto \\nu^{-\\beta})$ \\citep{sari, granot}, and so provide an opportunity to obtain not only effective reddening, but also absolute extinctions from distant galaxies (e.g. \\citealt{watson06}; \\citealt{schady}; \\citealt{zafar11}). The recent tremendous advances in this subject were enabled by the dedicated GRB satellite, \\emph{Swift} \\citep{gehrels}, with its precise localization and fast responding X-ray and UV/optical telescopes, detecting the afterglow within only a few tens of seconds after the $\\gamma$-ray trigger. The data from \\emph{Swift} have provided a very large sample of afterglows virtually complete in X-ray detections, and with a high completeness in the optical. It has also for the first time, enabled the discovery of GRBs at redshifts greater than or comparable to the most distant known galaxies and quasars (QSO) \\citep[see][]{kawai,greiner,tanvir,salvaterra}. In this paper we collect for the first time all the available afterglow photometry and spectra from near-infrared (NIR) to X-ray at different epochs for all GRBs at $z>6$: GRB\\,050904 ($z=6.295$), GRB\\,080913 ($z=6.69$) and GRB\\,090423 ($z\\sim8.2$). Since observations of these high-z GRBs are being made in the rest frame UV, they are in principle relatively sensitive to even small amounts of dust. We extract the SEDs at multiple epochs for all three GRBs and jointly fit the X-ray to NIR data to determine the properties of dust at $z>6$. In \\S2 we describe multi-wavelength observations of the afterglows at different epochs carried out with different instruments. In \\S3 we present our results from the SED fitting. Based on our results, in \\S4 we discuss the origin of dust at high redshift and compare to the first large spectroscopic sample of GRB-derived extinction curves, the first spectroscopic sample of absolute extinction curves outside the local group. In \\S5 we provide our conclusions. For all parameters we quote uncertainties at the 68\\% confidence level. We derive $3\\sigma$ upper limits for the cases where detection is less than $2\\sigma$ significant. A cosmology where $H_0=72$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_\\Lambda = 0.73$ and $\\Omega_{\\rm m}=0.27$ is assumed throughout. ", "conclusions": "} We have investigated cosmic dust at $z>6$ using the afterglows of the highest redshift GRBs known. We performed multi-epoch NIR--X-ray SED analysis of GRB\\,050904, 080913, and 090423. We infer from our analysis that there is no evidence of dust in GRB\\,050904 and GRB\\,090423. We find possible evidence of low extinction in GRB\\,080913 consistent with an SMC extinction curve, however the optical spectrum and NIR photometry alone are consistent with a power-law with no dust extinction. In no case did we find evidence for any other extinction curve (e.g.\\ the `SN-origin' curve of \\citealt{maiolino}). Comparing to a much larger spectroscopic sample of GRB extinctions, we find a distinct absence of extinction at high redshifts. While the high redshift sample is very small and some of the effect may be explained by restframe UV-selection bias, this is not the whole picture, as all the $z>6$ GRBs would have been detected at the typical $A_V$ of lower redshift reddened bursts. This hints that there is less dust along sightlines to the highest redshift GRBs, indicating less dust in the early Universe. From these results, we also infer an extremely low dust-to-metals ratio in GRBs at high redshift, suggestive either of efficient dust destruction, or a delay of at least several million years between the formation of metals and the formation and growth of dust." }, "1101/1101.1029_arXiv.txt": { "abstract": "{The most compelling and popular models for dark matter predict that it should congregate and annihilate in stellar cores. Stars where annihilation contributes substantially to the total energy budget look very different to those with which we are familiar. Here I explain the general features of stars modified by dark matter annihilation with the help of a series of grids of `dark' stellar evolutionary models, and describe the public code with which they were computed. I go on to discuss possible impacts of dark stars on the high-redshift Universe, including the history of reionisation. The preliminary reionisation calculations reproduced here are based on dedicated models for dark star atmospheres, and for the stellar populations to which dark stars would belong.} \\FullConference{Cosmic Radiation Fields: Sources in the early Universe\\\\ November 9-12, 2010\\\\ DESY, Germany} \\begin{document} ", "introduction": "Weakly-interacting massive particles (WIMPs) are amongst the most promising candidates for dark matter \\cite{Jungman96, Bergstrom00, Bertone05}. In the standard scenario, WIMPs are Majorana particles that self-annihilate with a cross-section of $\\sim3\\times10^{-26}$\\,cm$^3$\\,s$^{-1}$, leading to a thermal relic dark matter abundance very similar to the observed value. If such WIMPs are somehow confined within stellar cores, the resultant annihilation can have a substantial effect on stellar properties \\cite{SalatiSilk,BouquetSalati}. The idea that such `dark stars' could exist has been pursued rather vigorously in recent years, with attention given to white dwarfs \\cite{MoskalenkoWai,Bertone08,McCullough10,Hooper10}, neutron stars \\cite{Bertone08,Fairbairn10,Kouvaris10} and main-sequence stars \\cite{Scott07,Fairbairn08,Scott09,Casanellas09,Scott09proc} in our own Galaxy, as well as to the first stars (Pop III) \\cite{Spolyar08,Iocco08a,Freese08a,Iocco08b,Freese08b,Yoon08,Taoso08,Freese09,Natarajan09,Ripamonti09,Spolyar09,Iocco09,Umeda10,Sivertsson10,Ripamonti10,Gondolo10}. The effect of dark matter annihilation on Pop III stars has been examined in the context of broader astronomical observations, including reionisation \\cite{Schleicher, Venk11}, HST/JWST deep-field magnitudes and galaxy spectra \\cite{Zackrisson10a,FreeseSMDS,Zackrisson10b}. For the purposes of these proceedings, and more generally, I define a dark star as `any stellar object whose structure or evolution has been affected by dark matter annihilation'. Other authors prefer to use the term `dark star' to refer exclusively to the dark-matter dominated phase that may occur during the formation of Pop III stars (e.g.~\\cite{Spolyar08}). The distinction is essentially just semantic, but important for readers' understanding of the literature, should they choose to delve into it. These proceedings are not intended as a comprehensive review of dark stars as a field of active research, rather as simply a set of summarising notes of the (semi-review) presentation given at CRF2010. Dark matter can find its way into stellar cores via two distinct physical mechanisms, as outlined in Fig.~\\ref{fig1}. The first is by the gravitational contraction of a baryonic gas cloud, as it cools and collapses during star formation. The baryonic contraction steepens the gravitational potential, drawing dark matter into the centre of collapsing cloud. The contraction of the dark matter need not be entirely adiabatic for this to occur, although adiabaticity is typically assumed. This mechanism can only occur during formation of a star, so the resulting population of dark matter cannot be replaced by the same mechanism after it annihilates away. The second mechanism is by nuclear scattering, whereby WIMPs scatter weakly off nucleons in the star, lose kinetic energy and become gravitationally bound to the star. They then follow bound orbits to return and scatter repeatedly, losing more energy and settling down to the stellar core. This mechanism can continue to be efficient for as long as an appropriate population of dark matter exists for the star to capture from in the halo surrounding it; in this way, capture via nuclear scattering may in principle replenish a star's population of dark matter indefinitely. \\begin{figure} \\centering \\includegraphics[trim = 0 0 305 630, clip=true]{contrac_cap} \\caption{A schematic representation of the two physical mechanisms by which dark stars acquire dark matter. Gravitational contraction (left) occurs where baryonic collapse of a star-forming gas cloud steepens the gravitational potential, drawing dark matter in along with the baryons. Nuclear scattering (right) relies on weak scattering events between WIMPs and nucleons, whereby dark matter particles lose energy and become gravitationally bound to the star.} \\label{fig1} \\end{figure} \\begin{figure}[p] \\begin{minipage}[t]{0.48\\linewidth} \\centering \\includegraphics[width=0.96\\linewidth, trim = 0 0 0 30, clip=true]{HRPlotn3} \\includegraphics[width=0.96\\linewidth, trim = 0 0 0 30, clip=true]{HRPlotn1} \\includegraphics[width=0.96\\linewidth, trim = 0 0 0 30, clip=true]{HRPlot0} \\includegraphics[width=0.96\\linewidth, trim = 0 0 0 30, clip=true]{HRPlotp1} \\end{minipage} \\hspace{0.04\\linewidth} \\begin{minipage}[t]{0.48\\linewidth} \\centering \\includegraphics[width=0.96\\linewidth, trim = 0 0 0 30, clip=true]{TrhoPlotn3} \\includegraphics[width=0.96\\linewidth, trim = 0 0 0 30, clip=true]{TrhoPlotn1} \\includegraphics[width=0.96\\linewidth, trim = 0 0 0 30, clip=true]{TrhoPlot0} \\includegraphics[width=0.96\\linewidth, trim = 0 0 0 30, clip=true]{TrhoPlotp1} \\end{minipage} \\caption{Example evolutionary tracks (left) and central equation-of-state diagrams (right) for a grid of low-mass Pop II stars, when evolved from the main sequence with different amounts of dark matter capture by nuclear scattering. From top to bottom, models exhibit sufficient capture to provide $10^{-3}$, $10^{-1}$, $1$ and $10$ times as much energy from WIMP annihilation as from nuclear fusion. For details, see Ref.~\\cite{Scott09}. Models were computed using the \\textsf{DarkStars} code \\cite{Scott09proc}. Qualitatively similar changes occur for higher masses and other metallicities. Stars that are strongly affected by dark matter annihilation \\emph{before} reaching the main sequence approach similar Hayashi-like solutions, but descend from higher on the Hayashi track instead of ascending to them from the main sequence.} \\label{fig2} \\end{figure} ", "conclusions": "" }, "1101/1101.4867_arXiv.txt": { "abstract": "{A study of the variability of the broad emission-line parameters of 3C390.3, an active galaxy with the double-peaked emission-line profiles, is {presented}. Here we give a detail analysis of variation in the broad H$\\alpha$ and H$\\beta$ emission-line profiles, the ratios, and the Balmer decrement of different line segments.} {With investigation of the variability of the broad line profiles we explore the disk structure, that is assumed to emit the broad double-peaked H$\\beta$ and H$\\alpha$ emission lines in the spectrum of 3C390.3.} {We divided the observed spectra in two periods (before and after the outburst in 2002) and analyzed separately the variation in these two periods. First we analyzed the spectral emission-line profiles of the H$\\alpha$ and H$\\beta$ lines, measuring the peak positions. Then, we divided lines into several segments, and we measured the line-segment fluxes. The Balmer decrement variation for total H$\\alpha$ and H$\\beta$ fluxes, as well as for the line segments has been investigated and discussed. Additionally, we modeled the line parameters variation using an accretion disk model and compare our modeled line parameter variations with observed ones.} {We compared the variability in the observed line parameters with the disk model predictions and we found that the variation in line profiles and in the line segments corresponds to the emission of a disk-like BLR. But, also there is probably one additional emission component that contributes to the H$\\alpha$ and H$\\beta$ line center. We found that the variation in the line profiles is caused by the variation in the parameters of the disk-like BLR, first of all in the inner (outer) radius which can well explain the line parameter variations in the Period I. The Balmer decrement across the line profile has a bell-like shape, and it is affected not only by physical processes in the disk, but also by different emitting disk dimension of the H$\\alpha$ and H$\\beta$ line.} {The geometry of the BLR of 3C390.3 seems to be very complex, and inflows/outflows might be present, but it is evident that the broad line region with disk-like geometry has dominant emission.} ", "introduction": "{The broad emission lines (BELs) are often observed in optical and ultraviolet spectra of acitve galactic nuclei (AGN). The study of the profiles and intensities of BELs can give us relevant information about the geometry and physics of the broad line region (BLR). The physics and geometry of the BLR are uncertain and investigation of BEL shape variability in a long period is very useful for determination of the BLR nature. The profiles of the BELs in AGN can indicate the geometry of emitting plasma in the BLR \\citep[see e.g.][etc.]{sul00,pop04,ga09,za10}. Particularly, very interesting objects are AGNs with unusual broad emission-line profiles, where broad Balmer lines show double peaks or double ''shoulders'', so called double-peaked emitters. The double peaked line profiles may be caused by an accretion disk emission. On the other hand,} the presence of an accretion-disk emission in the BLR is expected, and double-peaked line profiles of some AGN indicate {this} \\citep{pe88,eh94,eh04,er09}. One of the well known AGN with broad double-peaked emission lines in its spectrum is the radio-loud active galaxy 3C 390.3. Although, the double-peaked line profiles can be explained by different hypothesis \\citep[see e.g.][]{vz91}, as e.g. super-massive binary black holes \\citep{gas96}, outflowing bi-conical gas streams \\citep{zh96}), it seems that in this case a disk emission is present in the BLR \\citep[][hereinafter Paper I]{sh10}. There is a possibility that a jet emission can affect the optical emission in 3C 390.3 \\citep{tig10}, and some perturbations in disk could be present \\citep{jov10} and they can also affect the double-peaked line profiles. Long-term variability in the line/continuum flux as well as in the line profiles is observed in objects with broad double peaked lines \\citep[see e.g.][]{di88,sh01,se02,se10,sh10,le10}. Long-term variability of broad line profiles is intriguing because it is usually unrelated to more rapid changes in the continuum flux, but probably is related to physical changes in the accretion disk, as e.g. brightness of some part of the disk, or changes in the disk size and distance to the central black hole. The double-peaked broad line profile variability can be exploited to test various models for the accretion disk (as e.g. circular or elliptical). {Moreover, the double-peaked broad line studies can provide important information about the accretion disk, as e.g. inclination, dimension and emissivity of the disk as well as probes of dynamical phenomena that may occur in the disk \\citep[see in more details][and reference therein]{er09}.} In Paper I we have presented the results of the long-term (1995--2007) spectral monitoring of \\object{3C~390.3}. We have analyzed the light curves of the broad H$\\alpha$ and H$\\beta$ line fluxes and the continuum flux in the 13-year period. We also found that quasi-periodical oscillations (QPO) may be present in the continuum and H$\\beta$ light curves. We studied averaged profile of the H$\\alpha$ and H$\\beta$ line in two periods (Period I from 1995 to 2002, and Period II from 2003 to 2007) and their characteristics (as e.g. peaks separation and their intensity ratio, or FWHM). From the cross-correlations (ICCF and ZCCF) between the continuum flux and H$\\beta$ and H$\\alpha$ lines we found the lag of $\\sim$95 days for H$\\beta$ and $\\sim$120 days for H$\\alpha$ (see Paper I for details). We concluded that the broad emission region has disk-like structure, but there could probably exists an additional component, non-disk or also disk-like, with different parameters that contributes to the line emission. {We found a differences in the H$\\alpha $ and H$\\beta $ line profiles before and after the beginning of the activity phase in 2002, consequently we divided our spectra into two periods (before March 05, 2002 -- Period I and after that -- Period II, see Paper I).} In this paper we study in more details the H$\\alpha$ and H$\\beta$ line profiles and ratios, taking into account the changes during the monitoring period. The aim of this paper is to investigate the changes in the the BLR structure of 3C 390.3 that cause the line profile variations. To perform this investigation we analyzed the peak separation variations, variations in line segments and variation in Balmer decrement. Using a relatively simple disk model, we try to explain qualitatively the changes in disk structure that can cause the line parameter variations. The paper is organized as follow: in \\S 2 we {describe} of our observations; in \\S 3 we {present} the analysis of the H$\\alpha$ and H$\\beta$ line profiles variability, the peak-velocity variability and Balmer decrement; in \\S 4 we study the line-segment variations; in \\S 5 the Balmer decrement variation is analyzed; in \\S 6 we {discuss obtained results} , and finally in \\S 7 we outline our conclusions. ", "conclusions": "In this paper, we present line profile variations of 3C 390.3 in a long period. Due to the change of line profiles, we divided the observations into two periods (before and after the minimum in 2002: Period I and II, respectively) and found difference in the line segments and Balmer decrement variations in these two periods. From our investigation, the main conclusions are the following: i) the line profiles during the monitoring period are changing, always showing the disk-like profile, with the higher blue peak. There is also the central peak that may come from the emission region additional to the disk, but as it was mentioned in \\citet{jov10}, it may also be caused by the perturbation in the disk. ii) the far-wings flux variation in the first period, where the far-red wing flux does not respond well to the continuum flux, is probably caused by some physical processes in the innermost part of the disk. The observed changes in H$\\alpha$ and H$\\beta$ may be interpreted in the framework of a disk model with changes in the location and size of the disk line emitting regions. In Period II, the good correlation between the continuum and line-segment flux suggests that the brightness of the disk is connected with the ionizing continuum, and that the structure of the disk does not significantly change. iii) Balmer decrement is also different for these two periods. In the first period the BD decreases with continuum flux, while in the second period the BD stays more-less constant (around 4.5). The segment BD shows two maxima (around $\\pm$6000 km s$^{-1}$) which do not correspond to the red and blue peak, but instead they are farther in the blue and red wing (than peaks velocities). Also one minimum around zero velocity is present. This minimum changed position between $\\pm$2000 km s$^{-1}$ around zero velocity. This central minimum (as well as shifted maxima) may be caused by the additional (to the accretion disk) emission that perhaps is present in the 3C 390.3 broad lines. These results suggest that, in addition to the physical conditions across disk, the size of the H$\\alpha$ and H$\\beta$ emitting regions of the disk plays an important role. We modeled bell-like BD vs. velocity profiles in the case when the H$\\alpha$ disk is larger than H$\\beta$ one. iv) the variation observed in the line parameters can be well modeled if one assumes changes in position of the emitting disk with respect to the central black hole. The emission of the disk-like region is dominant, but there is the indication of the additional emission. Therefore, to explain the complex line-profile variability one should consider a complex model that may have a disk geometry together with outflows/inflows (see Fig. \\ref{f20}). An important conclusion of this work is that, even if the disk-like geometry plays a dominant role, the variability of the H$\\alpha$ and H$\\beta$ line profiles and intensities (and probably partly in the continuum flux) has different nature for different periods. It seems that in Period I, the perturbation(s) in the disk caused (at least partly) the line and continuum amplification, while in Period II the ionizing continuum caused the line amplification without big changes in the disk-like structure." }, "1101/1101.1997_arXiv.txt": { "abstract": "The formation and evolution of a circumstellar disk in magnetized cloud cores is investigated from prestellar core stage until $\\sim10^4$\\,yr after protostar formation. In the circumstellar disk, fragmentation first occurs due to gravitational instability in a magnetically inactive region, and substellar-mass objects appear. The substellar-mass objects lose their orbital angular momenta by gravitational interaction with the massive circumstellar disk and finally fall onto the protostar. After this fall, the circumstellar disk increases its mass by mass accretion and again induces fragmentation. The formation and falling of substellar-mass objects are repeated in the circumstellar disk until the end of the main accretion phase. In this process, the mass of fragments remain small, because the circumstellar disk loses its mass by fragmentation and subsequent falling of fragments before it becomes very massive. In addition, when fragments orbit near the protostar, they disturb the inner disk region and promote mass accretion onto the protostar. The orbital motion of substellar-mass objects clearly synchronizes with the time variation of the accretion luminosity of the protostar. Moreover, as the objects fall, the protostar shows a strong brightening for a short duration. The intermittent protostellar outflows are also driven by the circumstellar disk whose magnetic field lines are highly tangled owing to the orbital motion of fragments. The time-variable protostellar luminosity and intermittent outflows may be a clue for detecting planetary-mass objects in the circumstellar disk. ", "introduction": "The observation of many star-forming regions has shown that stars form in molecular cloud cores. A star is born after a long journey through the gravitational collapse of a cloud core. The gas begins to collapse in a dense part of the cloud (i.e., the cloud core) and increases its density and temperature as the cloud core collapses further. When the collapsing gas reaches a sufficiently high density ($\\sim10^{21}\\cm$), the collapse stops and a protostar having almost a Jovian mass is born \\citep{larson69}. Then, the protostar increases its mass through mass accretion and finally evolves into a main-sequence star (or main-sequence phase). The star formation process can be divided into two phases: the early protostellar collapse phase and later accretion phase of the star formation. The early phase, also called the gas collapsing phase, is defined as the period before protostar formation after gas collapse is initiated in the cloud core. Thus, during this phase, the gas continues to collapse. On the other hand, the later phase is defined as the period before the main-sequence phase after the gas collapse stops and a protostar appears in the collapsing cloud core. In particular, during the later phase, the period when the protostar significantly increases its mass by mass accretion is called ``the main accretion phase.'' Since the star acquires almost all its mass in this main accretion phase, the final stellar mass is determined in this phase. Protostars during the main accretion phase are observed (or defined) as Class 0 or I objects by spectral energy distribution (SED). The observation of star-forming regions also has shown that almost all Class 0 or I objects drive a protostellar outflow. Because a certain fraction of mass and angular momentum are ejected by the protostellar outflows, they are closely related to the rate of mass accretion onto the protostar and final stellar mass. Moreover, the infalling gas with an angular momentum forms the circumstellar disk during the main accretion phase. The evolution of the circumstellar disk is related to the mass accretion and outflow rates. In addition, the circumstellar disk is the site of planet formation. Therefore, understanding the evolution of the circumstellar disk during the main accretion phase is essential to understanding both the star and planet formation processes. Since the protostar and circumstellar disk during the main accretion phase are veiled by the infalling envelope, it is difficult to observe them. Recently, however, the {\\it Spitzer Space Telescope} (SST) has been unveiling this phase. \\citet{enoch09a} discovered a massive circumstellar disk of $\\sim1\\msun$ comparable to a central protostar around a Class 0 object, indicating that (i) the disk already exists in the main accretion phase, and (ii) the disk mass is significantly larger than the theoretical prediction. \\citet{evans09} observed several star forming regions and identified several Class 0 objects using the SST. They showed that the bolometric luminosity of the objects is considerably dimmer than classical theoretical predictions, and has a dispersion over 2-3 orders of magnitude. They pointed out that recent observations aggravate the ``luminosity problem'' (the problem that the accretion luminosity of protostars is lower than theoretical predictions, see \\citealt{kenyon90}), and concluded that non-steady accretion is inevitably required to explain the observational results. The non-steady accretion may be attainable when the circumstellar disk is sufficiently massive and causes the gravitational instability \\citep{durisen07}. Another recent development in observation of star and planet formation is direct images of exo-planets \\citep{kalas08,marois08,thalmann09}, in which planets are located at $\\gtrsim 10$\\,AU from the central star. However, in the framework of the core accretion scenario \\citep{hayashi85}, there is less likehood of planets forming in a region that is so remote from the central star. Alternatively, the gravitational instability scenario \\citep{cameron78} may explain the formation of such planets. In summary, recent observations seem to indicate that the protostars have a considerably massive disk unlike what was previously believed. Although it seems that recent observational progress is unveiling problems for the early evolution stage of star formation, we cannot directly observe the circumstellar disk and protostar in the main accretion phase because they are embedded in the dense infalling envelope. Thus, theoretical study is necessary to understand the properties of the circumstellar disk and protostar. In particular, multi-dimensional simulations are necessary to investigate the evolution of the circumstellar disk, protostellar outflow, and so on. The evolution of the collapsing gas cloud core from the protostellar core stage (i.e., the gas collapsing phase) until protostar formation has been well investigated using multi-dimensional simulations \\citep[e.g.,][]{bate98,tomisaka02,whitehouse06,stamatellos07,banerjee06,machida06,machida07,machida08a,machida08b,machida09b}. On the other hand, only a few studies have focused on the main accretion phase immediately following the prestellar core stage, because it is difficult to calculate long-term evolution in the main accretion phase with sufficient spatial resolution. In unmagnetized cloud cores, the formation and evolution of the circumstellar disk in the main accretion phase through the gas collapsing phase were investigated by \\citet{walch09a,walch09b}, \\citet{vorobyov10}, and \\citet{machida10}. They found that the circumstellar disk is considerably massive to induce fragmentation or the gravitational instability that is related to a non-steady accretion flow onto the protostar. In reality, however, since molecular clouds are strongly magnetized \\citep{crutcher99}, the magnetic field may play an important role in the evolution of the circumstellar disk during the main accretion phase. \\citet{vorobyov06,vorobyov07} investigated the evolution in two dimensions of the circumstellar disk in a magnetized cloud core and showed the non-steady accretion onto the central protostar. In three dimensions, the formation and evolution of the circumstellar disk from prestellar core stage were investigated only by \\cite{inutsuka09}, in which they showed fragmentation and possible planet formation in the magnetically inactive region of the circumstellar disk during the main accretion phase. They also indicated non-steady mass accretion onto the protostar owing to the gravitational instability of the circumstellar disk. However, this study only calculated the evolution of the circumstellar disk about $\\sim1000$\\,yr after protostar formation. In this study, in a setting similar to \\cite{inutsuka09}, we investigate the evolution of the circumstellar disk for $\\sim10^{4}$\\,yr, which is $\\sim10$ times longer than the previous study. In addition to the model adopted by \\citet{inutsuka09}, we newly calculate the evolution of the circumstellar disk formed in a relatively stable initial cloud core. In both models, we compare the mass accretion rate, properties of protostellar outflow, and the fragmentation condition. The structure of the paper is as follows. The framework of our models and the numerical method are given in \\S 2. The numerical results are presented in \\S 3. We discuss the fragmentation condition of the circumstellar disk and its implication for the planet formation in \\S 4, and we summarize our results in \\S 5. ", "conclusions": "\\subsection{Can We Detect Planets or Substellar-Mass Objects in the Main Accretion Phase?} So far, it is considered that planets form in relatively quiet circumstellar disks after the gas accretion onto the circumstellar disk has almost ended. In this study, however, we showed the possibility of the formation of planets or substellar-mass objects by gravitational instability in a vigorous circumstellar disk during the main accretion phase, in which the circumstellar disk continues to increase its mass by gas accretion from the infalling envelope. When planets or substellar-mass objects appear in the circumstellar disk, they can affect the rate of accretion onto the protostar. As described in \\S\\ref{sec:mdot}, we may verify their existence by observing the time-variability of the rate of accretion onto the central protostar. As seen in Figure~\\ref{fig:9}, the protostar shows a relatively low time-variability of the accretion rate when neither planets nor substellar-mass objects exist in the circumstellar disk, while it shows a high time-variability of the accretion rate when fragmentation occurs and planets or substellar-mass objects appear. Figure~\\ref{fig:13} shows the time variation of accretion luminosity of the protostar, which is defined as \\begin{equation} L_{\\rm ps} = G\\, \\dfrac{M_{\\rm ps}\\, \\dot{M}_{\\rm ps}}{r_{\\rm ps}}, \\end{equation} where we adopted the protostellar mass $M_{\\rm ps}$ and mass accretion rate $\\dot{M}_{\\rm ps}$ as values derived from calculations at each time and used a constant protostellar radius $r_{\\rm ps} = 2\\rsun$ for simplicity. The figure indicates that the protostar for model A05 has an accretion luminosity in the range of $0.01 \\lesssim L_{\\rm ps}/\\lsun \\lesssim 1$ for $\\sim10^4$\\,yr after protostar formation. The accretion luminosity for this model gradually increases for a longer timescale of $\\sim10^4$\\,yr, while it oscillates with a period of $\\sim100$\\,yr, which corresponds to the Kepler timescale in the region of a developing spiral structure. The Kepler timescale $\\tau_{\\rm k}$ can be expressed as \\begin{equation} \\tau_{\\rm k} = 2\\pi \\left( {\\dfrac{r^3}{G M_{\\rm p}}} \\right)^{1/2} \\simeq 100 \\left( \\dfrac{r}{10\\,{\\rm AU}} \\right)^{3/2} \\left( \\dfrac{0.1\\msun}{M_{\\rm ps}} \\right)^{-1/2}\\, {\\rm yr}. \\end{equation} As shown in Figure~\\ref{fig:6}, the non-axisymmetric structure appears at $r\\lesssim 10-40$\\,AU, which corresponds to the magnetic non-active zone \\citep{machida07}. Thus, it is natural that the protostar shows time-variable accretion luminosity with a period of $\\sim100$\\,yr, because the angular momentum transfer that is closely related to the rate of mass accretion onto the protostar is caused by the spiral structure in the circumstellar disk. The variation in amplitude of the accretion luminosity for a short duration ($\\sim100$\\,yr) is about one order of magnitude. On the other hand, the accretion luminosity for model A03 is qualitatively and quantitatively different from that for model A05. The protostar for model A03 repeatedly shows a strong brightening. In this model, the accretion luminosity suddenly increases to reach $\\sim10-10^4\\lsun$ when planets or substellar-mass objects approach the protostar or fall onto the protostar, while the luminosity is lower, $\\lesssim 0.1-0.01\\lsun$ when planets or substellar-mass objects orbit far from the protostar. Thus, the strong brightness may be a clue for detecting planets or sub-stellar-mass objects. In addition, when the planets or substellar-mass objects orbit in the inner disk region of $\\lesssim 10-20$\\,AU, their orbital motion is linked to the accretion rate (or accretion luminosity). The accretion rate and orbital radius for $\\tc=1000-2800$\\,yr for model A03 are plotted in Figure~\\ref{fig:14}. In this time period, the first-generation planet orbits around the protostar for $\\sim10$ times. The figure indicates that the orbital motion of planets or substellar-mass objects clearly synchronizes with the mass accretion rate. The orbiting objects disturb the circumstellar disk and promote mass accretion onto the protostar. Since the circumstellar disk has a higher density and shorter timescale in the proximity of the protostar, a high mass accretion rate is realized when the planets or substellar-mass objects approach the protostar. As a result, the mass accretion rate is closely related to the orbital motion of planets and substellar-mass objects. We may estimate and verify the orbital motion of planets in the circumstellar disk by observing the time-variable mass accretion. Note that we ignored the effect of the infalling envelope when the accretion luminosity is estimated in Figure~\\ref{fig:13}. Note also that, to quantitatively estimate the accretion luminosity, we have to consider the infalling envelope that weakens the luminosity from the protostar. \\subsection{Mass Accretion Rate onto Disk and Masses of Fragments} \\label{sec:mass} In this study, we showed that fragmentation and subsequent formation of planet and sub-stellar mass object can occur in a massive circumstellar disk, while no fragmentation occurs in a relatively less massive disk. The mass of the circumstellar disk or the mass increase rate of the circumstellar disk is related to the initial stability of cloud core. We parameterized the ratio of thermal energy to gravitational energy, $\\alpha_0$, which is related to the mass accretion rate. In general, the mass accretion rate is described as \\begin{equation} \\dot{M} = \\dfrac{M_{\\rm J}}{t_{\\rm ff}} \\propto \\dfrac{c_s^3}{G}, \\end{equation} where $M_{\\rm J}$ is the Jeans mass and $t_{\\rm ff}$ is the freefall timescale. In this study, we made an equilibrium sphere (i.e., critical BE sphere) and increased the density (or mass) by a factor of $f$ ($=0.84/\\alpha_0$, see eq.~[\\ref{eq:alpha}]) to induce the collapse of the cloud core. Thus, the mass accretion rate in our setting can be described as \\begin{equation} \\dot{M} = \\dfrac{M_{\\rm cl}}{t_{\\rm cl,ff}} = \\alpha_0^{-3/2}\\, \\dfrac{c_s^3}{G}, \\end{equation} where the initial mass of cloud core $M_{\\rm cl}$ is proportional to the density enhancement factor $f$ and Jeans mass $M_J$ as $ M_{\\rm cl}\\propto f\\, M_{\\rm J} \\propto \\alpha^{-1}\\, M_{\\rm J}, $ and the freefall timescale of the cloud core $t_{\\rm ff,cl}$ is proportional to $ t_{\\rm ff,cl}\\propto f^{-1/2}\\, t_{\\rm ff,0} \\propto \\alpha^{1/2}\\, t_{\\rm ff,0}, $ where $t_{\\rm ff,0}$ is the freefall timescale of a critical BE sphere. Thus, a more unstable cloud core that has a smaller $\\alpha_0$ has a larger mass accretion rate and tends to show fragmentation due to gravitational instability. In reality, the mass increase rate for the model A03 is about three times higher than that for the model A05. Thus, the fragmentation and subsequent planets or substellar-mass objects may appear in relatively unstable cloud cores. In general, however, the initial clouds are characterized by three parameters, $\\alpha_0$, $\\beta_0$, and $\\gamma_0$, as described in \\S\\ref{sec:setting}. Although we fixed the angular velocity ($\\beta_0$) and magnetic field strength ($\\gamma_0$) of the initial cloud core in this study, fragmentation may occur in more stable cloud cores (i.e., cloud cores with large $\\alpha_0$) with different $\\beta_0$ and $\\gamma_0$. As shown in \\citet{hennebelle08b}, in the collapsing cloud core, the magnetic field suppresses fragmentation while the rotation promotes it. Thus, fragmentation tends to occur in the cloud core with larger $\\beta_0$ and smaller $\\gamma_0$ \\citep{machida08a}. Moreover, even in a cloud core with a very weak magnetic field, fragmentation in the circumstellar disk may be suppressed by the global spiral structure that effectively transports the angular momentum and lowers the surface density (i.e., increases the Toomre $Q$-parameter). On the other hand, when the circumstellar disk is strongly magnetized, the angular momentum is transferred by magnetic effects and the spiral structure hardly develops. Then, the mass accumulates in the inner disk region where Ohmic dissipation significantly weakens the magnetic field, and fragmentation can occur because no global spiral structure develops. Thus, magnetic field and rotation contribute to both promotion and suppression of fragmentation through a complicated process. In this study, we showed the possibility of fragmentation in the main accretion phase. However, to quantitatively determine fragmentation and subsequent formation of planets and substellar-mass objects in the circumstellar disk, we should investigate the evolution of the circumstellar disk from the molecular cloud core in a large parameter space. In the circumstellar disk, fragmentation occurs and planets or substellar-mass objects appear. In this study, we resolved the fragmentation process with sufficient spatial resolution \\citep{truelove97}. However, more spatial resolution and further developed numerical techniques may be necessary to determine the final mass of fragments. At its formation, a fragment's mass almost corresponds to the Jeans mass. In this study, we assumed the barotropic equation of state, in which the minimum Jeans mass is limited to $\\sim5\\mjup$. Thus, fragmentation does not occur with a fragment mass of $M \\lesssim 5\\mjup$. If the circumstellar disk cools to reaches lower temperatures, more less-massive fragments may appear. The barotropic approximation makes it possible to calculate very long-term evolution of the circumstellar disk and to investigate the evolution of the circumstellar disk by the end of the main accretion phase from the prestellar core stage, while radiation hydrodynamic calculations are necessary to qualitatively estimate the fragment mass. In addition, after fragmentation occurs in the circumstellar disk, we need a higher spatial resolution to investigate the acquisition process of the fragments' mass. Recent works about a gas giant planet formation showed that the Hill radius of the protoplanet should be resolved with sufficient spatial resolution for accurate estimate the rate of mass accretion onto the protoplanet \\citep{kley01,dangelo03,bate03,machida10b}. Note that in these calculations, the protoplanet mass was fixed, and the evolution of the circumstellar disk was calculated with a high spatial resolution for a shorter duration after a gas-giant formed. The Hill radius is expressed as \\begin{equation} r_{\\rm H} = \\left( \\dfrac{M_{\\rm p}}{3M_{\\rm ps}} \\right)^{1/3} r_{\\rm p}, \\end{equation} where $r_{\\rm p}$ is the orbital radius of the planet. When the protoplanet orbiting the solar-mass protostar has a mass of $M_{\\rm p}=10^{-3} M_{\\rm ps}$ ($10^{-2} M_{\\rm ps}$) and is located at $r=1$\\,AU (10\\,AU), the Hill radius is $r_{\\rm H} \\simeq 0.07$\\,AU ($1.5$\\,AU). When the Hill radius is resolved with $\\gtrsim10$ cells, we need a minimum spatial resolution (i.e., the size of each cell) of $\\sim0.007$\\,AU ($\\sim0.15$\\,AU). However, calculation with such high resolution needs much computing power, and thus we only calculated the evolution of the circumstellar disk for a very short duration. In reality, the evolution of the circumstellar disk was calculated for only $\\sim1000$\\,yr at best in previous studies, while we calculated the evolution of the circumstellar disk from prestellar core stage until $\\sim 10^4$\\,yr after protostar formation. The final mass of fragments is important to determine whether they are planets or binary companions. To understand both formation and evolution of (gas-giant) planets, we need to calculate the evolution of the circumstellar disk for a longer duration with higher spatial resolution. However, such calculations are impossible even using present supercomputers because a longer temporal calculation conflicts with higher spatial resolution. Thus, in this study, we covered the whole prestellar cloud core and investigated the evolution of the circumstellar disk with moderately coarser spatial resolution $\\Delta \\simeq 0.6$\\,AU, and showed that fragmentation occurs in the circumstellar disk and a planet or substellar-mass object forms. However, since we need a higher spatial resolution to determine the final mass of the fragment, we did not investigate this point in greater depth. In this study, we showed the possibility of the formation of planets or substellar-mass objects in a magnetized disk for the main accretion phase, and the precise mass of such objects should be investigated in future studies. \\subsection{Ohmic Dissipation and Ambipolar Diffusion} \\label{sec:ambipolar} As described in \\S \\ref{sec:property}, the recurrent planet formation and intermittent outflow shown in this paper are closely related to the effect of magnetic field and its dissipation in the circumstellar disk. Figure~\\ref{fig:10} shows that the magnetic field is well-coupled with the neutral gas in the outer disk region ($r_c > 10-100$\\,AU) where a relatively high ionization rate is realized owing to a relatively low gas density. Thus, the angular momentum in such region is effectively transferred by the magnetic braking and protostellar outflow, and the gas can flow into the inner disk region. On the other hand, the magnetic field is not well-coupled with the neutral gas in the inner disk region ($r_c < 10-100$\\,AU) where the gas density is high and a considerably low ionization rate is realized. Thus, the angular momentum transfer due to the magnetic effects is not so effective, and the gas flowing from the outer disk region accumulates in the inner disk region around the boundary between the magnetically active and inactive region at $r_c \\sim 5-50$\\,AU (see, Fig.~\\ref{fig:10}). In that region, the surface density continues to increase and fragmentation tends to occur. In summary, magnetic dissipation is essential in inducing recurrent fragmentation (or planet formation) in the circumstellar disk. In this study, we have considered only Ohmic dissipation as the dissipation mechanism of the magnetic field. This is because Ohmic dissipation is expected to be the most effective mechanism for magnetic field dissipation in high density region where most of the magnetic flux leakage occurs. However, ambipolar diffusion may also contribute to the magnetic flux leakage in some intermediate density region. Recent studies have shown that the Ohmic dissipation is more effective than the ambipolar diffusion in the range of $n\\gtrsim10^{12}\\cm$, while the ambipolar diffusion dominates in the range of $10^{11}\\cm\\lesssim n\\lesssim10^{12}\\cm$ \\citep{tassis07a,tassis07b,tassis07c,kunz10}. In the range of $n\\lesssim 10^{11}\\cm$ the magnetic field is well-coupled with the neutral gas in a magnetically super critical cloud \\citep{nakano02,kunz10}. Thus, it might be important to study how the inclusion of ambipolar diffusion would change the formation and evolution of the circumstellar disk, through the possible contribution in the limited range of density. With strong coupling approximation, the induction equation including the ambipolar diffusion can be written as \\begin{equation} \\dfrac{\\partial \\vect{B}}{\\partial t} = \\nabla \\times (\\vect{v} \\times \\vect{B}) + \\eta_{\\rm OD} \\nabla^2 \\vect{B} + \\nabla \\times \\left[ \\dfrac{\\vect{B}}{4\\pi \\gamma \\rho_n \\rho_i} \\times \\left[ \\vect{B} \\times (\\vect{\\nabla} \\times \\vect{B}) \\right] \\right], \\label{eq:ambipolar} \\end{equation} where $\\rho_n$, $\\rho_i$ and $\\gamma$ are the density of neutral gas and charged particles and a drag coefficient between charged particles and neutrals, respectively (for details, see, e.g., \\citealt{shu92}). The third term of equation~(\\ref{eq:ambipolar}) can be regarded as a non-linear diffusion term by the ambipolar diffusion. When we adopt an approximate relation between neutral and charged particle density, $\\rho_i = C \\rho_n^{1/2}$, where $C$ is a constant, the effective diffusion coefficient of the ambipolar diffusion ($\\eta_{\\rm ad}$) can be roughly estimated as \\begin{equation} \\eta_{\\rm ad} = \\dfrac{B^2}{4\\pi \\gamma C \\rho^{3/2}}. \\label{eq:coef} \\end{equation} In a realistic setting, \\citet{kunz10} showed that diffusion coefficient of the ambipolar diffusion ($\\eta_{\\rm ad}$) is 10-100 times larger than that of the Ohmic dissipation ($\\eta_{\\rm OD}$) in the range of $10^{11}\\cm \\lesssim n \\lesssim 10^{12}\\cm$. Strictly speaking, to include the effect of the ambipolar diffusion, we should solve equation~(\\ref{eq:ambipolar}) together with the equations determining charged particle density. Note that the realistic evolution of charged particle density sensitively depends on total surface area of dust grains, since the recombination of charged particles on the grain surface is very efficient. However, there is a considerable uncertainty in the evolution of dust grain properties, such as the size distribution, in the protostellar collapse process. Thus, our limited knowledge of grain properties make our accurate treatment of magnetic field dissipation impracticable. This is also the reason why we adopted parametric description of the Ohmic diffusion coefficient in our study of magnetic field dissipation \\citep{machida07}. In this paper we take the same approach to the effect of ambipolar diffusion. To roughly estimate the effect of the ambipolar diffusion, we calculate the evolution of the circumstellar disk with an artificially larger resistivity. In equation~(\\ref{eq:reg}), we replace $\\eta_{\\rm OD}$ by $c_\\eta\\, \\eta_{\\rm OD}$ as follows: \\begin{equation} \\dfrac{\\partial \\vect{B}}{\\partial t} = \\nabla \\times (\\vect{v} \\times \\vect{B}) + c_\\eta \\, \\eta_{\\rm OD} \\nabla^2 \\vect{B}. \\label{eq:od2} \\end{equation} We adopt $c_\\eta = 0$ (ideal MHD model), 1 (fiducial model), 10, 100. As a first test, we constructed a non-rotating Bonner-Ebert sphere with a central density of $n\\simeq 10^{3}\\cm$ immersed in the uniform magnetic field of $B_{z,0}=5\\times10^{-5}$\\,G. Then, we calculated the collapse of this cloud with different $c_\\eta$. Figure~\\ref{fig:15} shows the evolution of the magnetic field as a function of the central number density. The magnetic field in each model tracks the same path in the range of $n\\lesssim 10^9\\cm$, while the magnetic field for model with $c_\\eta = 100$ begins to depart from that for ideal MHD model ($c_\\eta=0$) at $n\\simeq10^9\\cm$. The figure indicates that non-ideal MHD effect becomes important at the lower density with the larger $c_\\eta$. When the central number density reaches $n=10^{14}\\cm$, the magnetic field strengths for model with $c_\\eta=0$ and $100$ are $B_z = 7$\\,G and $0.01$\\,G, respectively. Thus, the ideal MHD model has about 100 times stronger magnetic field than model with $c_\\eta=100$ at this epoch. At the same epoch, models with $c_\\eta=1$ and 10 have $B_z= 0.03$\\,G and 0.06\\,G of the magnetic field, respectively. Thus, the difference of the magnetic field strength between $c_\\eta=1$ and $100$ (10) is only factor of 6 (2) . Next, we calculated the evolution of the circumstellar disk with the initial condition same as model A03 but different resistivities $c_\\eta=0$, 10 and 100. The density distribution on the equatorial plane after the circumstellar disk formation for models with $c_\\eta = 0$ and $c_\\eta=100$ are plotted in Figure~\\ref{fig:16}. As seen in this figure, fragmentation does not occur in ideal MHD model ($c_\\eta=0$) until the end of the calculation, while, for model with $c_\\eta=100$, fragmentation occurs at $t_c \\simeq 2550$\\,yr after the protostar formation and two clumps appear in the circumstellar disk. For model with $c_\\eta=10$, the circumstellar disk also shows fragmentation and subsequent planet formation. Note that we have terminated the calculation with larger resistivity earlier, since it takes formidable time to follow the evolution with larger resistivity with our time-explicit integration scheme for diffusion term \\citep{machida07}. Thus, we cannot directly compare the results among models with different resistivities. However, Figure~\\ref{fig:16} indicates that fragmentation occurs even when the magnetic field dissipates in a relatively low-density region. Therefore we expect that recurrent planet formation and intermittent outflow can occur even with more realistic treatment of the ambipolar diffusion. \\subsection{Comparison with Previous Studies} Only a few studies focused on the formation and evolution of the circumstellar disk from the prestellar core stage. \\citet{walch09a} and \\citet{vorobyov10} investigated the circumstellar disk in unmagnetized clouds, and found that massive circumstellar disk formed at the early main accretion phase tends to show fragmentation. \\citet{machida10} and \\citet{bate10} pointed out that the disk forms before the protostar formation, and the circumstellar disk is inevitably massive than the protostar in the early accretion phase. This is because the circumstellar disk is originated from the first adiabatic core formed in the gas-collapsing phase before the protostar formation \\citep{bate98,machida10}. Such massive circumstellar disk is prone to show fragmentation. These studies indicate the possibility of the formation of substellar- or planet-mass object in the unmagnetized circumstellar disk by the gravitational instability. In the magnetized cloud core, only \\citet{vorobyov05,vorobyov06} investigated the formation and evolution of the circumstellar disk using two-dimensional simulation. Although, they adopted the magnetic evolution in an approximate way instead of solving the induction equation, they presented several interesting results. They pointed out that fragmentation and planet formation is possible even in the magnetized circumstellar disk. They also showed that the gravitational instability or fragments temporarily amplifies the mass accretion rate onto the protostar that may account for the FU Orionis outbursts. There are slight quantitative differences between \\citet{vorobyov05,vorobyov06} and ours, because we calculated the circumstellar disk with a realistic process of the magnetic dissipation in three dimensions. However, our results are qualitatively the same as their results. This is because the circumstellar disk always experiences gravitationally unstable phase after the protostar formation. \\citet{inutsuka09} showed that, even in the magnetized cloud core, the circumstellar disk is massive than the protostar and becomes gravitationally unstable. Such disk tends to show fragmentation in the early accretion phase. Even when fragmentation does not occur in the early accretion phase, such massive disk is expected to show fragmentation in the later accretion or subsequent phase. Thus, we expect that the planet formation due to the gravitational instability is common in the star formation process." }, "1101/1101.0396_arXiv.txt": { "abstract": "We study the global evolution of the magnetic field and interstellar medium (ISM) of the barred and ringed galaxies in the presence of non-axisymmetric components of the potential, i.e. the bar and/or the oval perturbations. The magnetohydrodynamical dynamo is driven by cosmic rays (CR), which are continuously supplied to the disk by supernova (SN) remnants. Additionally, weak, dipolar and randomly oriented magnetic field is injected to the galactic disk during SN explosions. To compare our results directly with the observed properties of galaxies we construct realistic maps of high-frequency polarized radio emission. The main result is that CR driven dynamo can amplify weak magnetic fields up to few $\\mu$G within few Gyr in barred and ringed galaxies. What is more, the modelled magnetic field configuration resembles maps of the polarized intensity observed in barred and ringed galaxies. ", "introduction": "To explain the observational properties of the magnetic field in barred and ringed galaxies the dynamo action is necessary. It is thought that the CR driven dynamo can be responsible for the following effects: amplification of galactic magnetic fields up to several $\\mu$G within a lifetime of a few Gyr; large magnetic pitch angles of about $-35$\\textdegree ; symmetry (even, odd); maintenance of the created magnetic fields in a steady state; magnetic field which does not follow the gas distribution, i.e. magnetic fields in NGC 4736 crossing the inner gaseous ring without any change of their direction (Chy\\.zy \\& Buta 2008) or magnetic arms in NGC 1365 which are located between gaseous spiral (Beck et al. 2002). ", "conclusions": "" }, "1101/1101.2446_arXiv.txt": { "abstract": "We present here findings for ${\\rm C}^{18}$O depletion in eight starless cores in Taurus: TMC-2, L1498, L1512, L1489, L1517B, L1521E, L1495A-S, and L1544. We compare observations of the ${\\rm C}^{18}$O J=2-1 transition taken with the ALMA prototype receiver on the Heinrich Hertz Submillimeter Telescope to results of radiative transfer modeling using RATRAN. We use temperature and density profiles calculated from dust continuum radiative transfer models to model the ${\\rm C}^{18}$O emission. We present modeling of three cores, TMC-2, L1489, and L1495A-S, which have not been modeled before and compare our results for the five cores with published models. We find that all of the cores but one, L1521E, are substantially depleted. We also find that varying the temperature profiles of these model cores has a discernable effect, and varying the central density has an even larger effect. We find no trends with depletion radius or depletion fraction with the density or temperature of these cores, suggesting that the physical structure alone is insufficient to fully constrain evolutionary state. We are able to place tighter constraints on the radius at which ${\\rm C}^{18}$O is depleted than the absolute fraction of depletion. As the timeline of chemical depletion depends sensitively on the fraction of depletion, this difficulty in constraining depletion fraction makes comparison with other timescales, such as the free-fall timescale, very difficult. ", "introduction": "Understanding the transformation of a starless core into a protostellar object is the first step in understanding the overall process of low-mass star formation. A key component in understanding low-mass star formation is determining the evolutionary states of starless cores. The competing models for star formation - ambipolar diffusion and gravo-turbulent fragmentation - each have very different timescales \\citep[see, e.g.,][]{Shuetal1987,gravfrag, maclow, Evansreview, hartmannetal}. Ambipolar diffusion posits that collapse of a starless core is slowed by neutrals slipping relative to ions and the decay of magnetic fields \\citep{Shuetal1987}. The ambipolar diffusion timescale is approximately ${5}\\times{10}^{6}$ years, assuming an ionization fraction of ${10}^{-7}$. Gravo-turbulent fragmentation proposes that large-scale turbulence causes fragmentation of individual cores, which then collapse on a free-fall timescale. For ${\\rm H}_{2}$ densities of ${5}\\times{10}^{5}$ ${\\rm cm}^{3}$, that implies a timescale of $5\\times{10}^{4}$ years. One approach is to use chemical evolution, or the process by which the chemistry of a core changes over time, as another piece of information to distinguish between these theories. In the cold ($\\leq$ 10K), well-shielded interiors of dense starless cores, significant freeze-out (depletion) of molecular species such as CO, CS, CCS, and HCO$^+$ are observed (e.g., \\cite{Tafalla416}). The more evolved cores are expected to have more freeze-out, or depletion \\citep{Tafalla416}. For example, \\cite{Crapsi2007} found that in L1544, 93\\% of the ${\\rm C}^{17}$O has frozen out in the center of this very centrally condensed core (central density of $8 \\times 10^5$ cm$^{-3}$). Another example is the work of \\cite{Redmanetal} toward L1689B (central density of ${2.8}\\times{10}^{5}$ cm$^{-3}$ \\citep{Bacmannetal}) where 90\\% of CO is depleted after ${4}\\times{10}^{4}$ years, or half the free-fall timescale at that density. However, this is a lower limit, as they assume a sticking coefficient of unity, and a constant density over time (which, as they point out, is unreasonable - the density would increase over time). The state-of-the-art theoretical studies are chemodynamical models that use a chemical network at each time step during the collapse of a starless cores to calculate abundances. For instance, the models of \\cite{Lee2004} shows significant chemical depletion happening at roughly ${10}^{6}$ years. It is therefore theoretically possible that the depletion profile can be used to distinguish between ambipolar diffusion and gravo-turbulent fragmentation. In this paper, we map the J=2-1 transition of ${\\rm C}^{18}$O toward a sample of eight starless cores in the nearby Taurus molecular cloud. By comparing observed integrated intensities with intensities obtained from radiative transfer modeling, we place constraints on the amount of depletion in terms of how much CO is depleted (fraction of depletion ${f}_{d}$), and where this occurs (the depletion radius ${r}_{d}$ ). Our work is unique because we use a larger sample set (8 cores), all located in the same molecular cloud and are then able to compare the amount of depletion in these cores to evolutionary parameters, such as central density. We do not assume a constant ${\\rm T}_{g}$=10~K profile for our models, as has been done in previous work, and investigate degeneracies in choice of temperature versus choice in abundance. We also investigate possible effects of the gas temperature and the dust temperature not being equal. Unless otherwise specified by ${\\rm T}_{dust}$, temperatures referenced are the kinetic gas temperature,${\\rm T}_{g}$ . This paper is organized as follows. In \\S~2 we discuss our sample selection and motivation. We also describe our observations and modeling methodology, including the implementation of the RATRAN (\\cite{ratran}) radiative transfer code. In \\S~3 we test our modeling procedure, including potential degeneracies of various parameters, and present results for each individual core. In \\S~4 we present our analysis and in \\S~5 we give our conclusions. ", "conclusions": "We find that all of our cores but one, L1521E, are best fit by models showing significant depletion, consistent with previous work in the field suggesting that unevolved cores are rare. This would imply that the depletion process happens very quickly. The lack of unevolved cores is likely an observational bias, as chemically unevolved cores may also be lower density than those in our sample set. Observations with \\textit{Herschel} and SCUBA2 may find many more lower density cores that may also be chemically unevolved. We find that assuming an isothermal temperature profile can reproduce the observations, although one needs further information to know the appropriate temperature to use. This is likely due to the small temperature variations seen in dust continuum radiative transfer models. T=10~K, however, rarely provides a good best match. We find assuming $T_d = T_{gas}$ gives similar results to energetics calculations with $T_d \\neq T_{gas}$. Our work also finds no correlation between the amount of depletion and central density or central temperature, nor do we find correlations between density or temperature at the radius of depletion. It is possible that these cores are evolving at different rates that may be affected by variations in the magnetic field strengths, turbulence, and bounding pressures across Taurus. Constraining the amount of depletion in these cores can be complicated by degeneracies between temperature and canonical abundance and difficulty in placing strong constraints on ${f}_{d}$. Constraining ${f}_{d}$ more strongly is necessary to compare chemical evolution time scales to free-fall time scales, an important step in distinguishing between the relative importance of ambipolar diffusion and gravo-turbulent fragmentation in the evolution of starless cores. Ultimately, observations with better angular resolution, as well as additional understanding of the desorption processes, dynamics, and surrounding environments of these cores, will help complete our understanding of depletion and chemical evolution in starless cores." }, "1101/1101.5099_arXiv.txt": { "abstract": "{ The ground-based Konkoly Blazhko Survey I and II aim to collect accurate, extended, multicolour light curves of galactic field RRab stars in order to determine the incidence rate of the modulation in the sample, to study the modulation in details, to study long-term changes in the modulation properties and to find changes in the mean global physical parameters of the stars with Blazhko phase. Here the diversity of light curve variations of Blazhko stars is demonstrated through a sub-sample of the modulated RRab stars found by the first part of this survey. } ", "introduction": "The cyclic variations in the shape of RR Lyrae light curves, the Blazhko effect, recognized by \\cite{blazhko} and \\cite{shapley}, is one of the most long-lasting riddle of stellar pulsation theory. Recent extended studies of these stars revealed that the phenomenon manifests itself in many different ways, which poses a great challenge against theories trying to explain the effect \\citep{kovacs}. In this paper the Konkoly Blazhko Survey I and II is introduced in brief and the diversity of modulated RRab light curves is demonstrated through a sub-sample of the Blazhko stars found by the survey. ", "conclusions": "" }, "1101/1101.0219_arXiv.txt": { "abstract": "We investigate the hydrostatic equilibrium of stellar structure by taking into account the modified Lan\\'{e}-Emden equation coming out from $f(R)$-gravity. Such an equation is obtained in metric approach by considering the Newtonian limit of $f(R)$-gravity, which gives rise to a modified Poisson equation, and then introducing a relation between pressure and density with polytropic index $n$. The modified equation results an integro-differential equation, which, in the limit $f(R)\\,\\rightarrow\\,R$, becomes the standard Lan\\'{e}-Emden equation. We find the radial profiles of gravitational potential by solving for some values of $n$. The comparison of solutions with those coming from General Relativity shows that they are compatible and physically relevant. ", "introduction": "Extended Theories of Gravity (ETG) \\cite{book} are a new paradigm of modern physics aimed to address several shortcomings coming out in the study of gravitational interaction at ultra-violet and infra-red scales. In particular, instead of introducing unknown fluids, the approach consists in extending General Relativity (GR) by taking into account generic functions of curvature invariants. These functions can be physically motivated and capable of addressing phenomenology at galactic, extragalactic, and cosmological scales \\cite{review}. This viewpoint does not require to find out candidates for dark energy and dark matter at fundamental level (not detected up to now), but takes into account only the observed ingredients (\\emph{i.e.} gravity, radiation and baryonic matter), changing the \\emph{l.h.s.} of the field equations. Despite of this modification, it is in agreement with the spirit of GR since the only request is that the Hilbert-Einstein action should be generalized asking for a gravitational interaction acting, in principle, in different ways at different scales but preserving the robust results of GR at local and Solar System scales (see \\cite{book} for a detailed discussion). This is the case of $f(R)$-gravity which reduces to GR as soon as $f(R)\\,\\rightarrow\\,R$. Other issues as, for example, the observed Pioneer anomaly problem \\cite{anderson} can be framed into the same approach \\cite{bertolami} and then, apart the cosmological dynamics, a systematic analysis of such theories urges at short scales and in the low energy limit. On the other hand, the strong gravity regime \\cite{psaltis} is another way to check the viability of these theories. In general the formation and the evolution of stars can be considered suitable test-beds for Alternative Theories of Gravity. Considering the case of $f(R)$-gravity, divergences stemming from the functional form of $f(R)$ may prevent the existence of relativistic stars in these theories \\cite{briscese}, but thanks to the chameleon mechanism, introduced by Khoury and Weltman \\cite{weltman}, the possible problems jeopardizing the existence of these objects may be avoided \\cite{tsu}. Furthermore, there are also numerical solutions corresponding to static star configurations with strong gravitational fields \\cite{babi} where the choice of the equation of state is crucial for the existence of solutions. Furthermore some observed stellar systems are incompatible with the standard models of stellar structure. We refer to anomalous neutron stars, the so called \"magnetars\" \\cite{mag} with masses larger than their expected Volkoff mass. It seems that, on particular length scales, the gravitational force is larger or smaller than the corresponding GR value. For example, a modification of the Hilbert-Einstein Lagrangian, consisting of $R^2$ terms, enables a major attraction while a $R_{\\alpha\\beta}R^{\\alpha\\beta}$ term gives a repulsive contribution \\cite{stabile_2}. Understanding on which scales the modifications to GR are working or what is the weight of corrections to gravitational potential is a crucial point that could confirm or rule out these extended approaches to gravitational interaction. The plan of paper is the following: In Sec. \\ref{classic}, we review briefly the classical hydrostatic problem for stellar structures. In Sec. \\ref{FEnewtonian} we derive the Newtonian limit of $f(R)$-gravity obtaining the modified Poisson equation. The modified Lan\\'{e}-Emden equation is obtained in Sec. \\ref{LEsection} and its structure is compared with respect to the standard one. In Sec.\\ref{solutions}, we show the analytical solutions of standard Lan\\'{e}-Emden equation and compare them with those obtained perturbatively from $f(R)$-gravity. With help of plot we can compare between them all results. Discussion and conclusions are drawn in Sec. \\ref{conclu} . ", "conclusions": "In this paper the hydrostatic equilibrium of a stellar structure in the framework of $f(R)$-gravity has been considered. The study has been performed starting from the Newtonian limit of $f(R)$-field equations. Since the field equations satisfy in any case the Bianchi identity, we can use the conservation law of energy-momentum tensor. In particular adopting a polytropic equation of state relating the mass density to the pressure, we derive the \\emph{modified Lan\\'{e}-Emden equation} and its solutions for $n\\,=\\,0,\\,1$ which can be compared to the analogous solutions coming from the Newtonian limit of GR. When we consider the limit $f(R)\\,\\rightarrow\\,R$, we obtain the standard hydrostatic equilibrium theory coming from GR. A peculiarity of $f(R)$-gravity is the non-viability of Gauss theorem and then the \\emph{modified Lan\\'{e}-Emden equation} is an integro-differential equation where the mass distribution plays a crucial role. Furthermore the correlation between two points in the star is given by a Yukawa-like term of the corresponding Green function. These solutions have been matched with those coming from GR and the corresponding density radial profiles have been derived. In the case $n\\,=\\,0$, we find an exact solution, while, for $n\\,=\\,1$, we used a perturbative analysis with respect to the solution coming from GR. It is possible to demonstrate that density radial profiles coming from $f(R)$-gravity analytic models and close to those coming from GR are compatible. This result rules out some wrong claims in the literature stating that $f(R)$-gravity is not compatible with self-gravitating systems. Obviously the choice of the free parameter of the theory has to be consistent with boundary conditions and then the solutions are parameterized by a suitable \"wave-length\" $m\\,=\\,\\sqrt{-\\frac{1}{3f''(0)}}$ that should be experimentally fixed. The next step is to derive self-consistent numerical solutions of \\emph{modified Lan\\'{e}-Emden equation} and build up realistic star models where further values of the polytropic index $n$ and other physical parameters, e.g. temperature, opacity, transport of energy, are considered. Interesting cases are the non-relativistic limit ($n\\,=\\,3/2$) and relativistic limit ($n\\,=\\,3$) of completely degenerate gas. These models are a challenging task since, up to now, there is no self-consistent, final explanation for compact objects (e.g. neutron stars) with masses larger than Volkoff mass, while observational evidences widely indicate these objects \\cite{mag}. In fact it is plausible that the gravity manifests itself on different characteristic lengths and also other contributions in the gravitational potential should be considered for these exotic objects. As we have seen above, the gravitational potential well results modified by higher-order corrections in the curvature. In particular, it is possible to show that if we put in the action (\\ref{FOGaction}) other curvature invariants also repulsive contributions can emerge \\cite{stabile_2,cqg}. These situations have to be seriously taken into account in order to address several issues of relativistic astrophysics that seem to be out of the explanation range of the standard theory." }, "1101/1101.2500_arXiv.txt": { "abstract": "Bound orbits have traditionally been assigned to the Large and Small Magellanic Clouds (LMC and SMC, respectively) in order to provide a formation scenario for the Magellanic Stream (MS) and its Leading Arm (LA), two prominent neutral hydrogen (HI) features connected to the LMC and SMC. However, Hubble Space Telescope (HST) measurements of the proper motions of the LMC and SMC have challenged the plausibility of bound orbits, causing the origin of the MS to re-emerge as a contested issue. { We present a new tidal model in which structures resembling the bifurcated MS and elongated LA are able to form in a bound orbit consistent with the HST proper motions. The LMC and SMC have remained bound to each other only recently in our model despite being separately bound to the Milky Way for more than 5 Gyr. We find that the MS and LA are able to form as a consequence of LMC-dominated tidal stripping during the recent dynamical coupling of the LMC and SMC. Our orbital model depends on our assumption that the Milky Way has a constant circular velocity of $V_{\\rm cir}=$ 250 km s$^{-1}$ up to 160 kpc, which implies a massive isothermal halo that is not completely rejected by observations.} ", "introduction": "The Magellanic Stream (MS) is a massive trail of HI gas which lags behind the LMC and SMC in their orbit about the Milky Way (MW). On the leading side of the orbit is a network of HI clumps known collectively as the Leading Arm (LA), and in the region between the LMC and SMC is the Magellanic Bridge. A number of numerical models have provided a dynamical origin for these gaseous features whereby gas is stripped away from the SMC disk by tidal gravitational forces: Murai \\& Fujimoto 1980 (MF80), Gardiner \\& Noguchi 1996 (GN96), Yoshizawa \\& Noguchi 2003, Connors et al. 2006 (C06). In these \\emph{traditional tidal models}, the Magellanic Stream is stripped away $\\sim$1.5 Gyr ago coinciding with a close encounter with the MW. This formation scenario has recently been cast into doubt by the HST proper motion study of Kallivayalil et al. (2006a, 2006b): whereas the traditional tidal models rely on a succession of close encounters between the L/SMC and MW, the HST data seem to imply that such bound orbits are implausible (Besla et al. 2007). This has led to a recent proposal that the MS may have formed within a first passage orbit about the MW (Besla et al 2010, hereafter B10). The conclusion of unbound L/SMC orbits, however, strongly depends on the adopted profile of the MW halo, as demonstrated by Shattow \\& Loeb (2009, hereafter SL09). SL09 revise the MW mass profile of Besla et al. 2007 to agree with new measured values of the MW circular velocity and find that the LMC is most likely bound to the MW with a long ($\\sim$6 Gyr) orbital period { (references for the increased value of $V_{\\rm cir}$ are given in Table 1)}. These conclusions must again be revised if one considers the steep MW mass profile adopted by the tidal models, i.e., an isothermal sphere. Here we present a numerical model based upon the traditional tidal model (e.g., MF80 and GN96) but with two important differences: we incorporate proper motions for the LMC and SMC within 1$\\sigma$ of the HST values, and motivated by SL09 we parametrize the MW isothermal halo with an increased circular velocity of $V_{\\rm cir}=$ 250 km s$^{-1}$. We retain an orbital scenario in which the LMC and SMC remain bound to the MW for more than 5 Gyr, and we furthermore demonstrate that structures resembling the MS, Bridge, and LA can plausibly form within the orbit. Unlike the traditional models, however, the pericentric passage about the MW in our orbit (also occurring at $\\sim$1.5 Gyr ago) does \\emph{not} contribute to the tidal stripping of the SMC disk. Instead, the tidal stripping occurs during the formation of a strong LMC-SMC binary pair $\\sim$1.2 Gyr ago. Because the present model is somewhat idealized { (e.g., neglect of self-gravity, neglect of gas physics, adoption of an isothermal sphere)}, one should regard the present work as a stepping stone between the traditional models of the past and the more robust and realistic models to be developed in the future. { We also stress that the latest proper motion measurements of the LMC, derived from the motions of 3822 stars by Vieira et al. (2010), implies that the LMC's present velocity is $\\sim 340$ km s$^{-1}$ ($\\pm48$ km s$^{-1}$), which is significantly lower than the earlier result ($\\sim 380$ km s$^{-1}$) derived from the motions of only 810 stars by Kallivayalil et al. (2006a). Thus, we may not need to adopt an LMC velocity as high as that of the present model ($\\sim 360$ km s$^{-1}$), but we retain this value in order to demonstrate that the tidal origin of the MS can still be preserved in a bound orbit with a large LMC velocity.} \\begin{table*} \\begin{minipage}[c]{200cm} \\caption{Observational Constraints} \\begin{tabular}{@{}lcp{6cm}} \\hline Parameter & Chosen Value & References \\\\ \\hline \\bf{Milky Way} & & \\\\ Circular Velocity $V_{ \\rm cir }$ (km s$^{-1}$) & 250 & Reid \\& Brunthaler (2004); Uemura et al. (2000); Reid et al. (2009); Sirko et al. (2004) \\\\ Distance to Sun$^\\itl{a}$ $R_{ \\odot }$ (kpc) & 8.5 & Gillessen et al. (2009); Reid et al. (2009) \\\\ Velocity of Sun$^\\itl{a}$ (km s$^{-1}$)& (10.0, 5.2 + $V_{ \\rm cir }$, 7.2) & Dehnen \\& Binney (1998) \\\\ \\hline \\bf{LMC} & & \\\\ Proper Motion ($\\mu_W$, $\\mu_N$) (mas yr$^{-1}$) & (-2.04, 0.48) & Kallivayalil et al. (2006) \\\\ Line-of-Sight Velocity $v_{ \\rm sys}$ (km s$^{-1}$) & 262.2 & van der Marel et al. (2002) \\\\ Position ($\\alpha$, $\\delta$) (degrees) & (81.9, -69.9) & van der Marel et al. (2002) \\\\ Distance Modulus & 18.50 & Freedman et al. (2001) \\\\ \\hline \\bf{SMC} & & \\\\ Proper Motion ($\\mu_W$, $\\mu_N$) (mas yr$^{-1}$) & (-1.31, -1.27) & Kallivayalil et al. (2006) \\\\ Line-of-Sight Velocity $v_{ \\rm sys}$ (km s$^{-1}$) & 146.0 & Harris \\& Zaritsky (2006) \\\\ Position ($\\alpha$, $\\delta$) (degrees) & (13.2, -72.5) & Stanimirovi\\'c et al. (2004) \\\\ Distance Modulus & 18.95 & Cioni et al. (2000) \\\\ \\hline $^\\itl{a}$ Represented in a galactocentric frame (GSR). & & \\\\ \\end{tabular} \\end{minipage} \\end{table*} ", "conclusions": "We find that a simple change of parameters permits a new formation scenario for the Magellanic Stream within the framework of the traditional tidal models. In particular, the MS can suitably form in a bound orbit about the Milky Way via LMC-dominated tidal stripping of the SMC disk. The B10 model supports a similar stripping mechanism, although they adopt an unbound orbit about the MW. Because both bound (this work) and unbound (B10) orbital scenarios appear plausible, the formation of the MS may not supply a strong condition on whether the L/SMC are bound to the Milky Way. Nevertheless, our bound model suggests that the presence of the Milky Way has two important impacts: first, the strong MW tidal field during the most recent MW pericentric passage is able to elongate the LA to $b\\approx30^{\\circ}$; and second, the combined weak tidal fields of the MW and LMC from $t=-3$ Gyr to $t=-1.2$ Gyr are responsible for creating bifurcated filaments within the MS. { Because a first passage orbit relegates the MW to a peripheral role, the B10 model failed to reproduce these two features.} Moreover, if the observed bifurcation of the Magellanic Stream (P03) can indeed be linked to the tidal field of the MW \\emph{prior} to the epoch of stripping, an obvious conflict arises with the first passage scenario. { We unfortunately cannot assess such claims here because the present model is too simple (e.g., test particle method, neglect of gas physics) to accommodate an intricate analysis of the MS. In our subsequent work (Diaz \\& Bekki 2011) we have overcome this simplicity by extending the present model to include self-gravity and a gas-dynamical drag term. The drag is proportional to the ram pressure encountered by the particles of the MS and LA as they orbit through the gaseous MW hot halo (Gardiner 1999). We find that the drag force is able to reduce the total number of particles in the LA and retard their velocities, thereby mitigating the discrepancy with observations for the MS-to-LA mass ratio and the velocity profile of Fig 3, respectively. Additionally, we find that the presence of drag enhances the on-sky bifurcation of the MS by reducing its velocity dispersion. Though this follow-up study is only preliminary, it suggests that our model exhibits a number of promising features which deserve deeper investigation. In particular, a fully hydrodynamical treatment may help illuminate how the LA is able to evolve from a continuous tidal feature into a discrete set of clumpy cloudlets (Bruns et al. 2005).} \\begin{figure} \\includegraphics[trim=0 0 60 0, width=8cm]{f4.jpg} \\caption{Three-dimensional distribution of test particles for the model which evolves from $t=-1.23$ Gyr (left), and for the { fiducial} model which evolves from $t=-3$ Gyr (right). The red and green particles are those within the LMC and SMC, respectively, and the lower and upper features are the MS and LA, respectively. { Structure} within the MS is evident for the model at right and noticeably absent in the model at left. The bounding box has a vertical dimension of 170 kpc.} \\end{figure} \\begin{figure} \\begin{centering} \\includegraphics[width=4cm]{f5.jpg} \\caption{Color-coding which highlights the individual components of the final test particle distribution of the fiducial model. The particles constituting the MS (blue) and the two branches of the LA (green $l>285^{\\circ}$; red $l<285^{\\circ}$) are shown separately. All other particles are colored black.} \\end{centering} \\end{figure} \\begin{figure} \\begin{centering} \\includegraphics[width=8cm]{f6.jpg} \\caption{Color-coding of the undisrupted SMC disk at $t=-1.23$ Gyr, indicating the particles which will eventually constitute the MS (blue) and the two branches of the LA (green $l>285^{\\circ}$; red $l<285^{\\circ}$). All other particles are colored black. The bounding box has a vertical dimension of 16 kpc.} \\end{centering} \\end{figure} Although we chose an isothermal MW halo to purposefully imitate the traditional tidal models, this choice threatens the credibility of our model. The isothermal halo produces a steep mass profile which artificially ensures that the rotation curve of the MW remains flat. These conditions are not very realistic, especially at the large radii (50 kpc to 150 kpc) probed by the LMC orbit about the MW. A more reasonable choice for the MW halo is, for instance, the NFW halo adopted by SL09. As a consequence, the SL09 orbits are more plausible: they are bound but with considerably longer periods about the MW ($\\sim$6 Gyr versus $\\sim$1.5 Gyr in our model). As discussed in B10, the traditional scenario for the formation of the MS (e.g., GN96) is incompatible with these long-period orbits. However, we have proposed a stripping mechanism which is independent of the MW pericenter and in principle does not conflict with the long-period orbits of SL09, i.e., strong tidal forces associated with the formation of a recent LMC-SMC binary pair. In this sense, our proposed MS formation scenario will likely survive even when the present model is replaced with more realistic MW halos and more plausible orbits. While our model and the B10 model both support an LMC-dominated stripping mechanism, there is a clear difference in the supplementary role played by the MW. In our bound model, the MW mediates the binary action of the LMC-SMC pair, guiding them into a recently formed short-period orbit. In the unbound B10 model, the binary action of the LMC and SMC is predetermined as a two-body system in isolation of the MW. The particular LMC-SMC orbit of the B10 model is furthermore curious in that the SMC is presently \\emph{approaching} the LMC. This is in apparent contradiction with the observed LMC-SMC relative velocity (Kallivayalil et al. 2006b). Our orbital model - in particular the formation of the recent LMC-SMC binary pair - may provide insights into a number of outstanding research questions. Harris \\& Zaritsky (2009) have determined that the star formation histories of the LMC and SMC are coupled, with two correlated epochs of star formation occurring $\\sim$2 Gyr ago and $\\sim$0.5 Gyr ago. { Although the timing of the two close LMC-SMC passages in our model are not exactly consistent with the timing of these burst epochs, the present study may well imply that such starbursts could be associated with two strong tidal interactions between the LMC and SMC during their recent dynamical coupling. Our orbital model may also naturally explain why star formation fell into a quiescent epoch from 12 Gyr to 5 Gyr ago (Harris \\& Zaritsky 2009) because the LMC and SMC were separated by large distances and were unable to induce star formation in their respective disks.} The formation of a recent LMC-SMC binary pair has also been proposed by Bekki et al. (2004) to explain the unique distribution of ages among globular clusters of the LMC. { Most of the LMC globular clusters were formed either at ancient times ($\\sim$13 Gyr ago) or at recent times ($< 3$ Gyr ago) (Da Costa 1991). This puzzling $\\sim$10 Gyr age gap was explained by Bekki et al. (2004) by invoking a recent epoch of globular cluster formation induced by strong interactions between the LMC and SMC.} Though the orbital model of Bekki et al. (2004) is inconsistent with HST proper motion data, our present model indicates that a similar orbital scenario is still possible. Drawing upon the framework of the traditional models, we have developed a well-constrained tidal model which reconciles { the latest proper motion data (e.g. Kallivayalil et al. 2006a, 2006b; Vieira et al. 2010)} with the formation of the Magellanic Stream in a bound orbit. Our model furthermore suggests that the LMC and SMC may have been separate entities when they formed, and that many of their observed properties, including the existence of the MS and LA, can be attributed to their recent activity as a binary pair. { Even though the LMC velocity adopted by the traditional tidal models (e.g., $\\sim 300$ km s$^{-1}$ for GN96 and C06) are inconsistent with the HST proper motion data, they are still valid to within 1-$\\sigma$ error of the LMC velocity $\\sim 340 \\pm 48$ km s$^{-1}$ implied by Vieira et al. (2010). Taken together with our present model, these results indicate that bound orbits for the LMC and SMC are still viable in self-consistently explaining the properties of the MS and LA. Furthermore, we have shown that the MW circular velocity adopted by the traditional models $V_{\\rm cir}=$ 220 km s$^{-1}$ must be increased in order to retain a bound orbital history. Ruzicka et al. (2010) come to a similar conclusion, thus providing another piece of evidence that the higher velocities of the LMC and SMC may still be consistent with a bound MS formation scenario. Lastly, these tidal models rely on the key assumption that the MS originated in the SMC disk, and the recent abundance study of Fox et al. (2010) provides a strong observational justification for this assumption: the low metallicity at the tip of the MS is consistent with metallicities observed in the SMC disk but \\emph{not} the LMC.} JD is supported by a Sir Keith Murdoch Fellowship and a SIRF scholarship. We are grateful to the anonymous referee for constructive and useful comments. We thank Lister Staveley-Smith for thoughtful discussions and access to data used in Figs 2 and 3." }, "1101/1101.2208_arXiv.txt": { "abstract": "All globular clusters (GCs) studied to date show evidence for internal (star-to-star) variation in their light element abundances (including Li, C, N, O, F, Na, Mg, Al, and probably He). These variations have been interpreted as evidence for multiple star formation episodes within GCs, with secondary episodes fueled, at least in part, by the ejecta of asymptotic giant branch (AGB) stars from a first generation of stars. A major puzzle emerging from this otherwise plausible scenario is that the fraction of stars associated with the second episode of star formation is observed to be much larger than expected for a standard IMF. The present work investigates this tension by modeling the observed anti-correlation between [Na/Fe] and [O/Fe] for 20 Galactic GCs. If the abundance pattern of the retained AGB ejecta does not depend on GC mass at fixed [Fe/H], then a strong correlation is found between the fraction of current GC stellar mass comprised of pure AGB ejecta, $f_p$, and GC mass. This fraction varies from 0.20 at low masses ($10^{4.5}\\Msun$) to 0.45 at high masses ($10^{6.5}\\Msun$). The fraction of mass associated with pure AGB ejecta is directly related to the total mass of the cluster at birth; the ratio between the initial and present mass in stars can therefore be derived. Assuming a star formation efficiency of 50\\%, the observed Na-O anti-correlations imply that GCs were at least $10-20$ times more massive at birth, a conclusion that is in qualitative agreement with previous work. These factors are lower limits because any mass-loss mechanism that removes first and second generation stars equally will leave $f_p$ unchanged. The mass-dependence of $f_p$ probably arises because lower mass GCs are unable to retain all of the AGB ejecta from the first stellar generation. Recent observations of elemental abundances in intermediate-age LMC clusters are re-interpreted and shown to be consistent with this basic scenario. The small scatter in $f_p$ at fixed GC mass argues strongly that the process responsible for the large mass loss is internal to GCs. A satisfactory explanation of these trends is currently lacking. ", "introduction": "\\label{s:intro} Evidence has been accumulating for the past thirty years that globular clusters (GCs) harbor internal (star-to-star) variation in their light element abundances \\citep[including Li, C, N, O, F, Na, Mg, and Al;][]{Cohen78, Kraft79, Smith82a, Smith83, Kraft94, Gratton04, Pasquini05, Smith05}. And yet, save the most massive GCs, they can still be considered mono-metallic in heavier elements including Ca, Si, and Fe \\citep[e.g.,][]{Carretta09a, Carretta10a}. While early work focused on giant branch stars, abundance variations have now been observed in main sequence turn-off stars \\citep[e.g.,][]{Gratton01, Briley02, Cohen02, Cannon98, Pancino10}, and so the observed variations cannot be attributed to non-canonical mixing in evolved stars\\footnote{Although non-canonical mixing cannot explain the totality of the observed star-to-star variation within GCs, there are well-documented correlations between some elemental abundances, such as carbon and nitrogen, and location along the giant branch \\citep[see e.g.,][]{Smith02b} that cannot be explained by state-of-the-art stellar evolutionary models. This fact significantly complicates interpretation of certain elements.}. Rather, the observed star-to-star variation in elemental abundances must be due to the fact that the stars formed from different material. Perhaps the most striking result emerging from these observations is that the number of stars within a GC that show anomalous abundance ratios is comparable to those stars that have normal ratios. In this context `normal' abundance ratios refers to abundances characteristic of field stars at the same [Fe/H] abundance, and `anomalous' refers to abundance ratios that differ markedly from the field. The comparable number of normal and anomalous stars is observed for all GCs studied to date, spanning a wide range in stellar mass and metallicity \\citep[e.g.,][]{Martell09, Carretta10c}. The anomalous stars display abundance ratios that provide important clues to the source of the raw material from which they formed. These stars show enhanced Na and Al and depleted O and Mg abundances. They are also CN-enhanced and CH-depleted. These peculiar abundance patterns arise naturally when matter is brought to very high temperatures (far exceeding $10^7$ K). At sufficiently high temperatures the CNO, Na-Ne, and Mg-Al nuclear reaction cycles are activated \\citep[the precise temperature required for activation of these cycles depends in detail on the site, whether e.g., the stellar interior or the at the base of the convective envelope;][]{Charbonnel06, Karakas07, Ventura08}. In intermediate-mass asymptotic giant branch (AGB) stars, the base of the convective envelope reaches temperatures that are necessary to ignite these nuclear cycles. Because the envelope is convective, material through the whole envelope can therefore participate in nuclear burning. This process is known as envelope, or hot bottom burning, and it occurs for stars of initial masses in the range $4\\Msun\\lesssim M\\lesssim8\\Msun$ \\citep[e.g.,][]{Renzini81, Ventura08}. Formation of the anomalous abundance ratios within the envelopes of AGB stars is appealing because these stars do not produce heavier elements such as Ca, Si, and Fe, which then explains the lack of observed variation in these elements in most GCs studied to date. These stars also produce large amounts of He, which seems necessary to explain the CMD morphology of many GCs. While AGB stars are plausible candidates for the source of the peculiar abundance patterns, other sites have been proposed. Alternatives include massive ($\\gtrsim 20\\Msun$) rotating stars \\citep{Decressin07} and massive binary stars \\citep{deMink09}. One of the key differences between the massive star and AGB scenarios is the timescales involved. In the massive star scenario the peculiar abundance patterns are created on a timescale comparable to the production of type II supernovae (SNe). This scenario therefore faces two additional difficulties that the AGB scenario does not: 1) how to retain the processed material within the natal GC in the face of energy injection from SNe; and 2) how to create peculiar abundance patterns that do not show evidence for type II SNe products such as Si and Fe. Other arguments against this scenario are discussed in \\citet{Conroy11b}. While the AGB scenario is currently the favored source of the peculiar abundance patterns, it should be noted that this scenario contains significant uncertainties and theoretical difficulties. In light of these considerations, an intricate scenario has emerged to explain the observations. The key ingredient is that multiple generations of star formation have occurred within GCs. Later generations of stars formed from gas enriched by the ejecta of AGB stars associated with earlier generations of star formation. Such a scenario has been recently reviewed and discussed at length by \\citet{Conroy11b}, to which the reader is referred for details \\citep[see also, e.g.,][]{Cottrell81, Smith87, DErcole08, Renzini08, Carretta10c, DErcole10b}. As discussed in that work, the number of star formation events in typical GCs was probably limited to two --- a first generation of stars that formed from material with abundance patterns similar to field stars, and a second generation formed from gas enriched by AGB ejecta from the first generation. The timescale between these two epochs of star formation is probably several $10^8$ yr, which is the timescale for several relevant physical processes including AGB evolution, the time it takes for UV photon production to drop enough to allow the gas to cool and form stars, and the onset of prompt type Ia supernovae. This small difference in age cannot be seen in the main sequence turn-off point in the ancient GCs, but can be observed in younger clusters. Indeed, the massive intermediate-age clusters in the Large Magellanic Cloud (LMC) have a spread in their turn-off point that is consistent within an internal age spread of several $10^8$ yr \\citep{Goudfroij09, Milone09}, suggesting that this scenario occurs at the present epoch as well as the distant past. Another important ingredient in this scenario is that pure AGB ejecta must be mixed with gaseous material that has abundance ratios similar to the first generation \\citep{Prantzos07, Ventura08b, Ventura09, DErcole10b}. This is required to reproduce the observed range in light element abundances, which extends smoothly from the very anomalous to the normal. A natural way to acquire additional material with normal abundance ratios is via accretion from the ambient interstellar medium \\citep[ISM;][]{Pflamm-Altenburg09, DErcole10b, Conroy11b}, although other scenarios are possible \\citep{DErcole08, Gratton10a}. As demonstrated in \\citet{Conroy11b}, significant accretion from the ambient ISM is possible for the physical conditions characteristic of young GCs (i.e., cold dense interstellar media and low relative velocities). Ram pressure is not important except for the lowest masses, where indeed anomalous abundance ratios are not observed, both in the Galactic open clusters \\citep{deSilva09, Martell09, Pancino10b} and the intermediate-age clusters in the LMC \\citep{Conroy11b}. Implicit in this requirement is that the accreted material remain incompletely mixed with the AGB ejecta --- the stars must form from material with a {\\it range} of abundances, not just some average of accreted material and AGB ejecta. The greatest challenge facing this otherwise plausible scenario is in explaining the roughly comparable number of first and second generation stars, which is observed for GCs spanning a wide range in mass. Under the assumption that second generation stars form from AGB ejecta plus a modest amount of accreted material, the ratio of first to second generation stars should be of order 10:1. That is, the standard scenario predicts a number of second generation stars lower by a factor of ten compared to observations. The standard prediction assumes 100\\% star formation efficiency and is based on stellar evolution theory and a canonical initial mass function (IMF), which implies that only $\\sim10$\\% of the mass of a stellar population ends up in AGB ejecta. A variety of solutions to this problem have been proposed, including a different IMF between first and second generation stars \\citep{Smith82b, Prantzos06}, and a substantially larger mass at birth of all GCs \\citep[][]{DAntona04, Bekki06, DAntona07, Decressin07b, DErcole08, Schaerer11}. The second scenario requires that GCs were factors of $10-100$ more massive at birth. If this is correct, it constitutes a dramatic revision of our understanding of GC formation and evolution. This tension provides the motivation for the present analysis. In the present work it is assumed that AGB stars with masses in the range $3-8\\Msun$ contribute to the formation of the second generation of stars. However, the current generation of AGB nucleosynthetic yields favors a narrower range of stellar masses, perhaps only $5-8\\Msun$, that are contributing to form the observed range of abundance patterns of second generation stars \\citep[e.g.,][]{Ventura09, DErcole10b}. Appealing to a smaller range of masses only exacerbates the problem noted above because an even larger initial GC is required to produce the same amount of polluted material (a factor of two larger if one considers only $5-8\\Msun$ AGB stars, rather than $3-8\\Msun$ AGB stars, everything else being equal). Moreover, if only massive AGB stars are allowed to contribute, then the timescale for the formation of second generation stars shrinks to $<10^8$ yr. If a second generation of stars forms within the potential well of a first generation GC, one might expect the second generation stars to be more spatially concentrated than the first generation, at least initially. $N-$body simulations have shown that if they begin more centrally concentrated, second generation stars will remain more concentrated than the first generation for a few central relaxation times \\citep{Decressin08, DErcole08}, after which time the relative number of first and second generation stars becomes constant within the half-mass radius. There is evidence that the second generation is more centrally concentrated than the first in NGC 1851 \\citep{Zoccali09}, $\\omega$Cen \\citep{Pancino03, Sollima07}, and NGC 3201 \\citep{Carretta10e}. The situation for other GCs is less clear \\citep{Carretta09b}, owing to the small numbers of stars observed, incomplete spatial coverage, and the short relaxation times for many GCs. Parallel to the advances in the chemical abundances of normal GCs has been the revelation of distinct stellar populations in the color-magnitude diagrams (CMDs) of the most massive GCs as revealed by the {\\it Hubble Space Telescope} \\citep[see][for a review]{Piotto09}. It is now clear that the massive GCs $\\omega$Cen and NGC 2808 harbor multiple distinct main sequences as seen in their CMDs. Many more GCs harbor distinct sub-giant branches and significant width in their red giant branches. Although it is clear that the spread in the CMDs is related to the elemental abundance variations \\citep{Yong08, Marino08, Carretta09b, Milone10}, a comprehensive understanding of the relation between these two phenomena is currently lacking. In particular, while much attention has been focused on the peculiar CMDs of the most massive GCs, it is far from clear that the features of these GCs are characteristic of the population as a whole. The present work aims to provide a quantitative understanding of the relative frequencies of first and second generation stars, the amount of material accreted from the ambient ISM, and the amount of AGB ejecta required to explain the observed elemental abundance variations, over a factor of 100 in GC mass. This work therefore aims to bridge our understanding of the properties of lower mass clusters with the most peculiar massive GCs. One of the major goals of the present analysis is to answer the following question: ``how much more massive did the progenitors of present day GCs need to be in order to explain the observed abundance variations within GCs?'' ", "conclusions": "\\label{s:disc} \\subsection{Origins \\& Implications} \\label{s:origin} The preceding analysis has revealed several important facts related to the multiple stellar population phenomenon in GCs. Perhaps most important is the result that the fraction of present day GC mass comprised of pure AGB ejecta, $f_p$, is large and strongly correlated with GC mass. The small scatter in $f_p$ at fixed mass suggests that whatever process shapes the elemental abundance trends is driven primarily by internal processes, rather than external ones such as tides from the host galaxy. It is puzzling why $f_p$ is so strongly correlated with mass and yet such a weak function of mass (varying by less than a factor of two over two decades in GC mass). There are at least two possible explanations for the observed correlation between $f_p$ and GC mass. First, the amount of AGB material retained within a GC may decrease with decreasing GC mass. This is plausible, as lower mass GCs have lower escape velocities, and yet the typical velocity of AGB ejecta is independent of GC mass. A second possibility is that GCs retain all of their AGB ejecta but this material is more strongly diluted by pristine material accreted from the ambient ISM in lower mass GCs. Lower mass GCs could more efficiently accrete pristine material if ISM sweeping is the dominant mechanism bringing in new material \\citep{Conroy11b}. The former explanation seems more natural and is therefore preferred. The derived fraction of mass in second generation stars presents one of the most significant outstanding puzzles. It is observed to be $>50$\\% of the total present GC mass and is relatively independent of mass. This fraction is derived in the present work, but has been found also by many authors through different techniques \\citep[e.g.,][]{Smith87, Carretta10c}. If GCs were $>10$ times larger at birth, then the early contribution of second generation stars to the total was small, of order a few percent. The fact that present day GCs always end up with roughly the same fraction of second generation stars strongly suggests that there is some mechanism that drives all clusters toward this state. One scenario that could achieve this is as follows. Imagine that some mechanism causes the first stellar generation to become unbound on the same timescale as the formation of the second generation (see below for an example). The second generation, being more tightly bound, would remain so, and might also be capable of capturing of order its own mass in first generation stars. The second major outstanding puzzle is the more fundamental issue of how and why the ancient GCs were initially so much more massive at birth. Standard dynamical effects cannot explain this, essentially because the relaxation time increases with increasing mass, and yet the mass enhancement factors are larger at higher masses. The only somewhat plausible scenario that has been proposed to explain this is {\\it primordial} mass segregation of the stars in the first stellar generation \\citep[e.g.,][]{DErcole08}. If a significant fraction of high mass stars were born near the cluster center, then as they evolve and die they will carry away a significant fraction of the total binding energy of the young GC, causing it to expand significantly. There is some circumstantial evidence for primordial mass segregation in nearby young clusters whose ages are comparable to their crossing times \\citep[e.g.,][]{Hillenbrand98, deGrijs02, Stolte06, Gennaro11}. If primordial mass segregation were strong enough, it could unbind young GCs on a timescale of $\\lesssim1$ Gyr \\citep{Vesperini09}, or at least lead to a substantial amount mass loss and cluster expansion \\citep{Marks10}. If this were a common feature of GC evolution, then perhaps GCs owe their very survival to the formation of tightly bound second generation stars \\citep{DAntona08b}. These second generation stars, born at the bottom of a deep potential well set by the first generation stars, should be relatively immune to the effects of primordial mass segregation. Detailed simulations will be required to validate this scenario, of which the simulations by \\citep{DErcole08} are an important first step. If confirmed, this would constitute a major revision to our understanding of the survival of massive star clusters. A novel implication of the fact that the ancient GCs were much more massive at birth is their potential contribution to the Galactic stellar halo \\citep[see also][]{Vesperini10, Schaerer11}. There are currently $\\approx150$ known GCs \\citep{Harris96}, and they have an average present mass of $10^{5.5}\\Msun$. If these GCs had on average 10 times more long-lived stars at birth, then in total they would have equaled a mass of $\\approx5\\times10^8\\Msun$. This value is within a factor of a few of the estimated stellar mass of the Galactic stellar halo \\citep[e.g.,][]{Siegel02}, and of the stellar halo in M31 \\citep[e.g.,][]{Kalirai06}. This is the contribution to the stellar halo of GCs that remain bound at the present epoch. \\citet{Martell10} have shown that approximately half of the Milky Way (MW) stellar halo could have formed from GCs that are now disrupted. An understanding of the early evolution of GCs may therefore be required in order to understand the hierarchical formation of the stellar halos around galaxies. If the average GC today was 20 times more massive at birth, then it would have had a mass of $\\approx6\\times10^6\\Msun$, and it would have formed out of a giant molecular cloud (GMC) with a mass of at least $\\sim10^7\\Msun$, assuming $\\esf=0.5$. Such large GMCs should form in the massive, gas-rich protogalactic disks common at high redshift \\citep{Escala08}. Indeed, observations of disk-dominated galaxies at $z\\sim2$ have identified large numbers of super-star forming clumps with masses of order $10^9\\Msun$ \\citep{Genzel08}. Each of these clumps could easily spawn several massive young GCs. Even at the present epoch young clusters have been found with dynamical masses in excess of $10^7\\Msun$ \\citep[e.g.,][]{Maraston04, Bastian06}. While rare at the present epoch, the conditions at high redshift may well have favored the formation of many such massive objects. M31 contains a large number of extended, low surface brightness GCs with half-light radii $>10$ pc \\citep{Huxor05} that have no counterparts in the Galaxy. The origin of these extended GCs is not known. In the context of the present discussion, they could arise from the rapid expansion caused by primordial mass segregation. If for some reason these clusters formed in a much more benign tidal field than the GCs in the Galaxy, then the loosely bound stars may not have been stripped by the present epoch. If this scenario is correct, then these extended GCs are comprised principally of first generation stars, and should contain a small core of second generation stars at their center. Radial gradients of the CN absorption feature in these GCs would be very interesting. The conclusion that GCs were much more massive in the past may also require some revision to dynamical models for the long-term evolution of the GC population \\citep{Gnedin97, Fall01, Marks10}. While the present analysis provides a self-consistent explanation for the variation in the extent of the Na-O anti-correlation from cluster to cluster, complications arise when considering other light elements. In particular, some clusters that show an Na-O anti-correlation do not harbor a corresponding Mg-Al anti-correlation \\citep{Carretta09c}. For example, the clusters NGC 6121 and NGC 6752 have very similar [Fe/H], total mass, and $f_p$ parameters, and yet the former shows no star-to-star variation in either Mg or Al, while the latter shows clear variations in both. The Mg and Al yields depend on AGB mass in a manner different from the Na and O yields \\citep{Ventura08}, and so it is possible that differences in Mg-Al anti-correlations in GCs that contain the same Na-O anti-correlation arise from a difference in the average AGB mass that donates ejecta for later star formation. See \\citet{Carretta09c} for further discussion. \\subsection{Alternative Explanations} A number of alternative explanations have been proposed for various aspects of the standard model considered herein. Several of these will be briefly discussed in this section; the reader is referred to \\citet{Renzini08} and \\citet{Conroy11b} for further discussion. While AGB ejecta is the favored source for the processed material, other proposed sources include winds from massive ($\\gtrsim 20\\Msun$) rotating stars \\citep{Decressin07} and massive binary star interactions \\citep{deMink09}. The most serious objection to these scenarios is that they are associated with massive stars with short lifetimes. It is therefore not natural for the processed material to remain free of supernovae contamination. For example, it is not obvious why the second stellar generation should have the same [Ca/Fe] abundances as the first generation in these scenarios, and yet this is observed to be so \\citep{Carretta10a}. The similar $10^8$ yr timescales for both intermediate-mass AGB stars and the drop in UV photons from the first generation provide a simple framework to understand not only the ancient GCs but also the observed age spread within intermediate-age LMC clusters. Invoking processed material from massive stars requires either a more nuanced or an altogether different mechanism to be at work in the LMC clusters compared to ancient GCs. This seems unappealing on the principle of simplicity. Several authors have considered the possibility that stars with anomalous abundance ratios simply had their surfaces polluted by AGB ejecta \\citep[e.g.,][]{DAntona83, Thoul02}. In this scenario there is only one generation of star formation and the observed range in light element abundances is due to the differing amount of surface pollution from star to star. A major advantage of this scenario is that the GCs need not have been substantially more massive at birth because a much smaller amount of AGB ejecta is needed to cover the surfaces of stars. This scenario is almost certainly ruled out by the lack of strong [O/Fe] and [Na/Fe] abundance variations between the main sequence and red giant branch (RGB). A $1\\Msun$ star on the RGB has a convective envelope comprising approximately 40\\% of its mass, while the same star on the main sequence has a negligible convective envelope. Therefore, this model would predict that as such a star evolves from the main sequence onto the RGB, surface pollution would be heavily diluted as the convective envelope deepens. This is not observed, which either means that the scenario is to be discounted \\citep[as suggested in][]{Cohen02}, or the amount of surface pollution was substantial, so that it could not be diluted even by deep convection. The latter possibility then requires large amount of AGB ejecta, which means that large mass enhancement factors are required. Recently, \\citet{Conroy11b} proposed a scenario for the formation of multiple stellar populations within GCs that contains many of the ingredients listed in the Introduction, e.g., AGB ejecta and matter accreted from the ambient ISM as the source of second generation stars. The principle difference between that model and the \\citet{DErcole08} model is that the former explicitly attempted to avoid the conclusion that the ancient GCs were much more massive at birth. Instead, the large amount of material needed to form the second generation stars came primarily from the ambient ISM. As found in previous work and confirmed herein \\citep[see also][]{DErcole11}, the difficulty with such a scenario in explaining the properties of the ancient MW GCs is that too few anomalous stars are produced. This is because the accreted ambient ISM, which has normal abundance patterns, dominates the mass budget. This does not preclude the possibility that star clusters in other systems formed a second generation of stars from primarily accreted ambient ISM material. Finally, several related scenarios appeal to highly unusual stellar configurations in order to explain the high number of anomalous stars \\citep[e.g.,][]{Bekki06b, Bekki07, Marcolini09}. These scenarios have two important properties in common. First, due to the special configurations required, they cannot be expected to operate at the present epoch, and therefore the multiple stellar population phenomenon observed in intermediate-age LMC clusters requires a second, distinct explanation. Second, they do not require present day GCs to have been substantially more massive at birth. The scenario outlined by \\citet{Bekki06b}, for example, envisions GC formation at the centers of their own dark matter halos. Stars throughout the dark halo are allowed to donate their AGB ejecta to the central regions where the GC is expected to form. Unfortunately, a number of indirect arguments disfavor GC formation at the centers of their own dark halos \\citep[see, e.g.,][for a review]{Conroy11b}, and direct dynamical searches for dark halos around present day isolated GCs strongly disfavor the presence of dark halos \\citep{Baumgardt09, Conroy11a}. Since the deep potential well provided by a dark halo should efficiently retain supernovae ejecta, it is also not clear why GCs in this scenario should not all have large internal variation in heavier elements such as Fe and Ca." }, "1101/1101.2178_arXiv.txt": { "abstract": "In this paper, using 2MASS photometry, we study the mass functions $\\phi(M) = dN/dM \\propto M^{-\\alpha}$ of a sample of nine clusters of ages varying from 4~Myr--1.2~Gyr and Galactocentric distances from 6--12~kpc. We look for evidence of mass segregation in these clusters by tracing the variation in the value of $\\alpha$ in different regions of the cluster as a function of the parameter $\\tau = t_{age}/t_{relax}$ (where $t_{age}$ is the age of the cluster and $t_{relax}$ is the relaxation time of the cluster), Galactocentric distance, age and size of the cluster. The value of $\\alpha$ value increases with age and $\\tau$ and fits straight lines with slopes $m$ and y-intercepts $c$ given by $m=0.40\\pm0.03$, $c=-1.86\\pm0.27$ and $m=0.01\\pm0.001$, $c=-0.85\\pm0.02$, respectively and is a clear indicator of the dynamical processes involved. The confidence level of the Pearson's product-moment correlation of $\\alpha$ with age is 0.76 with p=0.002 and with $\\tau$ is 0.71 with p=0.007. The value of $\\alpha$ also increases with Galactocentric distance, indicating the presence of a larger relative number of low mass stars in clusters at larger Galactocentric distances. We find two clusters, viz. IC~1805 and NGC~1893, with evidence of primordial or early dynamical mass segregation. Implications of primordial mass segregation on the formation of massive stars and recent results supporting early dynamical mass segregation are discussed. ", "introduction": "The distribution of mass amongst the stars born from a parent cloud is described by the initial mass function (IMF). It is a fundamental parameter not only in understanding the basic star formation process, but also in determining the properties and evolution of stellar systems, which are the basic building blocks of galaxies. The IMF estimated for different populations in which the stars can be observed individually show an extraordinary uniformity (Bastian et al. 2010). This uniformity appears to be present for stellar populations including present-day star formation in small molecular clouds, rich and dense massive star-clusters forming in giant clouds and also with old and metal-poor stellar populations that may be dominated by dark matter. The universality, origin and dependence on physical conditions of the IMF is a very active research area and is very crucial to understanding the basic physics of star formation (Kroupa 2002; Bonnell et al. 2007). The evolution of the IMF is influenced by the evolution of individual stars, the redistribution of stars of different masses and the loss of low mass stars by evaporation. Recent studies by (Goodwin and Kouwenhoven 2009) suggest that the same IMF can be derived from different modes of star formation and thus questioned if the IMF is a direct imprint of the star-formation process. Star clusters are an ideal test bed for studies of the IMF as they are a collection of coeval stars formed from the same parent cloud. Hence many uncertainties like reddening, distance, metallicity, etc in determination of stellar masses are minimised. They are suitable for studies on star formation and the dynamics of stellar systems formation and the dynamics of stellar systems (Lynga 1982;Janes and Phelps 1994; Kharchenko et al. 2005; Friel 1995; Bonatto and Bica 2005). The term \u2018ecology of star clusters\u2019, as coined by Heggie (1992), shows the close interplay between stellar dynamics, stellar evolution, the clusters\u2019 stellar content and the dynamics and properties of the host galaxy all which contribute to their structure and evolution. Mass segregation is the distribution of stars according to their masses, leading to the concentration of high mass stars near the centre and the low mass ones away from the centre. This can take place as a result of dynamical interac- tions between stars in young clusters or could be primordial in nature (Bonnell and Davies 1998; Gouliermis et al. 2004; de Marchi et al. 2006; Vesperini 2010; de Grijs et al. 2002, and references therein). For very young clusters, where the age of these clusters is small compared to their relaxation time, the process of dynamical segregation seems less likely, and this timescale argument has been used as evidence that primordial segregation has played a role (Hillenbrand and Hartmann 1998; Bonnell and Davies 1998; Raboud and Mer- milliod 1998). Examples of such clusters with ages less than 5 Myr include: Mon R2 (Carpenter et al. 1997); IC 1805 (Sagar et al. 1988); NGC 1893 (Sharma et al. 2007); NGC 6530 (McNamara and Sekiguchi 1986); NGC 6231 (Raboud and Mermilliod 1998); and the Orion Nebula Cluster (ONC) (Hillenbrand and Hartmann 1998). However, the simulations by Moeckel and Bonnell (2009) show that for such young sys- tems, star formation scenarios predicting primordial mass segregation are inconsistent with observed segregation lev- els. Recent work by Allison et al. (2009, 2010) showed that early mass segregation can be due to dynamical e\ufb00ects even in timescales as short as a Myr, thus not requiring the need of primordial mass segregation. Mass segregation has been studied using the variation of the slope \u03b1 of the mass function (MF) in di\ufb00erent regions of clusters (Bica et al. 2006; Hasan et al. 2008). The steepness of MF in the outer regions of the clusters compared to that of the inner regions, indicates the presence of mass segrega- tion in clusters. In an earlier paper, using the homogeneous data of the Two Micron All Sky Survey (2MASS), Hasan et al. (2008) studied a sample of four young clusters to test if the observed mass segregation is an imprint of the star formation process or is due to the dynamics of the clusters. They found that the observed mass segregation of the sam- ple of young clusters studied, could be explained on the basis of the dynamics. It was found by Bonatto and Bica (2005); Sharma et al. (2008), that the MF slopes (in the outer region as well as the whole cluster) undergo an exponential decay with the evolutionary parameter \u03c4 = tage /trelax and that the evaporation of low-mass members from outer regions of the clusters is not signi\ufb01cant at larger Galactocentric dis- tances of 9 \u2013 10.8 kpc. The parameter \u03c4 is an evolutionary parameter (Bonatto and Bica 2005) which indicates the ex- tent to which the cluster has relaxed. The relaxation time trelax is a characteristic time during which stars in a cluster tend to achieve equipartition of energy and the high mass stars with lesser kinetic energy sink to the core and the low mass stars move to the outer regions of the cluster (Binney and Tremaine 2008). To make inferences based on the properties and fundamental parameters of clusters, it is essential to use homogeneous samples of photometric data, coupled with uniform methods of data analysis. In this paper, we have selected a sample of nine clusters with varying ages, sizes and Galactocentric distances to study mass segregation and the change in $\\alpha$ in clusters in diverse environments. The clusters, viz. NGC~6704, NGC~6005, NGC~6200, NGC~6604, IC~1805, NGC~2286, NGC~2489, NGC~2354 and NGC 1893, are studied using photometric data from the 2MASS (Skrutskie et al. 2006). The 2MASS covers 99.99\\% of the sky in the near-infrared $J$~(1.25~$\\mu$m), $H$~(1.65~$\\mu$m) and $K_{s}$~(2.16~$\\mu$m) bands (henceforth $K_{s}$ shall be referred to as $K$). The 2MASS database has the advantages of being homogeneous, all sky (enabling the study of the outer regions of clusters where the low mass stars dominate) and covering near infrared wavelengths where young clusters can be well observed in their dusty environments. Many papers devoted to the study of clusters using the 2MASS have been presented in the past few years (Bonatto et al. 2006; Bica et al. 2003; Tadross 2008; Dutra et al. 2002) showing the potential of this database. We use the results of Hasan et al 2008) on four clusters and the results of this work on nine clusters to study the dependence of $\\alpha$ on $\\tau$, Galactocentric distance, age and size of the cluster. We study the structures and dynamical states of our sample of clusters and determine their MFs and degree of mass segregation in various regions of the clusters. We construct radial density profiles (RDPs), colour--magnitude diagrams (CMDs), colour--colour diagrams (CCs), luminosity functions (LFs) and MFs. The Galactocentric distance has been calculated based on the IAU-endorsed distance $R_{o}= 8.5$~kpc. The plan of the paper is as follows: Section 2 describes the clusters in our sample and shows the corresponding RDPs and the values obtained for the limiting radii for these clusters. Section 3 describes the method of selecting cluster members and the corresponding values of fundamental parameters obtained. LFs and MFs are described in Section 4 and a comparative study of these clusters is in the concluding Section 5. ", "conclusions": "In this paper, using 2MASS data, we have studied mass segregation in nine clusters in diverse environments to understand their structure and dynamics. The RDPs of the clusters have been plotted (Fig. \\ref{radall}) and the parameters for the clusters such as reddening, distance and age have been determined using isochrone fits (Table \\ref{allpar}). We have also plotted the LFs in the $J$, $H$ and $K$ bands and used the derived mass--luminosity relation to find the MFs using all three bands independently (see Figs \\ref{lfall}--\\ref{mf1893}). Clusters have been divided into three regions: core, inner and outer halo. The $\\alpha$ values have been determined for different regions and the overall clusters as a function of the parameter $\\tau$. We use the change in $\\alpha$ values for different regions to estimate the level of mass segregation of the clusters. The $\\alpha$ values of mass functions of the clusters under study range from 0.17 to 2.79. Figure \\ref{chi} shows the dependence of $\\alpha$ of clusters as a function of various parameters for 13 clusters (9 from this work and 4 from Hasan et al. 2008). Though our sample is small, it is homogeneous, in the sense of photometric data as well as methods of data analysis thus making it a controlled sample. Such studies are not suitable using heterogeneous datasets where unknown biases may be present. \\begin{figure} \\includegraphics[width=8cm,height=6cm]{chi.ps} \\caption{Dependence of $\\alpha$ on various parameters} \\label{chi} \\end{figure} It is interesting to note a very high confidence level in the correlation of $\\alpha$ with age and $\\tau$. As clusters age, they have steeper values of $\\alpha$. The value of $\\alpha$ value increases with age and $\\tau$ and fits straight lines with slopes $m$ and y-intercepts $c$ given by $m=0.40\\pm0.03$, $c=-1.86\\pm0.27$ and $m=0.01\\pm0.001$, $c=-0.85\\pm0.02$, respectively. The increase in the value of $\\alpha$ with age and $\\tau$, is a clear indicator of the dynamical processes involved where mass segregation can be explained by dynamics. The confidence level of the Pearson's product-moment correlation of $\\alpha$ with age is 0.76 with p=0.002 and with $\\tau$ is 0.71 with p=0.007. \\footnote{The p value shows at what level of confidence the null hypothesis (correlation) can be rejected. For example, p=0.05 shows a 95\\% probability that the hypothesis of a correlation is correct.} The value of $\\alpha$ increases with Galactocentric distance, indicating a larger number of low mass stars in clusters at larger Galactocentric distances due to lesser evaporation of stars. The cluster NGC~6704 had an $\\alpha$ value of 1.15$\\pm$0.33 for the overall cluster with an age exceeding 9 times the relaxation time. The cluster has dynamically relaxed, many of the less massive stars have moved to the outer regions of the cluster, some have been lost due to evaporation and hence halo2 has a flatter value of $\\alpha$ compared to halo1. NGC~6005 is an old cluster which has been mass segregated and has high values of $\\alpha$ due to the effect of both dynamics and evolution of stars, in which massive stars have evolved, lost mass and moved to the outer regions of the cluster. In the case of the cluster NGC~6200, the relaxation times for the core and cluster as a whole are 0.7~Myr and 13.8~Myr respectively and the cluster has partially relaxed. The $\\alpha$ value of the core is 1.25$\\pm$0.43 and it shows that the core has a larger number of high mass stars due to relaxation since the cluster has an age of 6.3~Myr ($> t_{relax}$ for the core). However, the inner and outer halos are in the process of relaxation and their $\\alpha$ values are 1.15$\\pm$0.31 and 1.18$\\pm$0.39 respectively. NGC~6604, though young, has an age of 6.3~Myr which exceeds the relaxation times for the core (0.1~Myr) and cluster (3.58~Myr) respectively. Hence, the cluster has relaxed and has $\\alpha$ values 0.53$\\pm$0.51, 0.46$\\pm$0.23, 0.35$\\pm$0.29 and 0.51$\\pm$0.17 for the core, halo1, halo2 and overall cluster respectively. IC~1805 has an age of 4~Myr and the relaxation times for the core and overall cluster are 0.2~Myr and 29~Myr respectively. It already shows mass segregation as earlier reported by Sagar et al 1988. The $\\alpha$ values of the mass function of the cluster, are core (0.59$\\pm$0.17), halo1 (0.88$\\pm$0.14), halo2 (0.68$\\pm$0.02) and overall cluster (0.69$\\pm$0.14). NGC~2286 has an age of 200~Myr which exceeds the core and overall relaxation times of 0.32~Myr and 18~Myr. The $\\alpha$ values of the core, halo1, halo2 and overall cluster are 1.37$\\pm$2.09, 2.79$\\pm$0.33, 1.91$\\pm$0.64 and 2.12$\\pm$0.38 respectively. Mass seggregation process must have taken place, but many of the high mass stars have moved away from the main sequence and have lost mass and the outer halo seems to have lost low mass stars and hence has a flatter $\\alpha$ . . NGC~2489 is an old relaxed cluster and many of the low mass stars from the core have moved to the outer regions of the cluster. This is evident from the flat $\\alpha$ value of the core (0.89$\\pm$0.54) and the larger $\\alpha$ values of halo1, halo2 and overall cluster (1.27$\\pm$0.24, 1.56$\\pm$0.39 and 1.11$\\pm$0.22 respectively). NGC~2354 is an old cluster and most massive stars have evolved away from the main sequence and the halos and overall cluster have similar $\\alpha$ values of 1.56$\\pm$0.36, 1.63$\\pm$0.45 and 1.48$\\pm$0.24 respectively. NGC~1893 is a very young cluster of age 4~Myr which shows signs of overall mass segregation not only in the core which has a relaxation time of 2.4~Myr, but also in the overall cluster whose relaxation time is very large (67~Myr). The $\\alpha$ values for the core, halo1, halo2 and overall cluster are 0.17$\\pm$0.24, 0.69$\\pm$0.19, 0.54$\\pm$0.18 and 0.68$\\pm$0.11 respectively. Of the nine clusters studied, two clusters (IC~1805 and NGC~1893), are too young to be dynamically relaxed and we speculate this as evidence for primordial mass segregation. Mass segregation by birth is a natural expectation because protostars near the density centre of the cluster have more material to accrete. The actual efficiency of this mechanism is still a matter of debate is still a matter of debate (Krumholz et al. 2005; Krumholz and Bonnell 2009). McMillan et al. (2007) presented an alternative scenario for a dynamical origin of early mass segregation in young clusters. Even if the clumps are not initially segregated, if their internal segregation timescale is shorter than the time needed for the clumps to merge, they will segregate through standard two-body relaxation and preserve this segregation after they have merged. The multiscale dynamical evolution of clumpy systems is, in this case, responsible for rapidly leading to mass segregation in young clusters without invoking any mechanism associated with the star-formation process. Recent simulations by the star-formation process. Recent simulations by Allison et al. (2009, 2010) showed that early mass segregation can be due to dynamical effects even in timescales as short as a Myr, thus not requiring the need of primordial mass segre- gation which would violate the universality of the IMF and set constraints on the origin of the IMF. Understanding the origin of mass segregation can also help di\ufb00erentiate between possible models of massive star formation. Do massive stars form in the centres of clusters, or do they migrate there over time due to gravitational in- teractions with other cluster members? In particular, are the masses of the most massive stars set by the mass of the core from which they form (Krumholz and Bonnell 2009) or by competitively accreting mass due to being located at a favourable position in the cluster (Bonnell and Davies 1998; Krumholz et al. 2005; Bonnell and Bate 2006)? Allison et al. (2009) showed that dynamical mass segregation can occur on a few crossing timescales suggests that massive stars could form in relative isolation in large cores and mass segregate later, possibly avoiding the need for competitive accretion as dominant process to form the most massive stars in the centre of a cluster. However, the simulations by Moeckel and Bonnell (2009) show that for such young systems, star formation scenarios predicting general primordial mass seg- regation are inconsistent with observed segregation levels. They found that a star-formation scenario in which only the most massive stars are primordially segregated is consistent with observations, and o\ufb00ers a way to account for compact groups of young, massive stars. Currently we cannot say conclusively if mass segrega- tion is a birth phenomenon (Gouliermis et al. (2004), or whether the more massive stars form anywhere throughout the proto-cluster volume. Star clusters that have already blown out their gas at ages of one to a few Myr are typi- cally mass- segregated (e.g. R136, Orion Nebula Cluster). Assuming primordial mass segregation would imply that massive stars ( > 10M\u2299 ) only form in rich clusters, and reject the possibility they can also form in isolation (see Li et al. (2003); Parker and Goodwin (2007)). A better understanding of the effects of dynamical evolution is required to clearly differentiate between present dynamically derived star cluster properties and those which were imprinted by star-formation processes." }, "1101/1101.2372.txt": { "abstract": "The two pulsating pre-main sequence (PMS) stars V 588 Mon and V 589 Mon were observed by CoRoT for 23.4 days in March 2008 during the Short Run SRa01 and in 2004 and 2006 by MOST for a total of $\\sim$70 days. We present their photometric variability up to 1000 $\\mu$ Hz and down to residual amplitude noise levels of 23 and 10 ppm of the CoRoT data for V 588 Mon and V 589 Mon, respectively. The CoRoT imagette data as well as the two MOST data sets allowed for detailed frequency analyses using Period04 and SigSpec. We confirm all previously identified frequencies, improve the known pulsation spectra to a total of 21 frequencies for V 588 Mon and 37 for V 589 Mon and compare them to our PMS model predictions. No model oscillation spectrum with $l$ = 0, 1, 2, and 3 $p$-modes matches all the observed frequencies. When rotation is included we find that the rotationally split modes of the slower rotating star, V 589 Mon, are addressable via perturbative methods while for the more rapidly rotating star, V 588 Mon, they are not and, consequently, will require more sophisticated modeling. The high precision of the CoRoT data allowed us to investigate the large density of frequencies found in the region from 0 to 300 $\\mu$Hz. The presence of granulation appears to be a more attractive explanation than the excitation of high-degree modes. %presence of granulation and its contribution to the background noise. This concept appears to us to be an extremely attractive model explaining the large density of frequencies found in photometric time series, which would be very difficult to interpret in terms of disk-integrated high degree ($l \\geq 4$) modes. Granulation was modeled with a superposition of white noise, a sum of Lorentzian-like functions and a Gaussian. Our analysis clearly illustrates the need for a more sophisticated granulation model. ", "introduction": "V 588 Mon (HD 261331, NGC 2264 2) and V 589 Mon (HD 261446, NGC 2264 20) are two pre-main-sequence (PMS) pulsating stars for which there exists strong evidence that they are members of the young open cluster NGC 2264. The proper motions for both stars are in agreement with the cluster\u00d5s average proper motion \\citep{hog00} and both fit the cluster\u00d5s HR- and color-magnitude diagrams well. A radial velocity measurement only exists for V 589 Mon \\citep{str71} but it is consistent with the values for other cluster members. %The pulsation periods of the two stars are too short to be identified as more luminous (i.e., background) $\\delta$ Scuti stars. Finally, the radial velocities of emission lines in the optical spectra of the two stars caused by interstellar gas match the radial velocities of the cluster suggesting that both stars are indeed embedded in gas clouds belonging to the cluster \\citep{kal08}. %{\\bf No detailed investigation about the two stars' memberships to the cluster is available in the literature up to now. But several indications exist that V 588 Mon and V 589 Mon are indeed members of NGC 2264. A radial velocity measurement only exists for V 589 Mon \\citep{str71} and is consistent with the values for other cluster members. The proper motions for both stars are in agreement with the cluster's average proper motion \\citep{hog00} and both fit the cluster's HR- and color-magnitude diagrams well. Their pulsation periods are too short if they would be background and hence more luminous $\\delta$ Scuti stars. Additionally, emission lines caused by interstellar gas with similar radial velocities indicate that both stars are embedded in the clouds of gas belonging to the cluster \\citep{kal08}. Therefore, there is strong evidence that V 588 Mon and V 589 Mon are indeed members of NGC 2264.} The cluster has a diameter of $\\sim$39 arcminutes and belongs to the Mon OB 1 association. The age of NGC 2264 is reported in the literature to lie between 3 \\citep[e.g.,][]{wal56,sun04} and 10 million years \\citep[e.g.,][]{sag86}. With such a young age, the cluster's main sequence only consists of massive O and B type stars, while stars of later spectral types are still in their pre-main sequence (PMS) phase. Therefore, V 588 Mon and V 589 Mon having spectral types of A7 and F2, respectively, have not arrived on the zero-age main sequence (ZAMS) yet. The \\dsct like variability of V 588 Mon and V 589 Mon was first reported by \\citet{bre72}. Hence, they were the first pulsating PMS stars discovered. In the meantime the number of known \\dsct like PMS pulsators has increased from 36 \\citep{zwi08} to $\\sim$ 60 (Zwintz, K., private communication) due to dedicated observations from ground and from space. Pulsating PMS stars have intermediate masses, i.e. between $\\sim$1.5 and 4 M$_{\\odot}$ and can become vibrationally unstable when they cross the instability region in the Hertzsprung-Russell (HR) diagram on their way to the ZAMS. Pre- and post-main sequence evolutionary tracks for the same stellar mass intersect several times close to the ZAMS which makes the determination of the evolutionary stage of a field star from only its effective temperature, luminosity and mass ambiguous. Additional information such as typical observational evidence for the PMS evolutionary stage (i.e. emission lines, IR excess, an X-ray flux, located in an obscured region on the sky etc.) or membership to very young open clusters is needed to resolve this ambiguity. Another way to distinguish the evolutionary stages can come from the asteroseismic interpretation of the observed pulsation frequencies \\citep{gue07}. Time series photometry for V 588 Mon and V 589 Mon has been obtained from a multi-site ground based campaign \\citep{kal08} and from two observing runs of the Canadian MOST space telescope \\citep{wal03} in 2004 and 2006 \\citep{gue09}. The 8 and 12 frequencies common to these three data sets for V 588 Mon and V 589 Mon, respectively, were submitted to a first asteroseismic analysis \\citep{gue09}. The accuracy of the MOST observations is higher than that of the ground based data. Hence it is not surprising that the two MOST data sets yield more significant frequencies at lower amplitudes that were not found in the ground based observations. The CoRoT observations of the two PMS pulsators are new and independent data sets of unprecedented accuracy that allow us for the first time to investigate other effects (e.g., granulation) in PMS stars. As both MOST data sets and the CoRoT data are available to us, we use them together for a detailed pulsational analysis. ", "conclusions": "The two pulsating PMS stars V 588 Mon and V 589 Mon were observed by MOST in 2004 and 2006 for in total $\\sim$70 days and in 2008 by the CoRoT satellite for 23.4 days during the Short Run SRa01. Detailed frequency analyses were conducted for all available data sets and compared to each other. Our analysis illustrates that the frequency dependent intrinsic background signal, i.e. granulation, can explain a large number of significant peaks detected in PMS stars. Granulation modeling was conducted using a first order model where two Lorentzian-like functions, white noise and a Gaussian are combined. The resulting number of frequencies is more consistent with the expected number of low-degree $p$-modes observed in integrated light. But it was also shown that the shape of the pulsational power excess is very likely to be more complicated than a Gaussian and that a more sophisticated approach would be needed in the future. This effect is illustrated by the fact that for V 588 Mon 7 frequencies at low amplitudes would have been suppressed by the background noise model, but appear significantly in all three satellite data sets, i.e. in data from MOST 2004, MOST 2006 and CoRoT. If frequencies are stable over a period of 4 years, they are unlikely to be caused by granulation. After a comparison of the independent analyses of the MOST 2004 and MOST 2006 data sets to the CoRoT data, for V 588 Mon and V 589 Mon 21 and 37 frequencies, respectively, can be attributed to pulsation, among those the 8 and 12 previously published frequencies \\citep{gue09}. %Summary from David Even after granulation filtering, both V 588 Mon and V 589 Mon have more frequencies observed than can be accounted for by $l$ = 1, 2, and 3 $p$-modes. We investigated the possibility that the extra frequencies are rotational split modes. For V 589 Mon the autocorrelation of the observed frequency spectrum and our model searches support the notion that rotational split frequencies, with splittings between 4 $\\mu$Hz and 5 $\\mu$Hz, are present. We are able to match more than half of the observed frequencies with model fits. For V 588 Mon the autocorrelation plot and our attempts to fit model spectra that included splittings were unsuccessful. We believe this is because V 588 Mon's rotation rate is high enough (as implied by its \\vsini) that its splittings are in the nonlinear regime as described by \\citet{esp04}, and further studied by \\citet{oua08}, \\citet{rees09b}, and \\citet{deu10}. How can we obtain unambiguous mode identifications for V 588 Mon and V 589 Mon when we do not know the interior rotation curve? Is inversion possible, i.e., determining the rotation curve from the splittings, simultaneously with identifying the modes? Unlike the slowly rotating Sun, where the splittings are in the linear regime, V 588 Mon and V 589 Mon are rotating rapidly enough that non-perturbative models are probably necessary, especially for V 588 Mon. Rather than to continue with more sophisticated models we are currently pursuing a different approach in which we apply a Bayesean approach that incorporates model independent prior knowledge about the split and unsplit modes to estimate model fit likelihoods. We will report on our efforts in a future paper. %We know that we cannot hope to fully model and interpret the oscillation spectrum of these two stars until we have successfully identified the $n$ and $l$ of each observed frequency. Needless to say with the relatively unconstrained possibilities for interior rotation curves, this task is nontrivial. We believe the best strategy is to first identify the non-rotationally split frequencies ($m = 0$) and then match up the split frequencies ($m = \\pm 1, \\pm 2$, etc.) with the unsplit frequencies. We are attacking this problem from several directions (e.g., Bayesian approach to the frequency extraction that includes priors allowing for rotational splitting; 2D stellar modeling of nonlinear rotational splittings) and will report on our efforts in a future paper." }, "1101/1101.2664_arXiv.txt": { "abstract": "{We analyze the impact of the choice rotation law on equilibrium sequences of relativistic differentially-rotating neutron stars in axisymmetry. The maximum allowed mass for each model is strongly affected by the distribution of angular velocity along the radial direction and by the consequent degree of differential rotation. In order to study the wide parameter space implied by the choice of rotation law, we introduce a functional form that generalizes the so called ``j-const. law\" adopted in all previous work. Using this new rotation law we reproduce the angular velocity profile of differentially-rotating remnants from the coalescence of binary neutron stars in various 3-dimensional dynamical simulations. We compute equilibrium sequences of differentially rotating stars with a polytropic equation of state starting from the spherically symmetric static case. By analyzing the sequences at constant ratio, $T/|W|$, of rotational kinetic energy to gravitational binding energy, we find that the parameters that best describe the binary neutron star remnants cannot produce equilibrium configurations with values of $T/|W|$ that exceed $0.14$, the criterion for the onset of the \\emph{secular} instability. } ", "introduction": "A neutron star (NS), during most of its life, is considered to be a stationary and rigidly rotating object, apart from a tiny lag between the rotation of the superfluid component and that of the normal fluid and the crust (e.g. \\citet{Baym1969} and \\citet{Pines1985}). In fact a nascent neutron star which is born in a supernova event is likely to rotate differentially at first before its angular velocity distribution evolves toward a uniform rotation. There are several mechanisms that account for the redistribution of the angular momentum. One such mechanism is the shear viscosity of neutron matter \\citep{Sawyer1989}. An estimate of the timescale in which the viscosity damps a neutron star's internal shear motion has received a lot of attention in the literature (e.g. \\citet{Cutler1987}). For a typical neutron star we expect a timescale of $10-100$ yrs. In the presence of magnetic fields, the magnetic braking \\citep{Spruit1999} or the magnetorotational instability (MRI) \\citep{Balbus1991} may drastically accelerate the redistribution of angular momentum to the order of 1s \\citep{Duez2006}. Differential rotation plays an important role in the beginning and at the end of the life of a NS. In ots early life, strong differential rotation of a massive core in a supernova may affect the collapse and bounce dynamics \\citep{Dimmelmeier2002, Ott2004}. A newly-born neutron star with strong differential rotation may lose its stability in the course of the rotational (and thermal) evolution, and may collapse to a black hole or other type of compact star (quark star, hybrid star). This may give a unique neutrino signal and a strong gravitational wave emission from the supernova explosion of massive stars (see, e.g., \\citet{Ott2007} and \\citet{Kotake2010} for recent studies on gravitational wave mechanism from stellar core collapse). An exciting way in which a neutron star's life can end is in the inspiral and consequent merger with a binary companion. Recent results of numerical 3-dimensional simulations of neutron star binary mergers (e.g., \\citet{Ruffert1996}; \\citet{Ruffert2001}; \\citet{ShibataUryu2002}; \\citet{Shibata2006}; \\citet{Anderson2008}; \\citet{Baiotti2008}; \\citet{Giacomazzo2010}) have shown that there are cases in which the remnant of the merger has a significantly high degree of differential rotation such that it can sustain a total mass considerably larger than that of a uniformly rotating star \\citep{Baiotti2008}. These hyper-massive neutron stars (HMNSs) remain stable over many dynamical timescales before collapsing to a BH (e.g. \\citet{ShibataUryu2002}, \\citet{Baiotti2008}). In order to study the effect of the centrifugal force in supporting the HMNS, most of these papers show the angular velocity profile of the merged object before the eventual collapse to a black hole. The typical profile extracted from simulations characteristically shows a plateau near the rotation axis and a certain distance from the axis shows a nearly power law behavior. The power index seems to differ from simulation to simulation, but generally does not agree with the so-called ``j-const. law\" \\citep{EriguchiMueller1985} of rotation, which has been so far the only choice available to construct equilibrium models of differentially rotating stars. Models of rapidly rotating stars in general relativity have been studied since the 70s, when large numerical computing facilities became available. The centrifugal deformation and general relativistic gravity make these investigations fully reliant on numerical methods. Since the pioneering work of \\citet{Ipser1976}, these studies have included progressively more sophisticated aspects, such as nuclear matter EOS, degrees of differential rotation, magnetic fields and quite recently meridional flows \\citep{Birkl_etal2010}. We here name the following references of such studies: \\citet{Ipser1976} ; \\citet{Komatsu1989a} ; \\citet{Bonazzola1993} ; \\citet{StergioulasFriedman1995} ; \\citet{Shapiro2000} ; \\citet{Ansorg_etal2002}. For more comprehensive collection of literatures, see \\citet{Stergioulas_LivingReview}. However, it is surprising to note that all of these studies have used just a single kind of ansatz on the rotational angular velocity profile (namely the ``j-const. law\", which include uniformly-rotating stars as a limiting case (Section \\ref{subsec:rotation profile})). In this work we introduce a new parametrized functional form of the rotation profile that enables us to investigate a broader class of differentially rotating stars. In particular, our new rotation law reproduces different power laws in the outer envelope of the angular velocity distribution (see Section\\ref{subsec:rotation profile}) with which it is possible to match the profile of the HMNS . To analyze the impact of the law of rotation on the maximum allowed mass for the differentially-rotating neutron star we construct equilibrium sequences by imposing a fixed value of $T/|W|$ close to what is the classical limit for the onset of the \\emph{secular} instability; rotating neutron stars are known to be destabilized by the effect of dissipative mechanisms like viscosity and gravitational radiation reaction (Chandrasekhar-Friedman-Schutz (CFS) instability; \\citet{Chandrasekhar1970,FriedmanSchutz1978}). We show that the classical secular instability criterion is satisfied only for a limited class of the rotation profiles. The paper is organized as follows. In Section \\ref{sec::formulation} we briefly review the formalism of the equations and the numerical method used to solve them. Then we describe the new functional form of the rotation profile. In Section \\ref{sec::results} we construct equilibrium sequences of rotating stars using the new rotation profile. Discussions are given in the final section. ", "conclusions": "We have introduced a new rotation profile to study equilibrium sequences of differentially rotating relativistic stars. Compared with the previous studies, which assume only one type of rotation profile, we are now able to investigate a broader class of rotating stars. As a first step towards systematic studies, we focus on the simplest neutron star model with a polytropic EOS with index $N=1$. For the rotation profiles that allow analytic expressions in the Eq. (\\ref{hydrostatic eq}), we computed sequences of spheroidal equilibria that start from non-rotating stars. Special attention is paid to the appearance of the critical point to the secular instability (CFS or viscous instability) to bar-shaped deformation of the star. We see that the appearance of the critical point may occur for rather strong differential rotation in the case of $\\alpha=-1$ and $-2$ in Eq.(\\ref{eq: functional form of g}), which corresponds respectively to so-called ``j-const.\" and ``v-const.\" rotation profiles in Newtonian stars \\citep{EriguchiMueller1985}. Even in these cases, the critical point does not appear before the configuration changes from a spheroidal topology to a toroidal one, unless the parameter $R_0$ in Eq.(\\ref{eq: functional form of g}) is sufficiently small. When $\\alpha=-4$ instead there seems no critical point of CFS instability before the equilibrium sequences reach their mass-shedding limit. This last case is relevant since it appears to mimic the approximate rotation profile of some of the quasi-stationary HMNSs seen in the numerical simulations of neutron star mergers. However these stars may be susceptible to the relatively new class of dynamical instability (so-called \"low $T/|W|$ instability\", e.g. \\citet{Centrella_etal2001, Shibata_etal2002, Watts_etal2003, SaijoYoshida2006, OuTohline2006, Baiotti2008}), whose existence relies on the strong shear flow (\\citet{Watts2005,Corvino2010}). Indeed, numerical simulations of binary NSs do indicate that the HMNS develops a bar deformation (e.g.\\citet{ShibataUryu2002}; \\citet{Baiotti2008}) on top of the background star with approximate axisymmetry. It would be interesting to study the appearance of low $T/|W|$ instability by using our equilibrium stars with different rotation profiles since the survival time of the merger remnant could be an important diagnostic. The observation of the delay between detection of the gravitational wave signal form a binary neutron stars merger and the possible observation of the short gamma ray burst counterpart would then help to put constraints on the internal structure of a NS." }, "1101/1101.4797_arXiv.txt": { "abstract": "Inspired by holographic principle, we suggest that the density of dark energy is proportional to the spatial Ricci scalar curvature (SRDE). Such model is phenomenologically viable. The best fit values of its parameters at $68\\%$ confidence level are found to be: $\\Omega_{\\rm m0}=0.259\\pm0.016$ and $\\alpha=0.261\\pm0.0122$, constrained from the Union+CFA3 sample of 397 SNIa and the BAO measurement. We find the equation of state of SRDE crosses $-1$ at $z\\simeq-0.14$. The present values of the deceleration parameter $q(z)$ for SRDE is found to be $q_{z=0}\\sim -0.85$. The phase transition from deceleration to acceleration of the Universe for SRDE occurs at the redshift $z_{q=0}\\sim 0.4$. After studying on the perturbation of each component of the Universe, we show that the matter power spectra and cosmic microwave background temperature anisotropy is slightly affected by SRDE, compared with $\\Lambda$CDM. ", "introduction": "In the last decade a convergence of independent cosmological observations suggested that the Universe is experiencing accelerated expansion. An unknown energy component, dubbed as dark energy, may be responsible for this phenomenon. Numerous dynamical dark energy models have been proposed in the literature, such as quintessence \\cite{quint,Zlatev}, phantom \\cite{phantom,Carroll2003}, $k$-essence \\cite{k,Gao2010,Yang2009,Yang2008a,yang10,yang10a}, tachyon \\cite{tachyonic,Bagla2003,Keresztes,yang}, DGP \\cite{Dvali,zhang}, Chaplygin gas \\cite{Chaplygin,Bento2002,Gorini05}, ect. However, the simplest and most theoretically appealing candidate of dark energy is the vacuum energy (or the cosmological constant $\\Lambda$) with a constant equation of state (EoS) parameter $w=-1$. This scenario is in general agreement with the current astronomical observations, but has difficulties to reconcile the small observational value of dark energy density with estimates from quantum field theories; this is the cosmological constant problem (see e. g. \\cite{weinberg,Sahni2000,Carroll2001,Copeland2006}). Recently it was shown that $\\Lambda$CDM model may also suffer from an age problem \\cite{Yang2010a}. According to the holographic principle \\cite{holography}, a method may help solve this problem was proposed that an unknown vacuum energy with density proportional to the Hubble scale, $\\Lambda\\propto l^{-2}\\sim H^2$, could be present \\cite{horova,Minic2000,Thomas2002}, here $l$ is a characteristic length. Unfortunately, the EoS for such vacuum energy is zero and the Universe is decelerating. Alternatively, the particle horizon can be introduced as the characteristic length $l$. However, the EoS is greater than $-1/3$ in this case; so it still could not explain the observed acceleration of the Universe \\cite{Fischler1998, Bousso1999,Hsu2004, Li2004}. In view of this, the future event horizon was proposed as the characteristic length $l$. This holographic dark energy model and its interacting version are successful in fitting the current observations \\cite{refHDE,Myung2009,Setare2009,Jamil2009,Mohseni2009, Guberina07,Elizalde2005,Karwan2008,Medved2009,Nayak2009,Cruz2008,Bisabr2009}. However, this holographic dark energy model has suffered some serious conceptual problems. Firstly, the present value of dark energy is determined by the future evolution of the Universe; this poses a challenge to causality. Secondly, the density of dark energy is a local quantity, while the future event horizon is a global concept of space-time. To avoid these short-comes, another possibility is considered: the characteristic length $l$ is given by the average radius of Ricci scalar curvature $R^{-1/2}$, in other words, the density of dark energy is proportional to the Ricci scalar curvature, $\\rho_{\\rm X}\\propto R$ (hereafter RDE for short) \\cite{Gao}. Recent studies on this model see e. g. \\cite{Granda,Kim08,Chen2009,Saridakis2008}. However, as holographic principle asserting \\cite{holography}, the number of possible states of a region of space is the same as that of binary degrees of freedom distributed on the boundary of the region. To reflect holographic principle more properly, we propose that the density of dark energy is proportional to the spatial Ricci scalar curvature, $\\rho_{\\rm X}\\propto R_{\\rm s}$ (hereafter SRDE for short), instead of Ricci scalar curvature. We find SRDE not only has the same properties, but also has properties which RDE doesn't have. This paper is organized as follows. In section II, we describe the SRDE model. Observations constraints and cosmic expansion will be discussed in section III. We discuss structure formation in section IV. Finally, conclusions and discussions are made in section V. ", "conclusions": "In this article, we have shown that replace Ricci scalar curvature in RDE model with spatial Ricci scalar curvature, $\\rho_X\\propto R_{\\rm s}$, the resulting SRDE model is viable phenomenologically. The dark energy component in SRDE contains a term evolves like nonrelativistic matter, a term increases with decreasing redshift, and a term serves as dark radiation which RDE doesn't have. Like RDE, SRDE avoided the casuality problem of holographic dark energy, because the dark energy is determined by the locally determined spatial Ricci scalar curvature, not the future event horizon. Because SRDE model is not associated Planck or other high energy physics scale, but with the size of space-time curvature, the fine tuning problem can be avoided. Assuming that the Universe is spatially flat, we have placed observational constraints on SRDE scenario with the Union+CFA3 sample of 397 SNIa and the BAO measurement from the SDSS. For SRDE, we have obtained the best fit values of the parameters at $68\\%$ confidence level: $\\Omega_{\\rm m0}=0.259\\pm0.016$ and $\\alpha=0.261\\pm0.0122$ with $\\chi^2_{\\rm min}=470.007$ (dof$=1.187$). Withe these values of parameters, we have studied evolutions of Hubble parameters, parameters of EoS and the deceleration parameters. SRDE evolves like nonrelativistic matter when $z\\longrightarrow2$, while RDE not. The present values of the deceleration parameter $q(z)$ for SRDE is found to be $q_{z=0}\\sim -0.85$. The phase transition from deceleration to acceleration of the Universe for SRDE occurs at the redshift $z_{q=0}\\sim 0.4$, comparable with that estimated from 157 gold data ($z_{\\rm t}\\simeq0.46\\pm0.13$) \\cite{Rie04}. We found the equation of state of SRDE crosses $-1$ at $z\\simeq-0.14$. We have discussed the dark matter/baryon density perturbations, the differences between SRDE and the $\\Lambda$CDM model are very small for the large scale perturbations, while a little larger for the small scale perturbations. As expected, due to the extra dust-like component, the growth rate of SRDE differs from that in $\\Lambda$CDM model, and the matter-radiation equality occurred at smaller $a_{\\rm eq}$. We also have shown that the CMB angular power spectra of SRDE is consistent with that of $\\Lambda$CDM on the large scale, while difference slightly on the small scale. This mean we could test SRDE from small scale data. The properties of SRDE, the presented dark radiation like as in brane dark energy and in Ho\\u{r}ava-Lifshitz cosmology and the behavior crossing $-1$ like as in quintom, may be interesting in future study. We note, however, we have been inspired by the holographic principle when constructing SRDE, SRDE does not necessarily have to be connected with the holographic principle. Although SRDE is phenomenologically successful and theoretically interesting, its physical mechanism is await further study." }, "1101/1101.1451_arXiv.txt": { "abstract": "Optical imaging and spectroscopic observations of the $z = 0.245$ double galaxy cluster Abell 2465 are described. This object appears to be undergoing a major merger. It is a double X-ray source and is detected in the radio at 1.4 GHz. The purpose of this paper is to investigate signatures of the interaction of the two components. Redshifts were measured to determine velocity dispersions and virial radii of each component. The technique of fuzzy clustering was used to assign membership weights to the galaxies in each clump. Using redshifts of 93 cluster members within 1.4 Mpc of the subcluster centres, the virial mass of the NE component is $M_{V} = 4.1 \\pm 0.8 \\times 10^{14} M_\\odot$ and $M_{V} = 3.8 \\pm 0.8 \\times 10^{14} M_\\odot$ for the SW. These agree within the errors with masses from X-ray scaling relations. The projected velocity difference between the two subclusters is 205 $\\pm$ 149 \\kms. The anisotropy parameter, $\\beta$, is found to be low for both components. Spectra of 37\\% of the spectroscopically observed galaxies show emission lines and are predominantly star forming in the diagnostic diagram. No strong AGN sources were found. The emission line galaxies tend to lie between the two cluster centres with more near the SW clump. The luminosity functions of the two subclusters differ. The NE component is similar to many rich clusters, while the SW component has more faint galaxies. The NE clump's light profile follows a single NFW profile with c = 10 while the SW is better fit with an extended outer region and a compact inner core, consistent with available X-ray data indicating that the SW clump has a cooling core. The observed differences and properties of the two components of Abell 2465 are interpreted to have been caused by a collision 2-4 Gyr ago, after which they have moved apart and are now near their apocentres, although the start of a merger remains a possibility. The number of emission line galaxies gives weight to the idea that galaxy cluster collisions trigger star formation. ", "introduction": "Fundamental questions of current astrophysics involve the roles of dark matter, baryonic matter, and dark energy as driven by gravity in the formation of the large-scale structure and galaxies. Double or multiple galaxy clusters can potentially provide information on the dynamics and structure formation on ($r \\ga$ 1 Mpc) scales and interest in them has grown from both the standpoints of modelling and observation. Although the presence of substructure in galaxy clusters has long been known, compared to single galaxy clusters, the properties of double and multiple clusters have received less attention owing to their added complexity. Interest in the observed substructure of galaxy clusters was pioneered by, \\eg Geller \\& Beers (1982) and studies employing the radial infall model (Beers, Geller \\& Huchra 1982; Beers \\etal 1991) was used for rough dynamical estimates. With the realization of their importance, a growing number of systems have now been more fully studied dynamically, from weak lensing, and in X-rays. A partial list includes: Abell 168 (Hallman \\& Markevitch 2004), Abell 399/401 (Sakelliou \\& Ponman 2004), Abell 520 (Girardi \\etal 2008), Abell 521 (Ferrari \\etal 2003), Cl0024+17 (Jee \\etal 2007), the ``bullet cluster,'' 1E0657-56 (Clowe et al. 2006), Abell 2146 (Russell \\etal 2010), RXJ1347.5-1145 (Brada\\v{c} \\etal 2008a), A399 and A401 (Yuan \\etal 2005), Abell 2163 (Maurogordato \\etal 2008), Abell 85 (Tanaka \\etal 2010), and Abell 901/902 (Heiderman \\etal 2009). Okabe \\& Umetsu (2008) studied seven merging clusters using weak lensing. Radio emission from merging clusters in the form of diffuse non-thermal radio halos or relics that arise from merger shocks in the interactions of the colliding galaxy clusters.has also been described by several groups including Slee, Roy, \\& Murgia (2001), Feretti (2002) who describe several objects, Bagchi \\etal (2006), Abell 3376, Orr\\'{u} \\etal (2007), Abell 2744 and Abell 2219, Bonafede \\etal (2009), Abell 2345, van Weeren \\etal (2009), A2256. Skillman \\etal (2010) summarize the the modeling situation. Modelling galaxy cluster mergers and collisions predict observable signatures (\\eg Roettiger \\etal (1996, 1997; Ricker 1998; Tazikawa 2000; Ricker \\& Sarazin 2001; Ritchie \\& Thomas 2002; Springel \\& Farrar 2007; Mastropietro \\& Burkert 2008; Poole \\etal 2008; Planelles \\& Quilis 2009). These simulations have mostly focused on the behavior of the baryonic and dark matter components and use a range of initial profiles and conditions and impact parameters which include both off-centre and head-on collsions. These calculations predict differing behaviors for the baryonic and dark matter components of the clusters at subsequent phases of the collisions. In a typical merger, the dark matter and the baryonic gas are elongated along the collision axis with a displacement between the baryonic and dark matter components. The gas, in addition, is shocked which results in multiple X-ray peaks and gas splashed perpendicularly to the direction of the merger. This produces non-isothermal temperature distributions and the increased ram pressure from the shocks could induce star formation in the member galaxies as well as 'sloshing' (Markevitch \\& Vikhlinin 2007). Several authors have attempted to extract information from double and multiple galaxy clusters on the nature of gravity and dark matter on galactic cluster ($\\sim$ 1 Mpc) distance scales and up to now this has mainly centred on analyzing the 1E0657-56 cluster. Farrar \\& Rosen (2006), Brownstein \\& Moffat (2007), Angus \\& McGaugh (2008), Schmidt, Vikhlinin, \\& Hu (2009), and De Lorenci, Faundez-Abans, \\& Pereira (2009) are among those who investigated whether modifications to gravity are needed to fit the available dynamical data. Springel \\& Farrar (2007), Pointecouteau \\& Silk (2005), and Hayashi \\&White (2006), however indicated that modifications are unneccesary. For studying the properties of dark matter, the situation is somewhat more definite. Clowe \\etal (2006) used weak lensing measurements of the bullet cluster to indicate direct proof of the the presence of dark matter from the offset between the X-ray gas and the lensing centres. Shan \\etal (2010) have studied further offsets between dark and ordinary matter in a further 38 lensed galaxy clusters. Galaxy clusters have been employed to place limits on neutrino masses (\\eg Tremaine \\& Gunn 1979; Angus, Famaey, \\& Diaferio 2010; Natarajan \\& Zhao 2008) and discussed whether or not such particle masses are needed to save the modified Newtonian dynamics formula. Even if one dismisses such claims, a considerable amount of more conventional information is obtainable from double galaxy clusters. This includes possible modifications to luminosity functions, mass profiles and velocity dispersion anisotropy measures (the $\\beta$ parameter) as a result of their interactions. Luminosity functions contain information on the galaxy formation history (\\eg Bingelli, Sandage \\& Tammann 1988) and have been studied in detail at a range of redshifts and environments, mostly for single systems (\\eg Wilson \\etal 1997; De Propis \\etal 2004; Blanton \\etal 2003; Christlein \\& Zabludoff 2003; Goto \\etal 2005). Generally single, and double Schechter functions (Schechter 1976), and Gaussian functionshave been used to fit the LFs. Collisions may modify these properties compared to isolated single clusters at some level, but this question about the effects of merging in double galaxy clusters, whether or not their interactions produce or lower star formation along with AGN activity, has not been answered yet. Hwang \\& Lee (2009) have reviewed empirical and theoretical evidence for this and conclude that observations support the importance of mergers. Haines et al. (2009) and Chung (2010) have reported evidence of enhanced star formation rates in interacting clusters including the bullet cluster. Mergers can distort galaxy cluster mass profiles. Many investigators have compared theoretical mass profiles with observations (\\eg Biviano \\& Girardi 2003; Katgert, Biviano, \\& Mazure 2004; Pointecouteau, Arnaud, \\& Pratt 2005; Kubo \\etal 2007; Okabe \\& Umetsu 2008). Although not all details of these models are agreed upon, the NFW profile (Navarro, Frenk, \\& White 1997) fits most observed profiles within the virial radius with a concentration parameter, $c$ for galaxy clusters is in the range of $c = 4-6$ in agreement with theoretical results (\\eg Zhao \\etal 2003). In addition, for a spherical system with the NFW profile, the anisotropy parameter $\\beta = 1-\\sigma_{\\theta}/\\sigma_{r}$ (where $\\sigma_{\\theta}$ is the azimuthal velocity dispersion and $\\sigma_{r}$ the radial velocity dispersion), is predicted to be near 0 at the centre and to increase to about 0.3 beyond the virial radius and can provide information on the properties of the dark matter (Host 2009). Many of the systems described in the literature are multiple and complex or minor mergers where the mass of one component is considerably larger than the other.The Abell 2465 double cluster discussed in this paper has a relatively uncomplicated substructure and shows some evidence for either a past collision or a commencing merger between the two components. The mean redshift is $z = 0.245$. ROSAT (Perlman et al. 2002), XMM (2008) data, and redshifts show two physically related X-ray sources 5.5$\\arcmin$ (1.2 Mpc) apart (hereafter the NE and SW clumps). As well, it is a 1.4 GHz radio source in the NVSS (Condon \\etal 1998). Both the virial and X-ray masses obtained in this paper indicate that the mass ratio of the two clumps is close to 1:1. Therefore the Abell 2465 is an example of a relatively rare major merger. This paper surveys the optical properties of Abell 2465 cluster and is organized as follows: Section 2 describes the new imaging and spectroscopy, Section 3 gives estimates from spectroscopy of virial masses and radii, values of $\\beta$, and discusses available X-ray and radio data and emission line galaxies in the two subclusters found from the spectra. Section 4 describes the determination of the luminosity functions, and Section 5 compares the estimates of the light profile and the corresponding mass profile. Section 6 discusses the results in relation to a collision, and Section 7 lists the conclusions. The WMAP 5 year cosmological parameters are used throughout this paper. \\begin{figure*} \\begin{center} \\leavevmode \\includegraphics[width=0.83\\textwidth]{cluster.eps} \\caption{A portion of the combined CFHT $i'$ Megaprime images of the centre of Abell 2465 showing the NE (upper left) and SW (lower right) clumps. The vertical edges of the picture are $8 \\farcm 1$ in length. North is to the top and East is to the left. (Plot made using ESO's SkyCat, http://archive.eso.org/cms/tools-documentation/skycat).} \\label{cluster.eps} \\end{center} \\end{figure*} ", "conclusions": "Spectroscopic and photometric observations of the double galaxy cluster Abell 2465 are presented. There are five main conclusions that can be drawn. 1. Concerning the cluster dynamics, the virial masses of the two subclusters are found from fuzzy clustering, which is used to estimate the probability of a galaxy's membership in each clump, with the result that $M_v = 4.0 \\pm 0.8 \\times 10^{14} {\\rm M}_\\odot$ for the NE member and $M_v = 3.8 \\pm 0.8 \\times 10^{14} {\\rm M}_\\odot$ for the SW member and the virial radii are $r_{200} = 1.21 \\pm 0.11$ Mpc and $1.24 \\pm 0.09$ Mpc for NE and SW respectively. The masses compare well with those from X-ray scaling relations that also give temperatures of 4.1 $\\pm 0.3$ and 3.75 $\\pm 0.2$ keV respectively. The velocity difference between the two subclusters is found to be $\\Delta{V}$ = 205 $\\pm$ 149 \\kms which confirms that they are related. Measurement of the clusters' velocity dispersions with radius assuming spherical symmetry indicate that the anisotropy parameter, $\\beta$, is low. 2. There is an excess of star forming galaxies showing emission lines. Of cluster members observed spectroscopically in Figure~\\ref{diagnostic.eps}, 37\\% have detectable H$\\alpha$ emission. These have the properties of star forming galaxies. There are more emission line objects in the SW clump than in the NE clump and there appears to be more emission line galaxies than non-emission between the two clumps. This does not seem to be explained by a selection bias. There is no evidence for strong AGN activity in Abell 2465. This number of emission line objects between the clump centres is unusual when compared to single galaxy clusters. 3. The $r'$ and $i'$ magnitudes show well defined red sequences in each subcluster. The luminosity functions determined within the central 0.6 Mpc of each clump indicate a normal mixture of galactic types. However, the SW region has more galaxies fainter than $M_I = -20.0$ than its NE companion. This could result from their collision or otherwise would suggest different formation histories. The possibility of a background cluster needs to be further checked. 4. The light profiles of both components measured as growth curves were fitted using NFW profiles. The NE clump is fit with a somewhat high concentration parameter $c = 10$, although this depends on the adopted virial radius. The SW clump is fit rather badly with $c = \\sim 4$ and needs a profile with a more compact core. A better fit is a sharp core ($c$ = 120) surrounded by an extended outer region ($c$ = 1.0). This is consistent with Figure~\\ref{XMMNVSS} and published ROSAT data showing that the X-ray core radii differ with $r_c$ of NE being about three times larger than that of the SW and indicates that SW has a cooling core. The derived $I$-band mass to light ratios are $\\Upsilon_I = 84 \\pm 12$ and $112 \\pm 20$ which puts them in the normal range for galaxy clusters. 5. A consistent picture of the collision of the Abell 2465 components is discussed. It is possible that the pair collided 2-4 Gyr ago and are now near maximum separation. The small displacements of the dark matter and baryonic matter as judged by the X-ray data and distribution of the galaxies is consistent with their re-merging after the collision, The high percentage of emission line galaxies in the spectroscopic sample may be a consequence of the collision and are the strongest argument for a past interaction, but this might also be the case if the merger is just starting and interaction occurs along the interface between the two clusters. More models that include the dynamics of the galaxies would be helpful. A weak lensing study of the two components of Abell 2465 is under way." }, "1101/1101.1721_arXiv.txt": { "abstract": "{\\Planck\\!\\!'s all-sky surveys at 30--857 GHz provide an unprecedented opportunity to follow the radio spectra of a large sample of extragalactic sources to frequencies 2--20 times higher than allowed by past, large-area, ground-based surveys. We combine the results of the \\Planck Early Release Compact Source Catalog (ERCSC) with quasi-simultaneous ground-based observations as well as archival data at frequencies below or overlapping \\Planck frequency bands, to validate the astrometry and photometry of the ERCSC radio sources and study the spectral features shown in this new frequency window opened by \\Planck\\!\\!. The ERCSC source positions and flux density scales are found to be consistent with the ground-based observations. We present and discuss the spectral energy distributions (SEDs) of a sample of ``extreme'' radio sources, to illustrate the richness of the ERCSC for the study of extragalactic radio sources. Variability is found to play a role in the unusual spectral features of some of these sources.} ", "introduction": "\\label{sec:introduction} This paper is one of a series based on observations of compact sources by the \\Planck\\!\\!\\footnote{\\Planck (\\url{http://www.esa.int/Planck}) is a project of the European Space Agency (ESA) with instruments provided by two scientific consortia funded by ESA member states (in particular the lead countries France and Italy), with contributions from NASA (USA) and telescope reflectors provided by a collaboration between ESA and a scientific consortium led and funded by Denmark.}~satellite that are included in the ERCSC \\citep{Planck2011-1.10}. Among these ``Early Results''~\\Planck papers there are three that address the extragalactic radio source population. \\citet{Planck2011-6.1} examines statistical properties such as number counts and spectral index distributions, but only at frequencies $\\ge$ 30~GHz.~\\citet{Planck2011-6.3a} incorporates \\Planck measurements and supporting ground-based and satellite observations to refine models for the physical properties of a sample of $\\sim$100 bright blazars. Here we address the observed diverse, sometimes peculiar, spectral properties of sources in the ERCSC, which include peaked-spectrum, flat-spectrum, upturn-spectrum and multicomponent-spectrum sources. We combine archival data as well as new, ground-based, radio observations with the \\Planck data to construct SEDs from $\\sim$3 to $\\sim$200~GHz (the exact frequency coverage varies case by case), to validate the astrometry and photometry of radio sources in the ERCSC. The \\Planck data are also valuable for studying, for instance, high frequency peaked-spectrum sources that were previously underrepresented in radio source populations due to the lack of observations in the sub-millimeter regime. On the other hand, the ground-based data complement the \\Planck data, and are crucial in defining the spectral shape of some sources by extending the observed SEDs. We present a sample of sources with near-simultaneous \\Planck and ground-based observations (primarily employing the VLA, Effelsberg, IRAM and Mets\\\"ahovi telescopes in the northern hemisphere, and ATCA in the southern hemisphere), to control for variability. We also investigate a small number of ERCSC sources at 30 to 70~GHz without clear identification in existing radio surveys. \\subsection{The \\Planck mission} \\Planck \\citep{tauber2010a, Planck2011-1.1} is the third generation space mission to measure the anisotropy of the cosmic microwave background (CMB). It observes the sky in nine frequency bands covering 30 to 857~GHz with high sensitivity and angular resolution from 32$'$ to 5$'$. The Low Frequency Instrument (LFI; \\citealt{Mandolesi2010, Bersanelli2010, Planck2011-1.4}) covers the 30, 44 and 70\\,GHz bands with amplifiers cooled to 20~\\hbox{K}. The High Frequency Instrument (HFI; \\citealt{Lamarre2010, Planck2011-1.5}) covers the 100, 143, 217, 353, 545 and 857\\,GHz bands with bolometers cooled to 0.1~\\hbox{K}. Polarisation is measured in all but the highest two bands \\citep{Leahy2010, Rosset2010}. A combination of radiative cooling and three mechanical coolers produces the temperatures needed for the detectors and optics \\citep{Planck2011-1.3}. Two Data Processing Centers (DPCs) check and calibrate the data and make maps of the sky \\citep{Planck2011-1.7, Planck2011-1.6}. \\Planck\\!\\!'s sensitivity, angular resolution and frequency coverage make it a powerful instrument for Galactic and extragalactic astrophysics as well as cosmology. The scan strategy employed in the \\Planck mission is described in \\citet{Planck2011-1.1}. As the satellite spins, sources are swept over the focal plane, as indicated schematically in Figure~\\ref{fig_focalplane}. In the course of each day, the pointing axis of the telescope is adjusted by $\\sim$1$\\degr$, so a given source will follow a slightly different track across the focal plane; thus its flux density is the average of many such scans. In addition, since the data included in the ERCSC amount to $\\sim$1.6 full-sky surveys, some sources have been covered twice with a time separation of $\\sim$6 months. Finally, sources near the ecliptic poles, where the scan circles intersect, are often covered multiple times. It is thus important to keep in mind that the flux densities cited in this paper (and indeed in the ERCSC as a whole) are {\\it averaged}. Figure~\\ref{fig_focalplane} illustrates that flux measurements at 44~GHz are particularly susceptible to time-dependent effects, because of the wide spacing of the 44\\,GHz horns in the focal plane (see further discussion in \\S\\,\\ref{sec:artifacts}). \\begin{figure} \\centering \\includegraphics[scale=0.20]{16475fig1.eps} \\caption{Focal plane, showing spacing of the \\Planck receivers. Each day, the pointing is adjusted by $\\sim$1$\\degr$ in the vertical direction. Note that the wide separation of the three 44\\,GHz horns (two at the top, one at the bottom) causes the 44\\,GHz observations of a given source to take place at two times separated by 7--10 days for each scan.} \\label{fig_focalplane}% \\end{figure} \\subsection{The ERCSC} The \\Planck Early Release Compact Source Catalog \\citep{Planck2011-1.10} provides lists of positions and flux densities of compact sources at each of the nine \\Planck frequencies. For frequencies from 30 to 143~GHz (those mostly cited in this paper), sources were detected using Powell Snakes techniques \\citep{carvalho2009}. In the four highest frequency channels, sources were located using the SExtractor method \\citep{bertin1996}. A set of selection criteria was further applied to select sources that are included in the ERCSC. The primary criterion was a Monte Carlo assessment designed to ensure that $\\ge\\,$90\\,\\% of the sources in the catalogue are reliable and have a flux density accuracy better than 30\\,\\%. External validation (also discussed in \\citealt{Planck2011-1.10}) shows that ERCSC met its reliability criterion, and we show evidence in \\S\\,\\ref{sec:validation} that the flux density scale of ERCSC is accurate. Secondary selection criteria, including the elimination of extended sources, were also applied. Virtually all of the sources discussed in this paper, and the vast majority of extragalactic ERCSC sources, were unresolved by \\Planck and in many cases even at the much higher angular resolution of the VLA or other ground-based instruments. The flux densities in the ERCSC are calculated using aperture photometry. The effective band centers or corresponding colour corrections depend to some degree on the spectrum of the source being observed. This relatively small dependence is discussed in the LFI and HFI instrument papers \\citep{Planck2011-1.4, Planck2011-1.5, Planck2011-1.6, Planck2011-1.7}. We adopt here (and list in Table~\\ref{table_bands}) the effective central frequencies defined in those papers, and the colour corrections defined for a source with a spectral index $\\alpha = -0.5$ (using the convention $S_{\\nu} \\propto \\nu^{\\alpha}$). To obtain the correct flux density for an assumed narrow band measurement, we divide the tabulated ERCSC flux densities by the colour correction factor at the corresponding central frequency. After the release of the ERCSC, it was found that the correction for aberration introduced by the motion of the satellite had not been correctly made. This introduces small ($<$\\,0.35$'$) errors in the catalogued positions of sources; these can in turn produce very small errors ($<$\\,1\\,\\%) in flux densities. We have not corrected the ERCSC flux densities for this small effect. In the next section, we describe briefly some general properties of the extragalactic radio sources in the ERCSC. We describe in \\S\\,\\ref{sec:observation} related ground-based observations. In \\S\\,\\ref{sec:validation}, we discuss the identification of sources, positional accuracy and flux density scales, comparing those obtained from \\Planck with ground-based measurements. This parallels the validation work discussed in \\cite{Planck2011-1.10}. Variability of source luminosity and the issue of different resolutions are also discussed. We present and discuss in \\S\\,\\ref{sec:ers} several examples of interesting classes of extragalactic radio sources. We conclude in the final section and point towards further research on these radio sources and many others contained in the ERCSC. \\begin{table} \\caption{Parameters of \\Planck bands employed in this paper.} \\label{table_bands} \\centering \\begin{tabular}{c c c c } \\hline\\hline Planck & Central Frequency & Colour & Beam FWHM \\\\ Band & [GHz] & Correction & [arcmin] \\\\ \\hline\\\\ 30 & 28.5 & 1.037 & 32.6 \\\\ 44 & 44.1 & 1.018 & 27.0 \\\\ 70 & 70.3 & 1.031 & 13.0 \\\\ 100 & 100 & 0.999 & 9.94 \\\\ 143 & 143 & 1.006 & 7.04 \\\\ 217 & 217 & 0.993 & 4.66 \\\\ 353 & 353 & 0.990 & 4.41 \\\\ \\hline \\end{tabular} \\end{table} ", "conclusions": "We summarise in this section the primary conclusions to be drawn from \\Planck\\!\\!'s study of extragalactic radio sources. We stick closely to the observational results, and provide comments on the fit between these observational results and current theories of the physics of extragalactic radio sources. \\subsection{Overall properties} \\Planck has demonstrated that the high frequency counts (at least for frequencies $\\leq$\\,143 GHz) of extragalactic sources are dominated at the bright end by synchrotron emitters, not dusty galaxies. This finding is in agreement with conclusions reached by the South Pole Telescope (SPT) and Atacama Cosmology Telescope (ACT) teams (\\citealt{vieira2010, marriage2010}, respectively) from the properties of sources at much lower flux densities. An inference from this result is that the cores of extragalactic sources, and not extended structure, dominate the high frequency emission, as suggested by earlier work at frequencies below most of the \\Planck bands (see recent discussions in \\citealt{lin2009} and \\citealt{murphy2010}). The conclusion that the core dominates the high frequency emission is supported by the close agreement between \\Planck and VLA flux densities (\\S\\,\\ref{sec:validation}), implying that the emission is found within the VLA beam of order arcseconds in size. The emerging dominance of emission from a flat spectrum core implies that extrapolation of flux densities and/or counts of sources to frequencies $\\geq$\\,30\\,GHz cannot be reliably made from low frequency catalogues, where the emission is dominated by lobes. This has been recognised for some time (see \\citet{lin2009} and \\citet{murphy2010}); \\Planck strongly confirms this conclusion. If we look at the SEDs of the sources, including all components, the spectra grow flatter as frequency increases and the steep-spectrum lobes fade away. On the other hand, \\Planck observations allow us to follow the SEDs of a statistically significant sample of sources to much higher frequencies than can generally be employed in ground-based observations. As noted in \\S\\,\\ref{sec:validation} (and described in more detail in \\citealt{Planck2011-6.1}), \\Planck provides clear evidence of a spectral steepening in the radio and millimeter wave emission from extragalactic cores; the steepening sets in at frequencies above $\\sim$44--70\\,GHz. The observed spectral steepening in turn means that radio sources contribute less foreground noise to increasingly sensitive searches for small angular scale anisotropies in the CMB than earlier models had suggested (e.g., \\citealt{dezotti2005}). This finding, too, is consonant with results from SPT and ACT. \\subsection{Properties of individual sources} \\Planck allows us for the first time to investigate sub-mm spectra beyond 200~GHz, which is the limit of most ground-based monitoring programs. In this regime, the spectra of radio sources are usually expected to be optically thin synchrotron emission, and our results largely confirm that this part of the spectrum is represented by a single power law. The vast majority of extragalactic sources in the ERCSC lists at 30--100 GHz are flat spectrum radio sources of the sort discussed in detail in \\citet{Planck2011-6.3a}, with a scattering of bright steep spectrum sources strong enough to register in the lowest frequency bands of \\Planck\\!\\!. A small fraction ($\\lsim$\\,10\\,\\%; Table~\\ref{table_candi_gps}) show peaked or convex spectra, but even these sources are mostly identified as known blazars. As blazars are known to be variable, care has to be taken in the interpretation of their spectra. For most sources, ERCSC flux densities represent the average over two scans, and even during one scan, spectral artifacts may occur due to the fact that it requires 7--10 days for the \\Planck focal plane to cross a point source. In contrast, most known examples of compact, newborn radio galaxies (CSOs), which were originally thought to be the dominant class of sources with Gigahertz-peaked spectra, are mostly too faint to be detected by {\\it Planck}. A new population of bright, very compact, high-peaked CSOs has not been found. Likewise, very few of the extragalactic radio sources found by \\Planck show evidence of the sharp spectral upturn expected from dust reemission at high frequencies. Only NGC253 shows a clear upturn at a frequency $\\leq$\\,143 GHz; sources that show dust reemission dominating at 217\\,GHz are in many cases Galactic sources, in the Magellanic clouds, or nearby known star-forming galaxies. Although we have not investigated in detail all the sources with no obvious matches in lower frequency radio catalogues (\\S\\,\\ref{sec:unmatched}), let alone every extragalactic source in the ERCSC, we find no convincing evidence for the emergence of a new and unexpected population of sources. The 26 ERCSC sources with no match in radio catalogues appear to be a heterogeneous mixture of conventional radio sources, many of them Galactic. \\subsection{Future observations and analysis} Two sets of future observations will help clarify the status of some of the sources listed in \\S\\,\\ref{sec:peaked} and \\S\\,\\ref{M-C-spec}. Careful monitoring and/or VLBI observations of sources with peaked spectra would determine whether they are highly compact radio galaxies (CSOs), compact jet objects with stable GPS type spectra, or are instead merely flaring sources with temporarily convex spectra. In fact, our results seem to suggest most extreme spectral features seen in \\Planck sources may be associated with flares. With at least four full sky surveys, \\Planck is a powerful tool to distinguish these possibilities, by deriving catalogues from individual, six-month, sky surveys and comparing flux densities of all bright point sources at six-month intervals. Further ground-based observations, now underway at frequencies below and overlapping the \\Planck frequency bands, will support this effort. Likewise, for the non-variable extreme radio sources like CSOs, the addition of more surveys to the \\Planck maps will allow us to extract deeper catalogues, with the potential to find some more of these usually faint objects. In the case of Galactic sources, the role of CO emission lines is important. The situation for extragalactic sources is more complicated, because the CO lines redshift in and out of the \\Planck bands. We have underway a study to look for the influence of CO emission in the ERCSC spectra. We expect the effect to be small, since most of the ERCSC sources are extremely bright, non-thermal emitters, with strong continua. There is no obvious evidence in the SEDs we have examined for the presence of CO emission. In this regard, we again warn readers to be careful in interpreting anomalous 44\\,GHz observations; a spectral bump at 44 GHz is not necessarily evidence for redshifted CO. We used ground-based observation data for both validation purposes, specifically cross-calibration between ground-based instruments and \\Planck\\!\\!, and in the study of the spectral and variability properties of extreme radio sources. Most of these observation programs, like the \\Planck mission itself, are still ongoing." }, "1101/1101.4935_arXiv.txt": { "abstract": "One of the main achievements in modern cosmology is the so-called `unified model', which successfully describes most classes of active galactic nuclei (AGN) within a single physical scheme. However, there is a particular class of radio-luminous AGN that presently cannot be explained within this framework -- the `low-excitation' radio AGN (LERAGN). Recently, a scenario has been put forward which predicts that LERAGN, and their regular `high-excitation' radio AGN (HERAGN) counterparts represent different (red sequence vs.\\ green valley) phases of galaxy evolution. These different evolutionary states are also expected to be reflected in their host galaxy properties, in particular their cold gas content. To test this, here we present CO(1$\\rightarrow$0) observations toward a sample of 11 of these systems conducted with CARMA. Combining our observations with literature data, we derive molecular gas masses (or upper limits) for a complete, representative, sample of 21 $z<0.1$ radio AGN. Our results yield that HERAGN on average have a factor of $\\sim7$ higher gas masses than LERAGN. We also infer younger stellar ages, lower stellar, halo, and central supermassive black masses, as well as higher black hole accretion efficiencies in HERAGN relative to LERAGN. These findings support the idea that high- and low-excitation radio AGN form two physically distinct populations of galaxies that reflect different stages of massive galaxy build-up. ", "introduction": "\\label{sec:intro} Over the past two decades a standard model of AGN has emerged. In this `unified' model efficient disk accretion of cold matter on the central supermassive black hole (BH) provides the radiation field that photoionizes emission-line regions. However, there is a certain fraction of AGN identified by radio observations that poses a challenge to the unified model, the so-called low-excitation radio AGN (hereafter:\\ LERAGN). The main difference between {\\em high-excitation radio AGN (HERAGN)} and these {\\em LERAGN} is that the latter do not exhibit strong emission lines in their optical spectra (Jackson \\& Rawlings 1997; Evans et al.\\ 2006). Recently, Hardcastle et al.\\ (2006) have suggested that high- and low-excitation radio AGN may represent a principal separator between populations fundamentally different in their black hole accretion mechanisms (see also Evans et al. 2006; Allen et al. 2006; Kewley et al. 2006). They developed a model in which central supermassive black holes of HERAGN accrete in a standard (radiatively efficient) way from the cold phase of the intragalactic medium (IGM), while those of LERAGN are powered in a radiatively inefficient manner by Bondi accretion of the hot IGM. \\smo\\ (2009) showed that low- and high excitation radio AGN exhibit not only systemic differences in their black hole masses and accretion rate properties, but also in their host galaxy properties, such as stellar masses and stellar populations. This is consistent with these two classes of radio AGN dividing in a stellar mass vs.\\ color plane in such a way that LERAGN occupy the red sequence and HERAGN inhabit the so called ``green valley'', a sparsely populated region between the blue-cloud and the red-sequence (\\smo\\ 2009). The stellar mass vs.\\ color plane can be interpreted as a time-sequence for galaxy evolution. Galaxies are thought to evolve from an initial star-formation-dominated state with blue optical colors into the most massive \u201cred-and-dead\u201d galaxies through a transition phase reflected in the green valley (Bell et al. 2004a, 2004b; Borch et al. 2006; Faber et al. 2007; Brown et al. 2007). In recent years it has been suggested that radio outflows from AGN likely play a crucial role in this massive galaxy build-up \\citep{croton06,bower06,sijacki06,sijacki07,fanidakis10}. In this context the radio-AGN feedback (often called the ``radio'' or ``maintenance'' mode), which is thought to limit stellar mass growth in already massive galaxies, is expected to occur only in LERAGN (\\smo\\ 2009). Furthermore, it has been shown that the cosmic evolution of the space density of various types of radio AGN is significantly different (e.g.\\ Peacock et al.\\ 1985, Willott et al.\\ 2001; Smolcic et al.\\ 2009). Based on a study of the evolution of the radio AGN luminosity function out to $z=1.3$, \\smo\\ et al.\\ (2009) have shown that the comoving space density of low-luminosity radio AGN (predominantly LERAGN) only modestly declines since $z=1.3$, while that of powerful AGN (predominantly HERAGN) dramatically diminishes over the same cosmic time interval. This suggests that LERAGN and HERAGN not only represent physically distinct galaxy populations, but also populations in different stages of massive galaxy build-up. If this is the case, the molecular gas masses and fractions in low- and high- excitation radio AGN are expected to directly reflect this trend. We here investigate this idea by observing CO($J$=1$\\to$0) emission of a carefully selected, representative sample of nearby ($z<0.1$) HE- and LERAGN with CARMA. We adopt a $\\Lambda$CDM cosmology with $H_0=70$, $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$. ", "conclusions": "\\label{sec:discussion} Our main result is that HERAGN have systematically higher molecular gas masses (a factor of $\\sim7$; see \\t{tab:averageprops} ), compared to LERAGN. Flaquer et al.\\ (2010) have found a similar trend by dividing their sample ($\\sim50$ radio AGN observed with the IRAM 30m telescope, partially overlapping with our sample) into FR class I and II objects. They find that the molecular gas mass in FR~IIs is a factor of $\\sim4$ higher than that in FR~Is. The FR class can be taken to roughly correspond to the low- and high- excitation classification.\\footnote{Almost all FR I \u2013- low power -\u2013 radio galaxies are LERAGN, while optical hosts of FR IIs, which are typically more powerful than FR Is (Fanaroff \\& Riley 1974; Owen 1993; Ledlow \\& Owen 1996), usually have strong emission lines. Note however that the correspondence between the FR class and the presence of emission lines is not one-to-one.} Flaquer et al.\\ (2010) have, however, concluded that the systematic differences they find are likely a result of a Malmquist bias, i.e.\\ simply due to a systematically higher redshift of their FR-II sources. Although our HERAGN lie on average at a slightly higher redshift, compared to our LERAGN (0.046 vs.\\ 0.030, resp.) in the following we argue that the systematic differences we find in molecular gas mass are not due to a Malmquist bias. Mori\\'{c} et al.\\ (2010) have shown that the redshift distributions of carefully selected samples of radio-selected LINERs and Seyferts are approximately the same (see their Fig.~6). This eliminates Malmquist bias from their results. They find that the detection fraction in the FIR is significantly lower for LINERs than for Seyferts (6.5\\% vs.\\ 22\\%, resp.) in their sample. Assuming that the star formation law parameterized by $L'_{\\rm CO}$ (as a proxy for total gas mass) and $L_{\\rm FIR}$ (as a proxy for star formation rate; e.g., Kennicutt 1998; Solomon \\& Vanden Bout 2005; Bigiel et al.\\ 2008), on average, correctly represents the star formation properties of these samples (as confirmed by the CO/FIR luminosities of the IRAS detected sources analyzed here; see \\f{fig:cofir} ), the lower average FIR luminosity in low-excitation sources (i.e.\\ LINER) implies lower gas masses than in high-excitation (i.e.\\ Seyfert) types of galaxies. A similar result is obtained based on average (optically derived) star formation rates\\footnote{Mori\\'{c} et al.\\ (2010) derived star formation rates for each galaxy in their sample via stellar population synthesis model fitting to the SDSS photometry of the host galaxy (see also \\smo\\ et al.\\ 2008).}, suggesting that those in LINERs are by about a factor of 3 lower than in Seyferts in a redshift-matched sample. These findings suggest that the systematic differences in molecular gas mass in high- and low-excitation radio AGN are physical, and not due to Malmquist bias. \\begin{figure} \\includegraphics[bb=94 400 410 712,scale=1.5,width=\\columnwidth]{f4.eps} \\caption{ CO vs.\\ FIR luminosity for our local AGN sources detected with IRAS. The lines represent the $L'_\\mathrm{CO} - L_\\mathrm{FIR}$ correlation derived by \\citet{riechers06}. } \\label{fig:cofir} \\end{figure} The systematically higher molecular gas masses that we find in HERAGN, relative to LERAGN in our $z<0.1$ radio AGN sample, are in excellent agreement with the systematic differences in various properties of high- and low- excitation radio AGN, both on pc- and kpc- galaxy scales (see \\s{sec:intro} \\ and \\t{tab:physprops} ). We find that, on average, HERAGN have lower stellar masses and stellar ages compared to LERAGN (see \\t{tab:averageprops} ; see also \\smo\\ 2009). This is consistent with HERAGN and LERAGN being green valley and red sequence sources, respectively. Furthermore, we show that HERAGN have on average higher radio luminosities than LERAGN, consistent with the results presented in Kauffmann et al.\\ (2008). Kauffmann et al.\\ have shown that the fraction of radio AGN with strong emission lines in their spectra significantly rises beyond $\\sim10^{25}$~\\wh . In general, the comparison of the black hole and host galaxy properties inferred for our 21 $z<0.1$ AGN with much larger samples of radio AGN (Kauffmann et al.\\ 2008; \\smo\\ 2009) suggests that our AGN sample is representative of high- and low-excitation radio AGN in the nearby universe. From the average stellar masses that we infer for our high- and low excitation sources we extrapolate that they occupy $\\sim3\\times10^{13}$~\\msol\\ and $\\sim5\\times10^{14}$~\\msol\\ halos, respectively \\citep[e.g.][]{behroozi10, moster10}. Compared to the systematic molecular gas mass difference, this yields an even more dramatic discrepancy of more than 2 orders of magnitude in the average molecular gas fractions in HE- ($\\sim10^{-5}$) and LE-RAGN ($\\sim9\\times10^{-8}$). The discrepancy remains significant (about an order of magnitude) if the average gas-to-stellar mass fraction (which can be interpreted as star formation efficiency) is considered. On small scales, the average black hole accretion efficiencies in HE- and LE-RAGN suggest different supermassive black-hole accretion mechanisms (standard disk accretion of cold gas in HERAGN vs.\\ Bondi accretion of hot gas in LERAGN; see Evans et al.\\ 2006; Hardcastle et al.\\ 2006). Furthermore, the higher black hole masses in LERAGN suggest a later evolution stage of their host galaxies, compared to that of HERAGN. This is further strengthened by the higher stellar masses in LERAGN, as well as older stellar ages, and less massive gas reservoirs. In the blue-to-red galaxy formation picture, blue gas rich galaxies are thought to transform into read-and-dead gas-poor galaxies, the stellar populations in the host galaxies of HERAGN are expected to be younger and have lower masses, while their molecular gas reservoirs -- fueling further stellar mass growth -- are expected to be higher than those in LERAGN. This is in very good agreement with the results presented here. Thus, in summary, our results strengthen the idea that low- and high-excitation radio AGN form two physically distinct galaxy populations that reflect different stages of massive galaxy formation." }, "1101/1101.5178_arXiv.txt": { "abstract": "{} {We investigate the new and still poorly studied matter of so-called multiple stellar populations (MSPs) in Galactic globular clusters (GGCs). Studying MSPs and their accumulated data can shed more light on the formation and evolution of GGCs and other closely related fundamental problems. We focus on the strong relation between the radial distribution of evolutionary homogeneous stars and their $U$-based photometric characteristics in the nearby GGC NGC 6752 and compare this with a similar relation we found in NGC 3201 and NGC 1261.} {We use our new multi-color photometry in a fairly wide field of NGC 6752, with particular emphasis on the $U$ band and our recent and already published photometry made in NGC 3201 and NGC 1261.} {We found and report here for the first time a strong difference in the radial distribution between the sub-populations of red giant branch (RGB) stars that are bluer and redder in color $(U-B)$, as well as between sub-giant branch (SGB) stars brighter and fainter in the $U$-magnitude in NGC 6752. Moreover, the fainter SGB and redder RGB stars are similarly much more centrally concentrated than their respective brighter and bluer counterparts. Virtually the same applies to NGC 3201. We find evidence in NGC 6752 as in NGC 3201 that a dramatic change in the proportion of the two sub-populations of SGB and RGB stars occurs at a radial distance close to the half-mass radius, $R_h$, of the cluster. These results are the first detections of the radial trend of the particular photometric properties of stellar populations in GGCs. They imply a radial dependence of the main characteristics of the stellar populations in these GGCs, primarily of the abundance, and (indirectly) presumably of the kinematics.} {} \\keywords {globular clusters: general -- globular clusters: individual: MGC 6752, NGC 3201, NGC 1261} ", "introduction": "\\label{introduc} The southern Galactic globular cluster (GGC) NGC 6752 is one of the nearest and most frequently studied GGCs. However, we here deal with a fairly new aspect, which is related to photometric manifestations of the so-called multiple stellar populations (MSPs) in GGCs. Beside $\\omega$ Cen they were initially revealed thanks to accurate Hubble space telescope photometry in most massive GGCs, such as NGC 2808 (Piotto et al. \\cite{piottoetal07}) NGC 1851 (Milone et al. \\cite{milonetal08}) among others, caused by the splitting of the sub-giant branch (SGB) or/and of the main sequence (MS). The splitting of the red giant branch (RGB) in most massive GGCs was also shown by Lee et al. (\\cite{leetal09}) using ground-based observations. Soon after that, evidence of MSPs have been revealed in lower mass GGCs, similar to NGC 6752. A recent paper by Milone et al. (\\cite{milonetal10}) is the first publication devoted to MSPs in NGC 6752. They argue that the broadening of the MS, visible in the color-magnitude diagram (CMD) obtained from the high-precision Hubble space telescope photometry of the cluster cannot be attributed to photometric errors or to binary stars. Also, a spread in the $(U-B)$ color of RGB stars is found to correlate with variations in Na and O abundances. We find and report on direct evidence for the inhomogeneity of the clusters' stellar population caused by the obvious manifestation, found for the first time in NGC 6752, of strong radial segregation between photometrically differing sub-populations of evolutionary homogeneous cluster stars. We also compare the obtained results with our recent similar findings in NGC 3201 (Kravtsov et al. \\cite{kravtsovetal10a}) and NGC 1261 (Kravtsov et al. \\cite{kravtsovetal10b}). ", "conclusions": "Based on a new multi-color photometry in a fairly wide field of NGC 6752, we found and report the following direct evidence of the inhomogeneity (multiplicity) of the cluster's stellar population. There is an essential radial segregation of SGB stars in the cluster, depending on their brightness in the $U$ band: the fainter sub-population of sub-giants is obviously centrally concentrated, and the bulk of these are confined within the cluster region with the half-mass radius, $R_h$, while their brighter counterparts show a flatter radial distribution within the observed field. The sub-populations of photometrically distinct RGB stars also exhibit essential radial segregation in NGC 6752: RGB stars redder in the color $(U-B)$, like the sub-population of the fainter SGB stars, are centrally concentrated mostly within the region with radius $R \\sim R_h$, as opposed to the sub-populations of both RGB stars bluer in the $(U-B)$ color and SGB stars brighter in the $U$ band. We note virtually the same radial segregation between photometrically distinct sub-populations both of RGB and SGB stars in NGC 3201. In NGC 1261 (Kravtsov et al. \\cite{kravtsovetal10b}), we also found similar trends, but we are not able to adequately judge what exactly happens in its central part. The revealed and demonstrated strong radial segregation between the sub-populations in NGC 6752 and NGC 3201 is closely related with the difference of their photometric characteristic either in the $U$ magnitude (SGB stars) or in the $(U-B)$ color (RGB stars), which are known to be sensitive to metallicity. Therefore, the obtained results not only provide evidence for the radial segregation itself between the sub-populations, but also imply (1) the radial dependence of the abundance of those elements, which are mainly responsible for this photometric difference between stars in the clusters, and (2) presumably different kinematics between the sub-populations. In this context, it would be important to study whether there is a radial trend (and what it is exactly) of the elemental abundance in NGC 6752 and other GGCs." }, "1101/1101.0870_arXiv.txt": { "abstract": "Observational evidence for black hole spin down has been found in the normalized light curves of long GRBs in the BATSE catalogue. Over the duration $T_{90}$ of the burst, matter swept up by the central black hole is susceptible to non-axisymmetries producing gravitational radiation with a negative chirp. A time sliced matched filtering method is introduced to capture phase-coherence on intermediate timescales, $\\tau$, here tested by injection of templates into experimental strain noise, $h_n(t)$. For TAMA 300, $h_n(f)\\simeq 10^{-21}$ Hz$^{-\\frac{1}{2}}$ at $f=1$ kHz gives a sensitivity distance for a reasonably accurate extraction of the trajectory in the time frequency domain of about $D\\simeq 0.07-0.10$ Mpc for spin fown of black holes of mass $M=10-12M_\\odot$ with $\\tau=1$ s. Extrapolation to advanced detectors implies $D\\simeq 35-50$ Mpc for $h_n(f)\\simeq 2\\times 10^{-24}$ Hz$^{-\\frac{1}{2}}$ around 1 kHz, which will open a new window to rigorous calorimetry on Kerr black holes. ", "introduction": "Calorimetry and spectroscopy on all radiation channels will be key to identifying Kerr black holes in the Universe \\citep{van02}. They may be central to some of the core-collapse supernovae (CC-SNe; \\cite{woo93}) and mergers of neutron stars with another neutron star (NS-NS) or with a companion black hole hole (BH-NS; \\cite{pac91})--the currently favored astronomical progenitors of cosmological gamma-ray bursts (GRBs). While all short GRB are probably produced by mergers, the converse need not hold in general. Black holes with a neutron star companion should have a diversity in spin \\citep{van99}. These mixed binaries can be short and long lived, depending on the angular velocity of the black hole. It predicts that some of the short GRBs, produced by slowly spinning black holes, also feature X-ray afterglows \\citep{van01a} as in GRB050509B \\citep{geh05} and GRB050709 \\citep{vil05,fox05,hjo05}. Long GRBs, produced by rapidly rotating black holes, are expected to form in CC-SNe from short, intra-day period binaries (\\cite{pac98,van04}), but also from mergers with a companion neutron star \\citep{van99} or out of the merger of two neutron stars \\citep{bai08,van09b}. A diversity in the origin of long GRBs in CC-SNe and mergers naturally accounts for events with and without supernovae, notably GRB060614 \\citep{van08a,caito09}, with and without pronounced X-ray afterglows \\citep{van09} and in wind versus constant density host environments, that are relevant to recent studies of extraordinary {\\em Swift} and {\\em Fermi}-LAT events \\citep{cen10,cen10b}. Rapidly rotating black holes can sweep up surrounding matter and induce the formation of multipole mass-moments. In this process, spin energy is catalytically converted to a long duration gravitational wave burst (GWB; \\cite{van01b,van03a}). For stellar mass black holes, this output may be detected by advanced gravitational wave detectors for events in the local Universe. As candidate inner engines to long GRBs, evidence for the associated spin down of the black hole has been found in the normalized light curve of 600 long GRBs in the BATSE catalogue \\citep{van09}. The high frequency range of the planned advanced detectors LIGO-Virgo \\citep{bar99,arc04}, the Large-scale Cryogenic Gravitational-wave Telescope (LCGT, \\cite{lcgt}) and the Einstein Telescope (ET, \\cite{et08}) covers the quadrupole emission spectrum of orbital motions around stellar mass black holes, thus establishing a window to rigorously probe the inner most workings of GRBs and some of the CC-SNe. We anticipate an event rate for long GWBs of 0.4-2 per year within a distance of 100 Mpc from the local event rate of long GRBs \\citep{gue07}. It compares favorably with that of mergers of binary neutron stars (e.g \\cite{osh08,aba10}). Since the local event rate of type Ib/c supernovae is $\\sim 80$ per year within a distance of 100 Mpc, the branching ratio of Type Ib/c supernovae into long GRBs is therefore rather small, about 0.5 \\citep{van04} up to $\\sim 2.5\\%$ \\citep{gue07}. It suggests the existence of many failed GRB-supernovae, notably supernovae with relativistic ejecta, supernovae with pronounced aspherical explosions and radio- loud supernovae (e.g. \\cite{del10} and references therein). Therefore, the rate of events of interest to potential bursts in gravitational waves appears to be 1-2 orders of magnitude larger than the event rate of successful GRB-supernovae. In addition, also type II SNe, whose event rate is 3-4 times larger than that of type Ib/c \\citep{cap99,man05}, may explode and expand asymmetrically \\citep{hoe99,ish92}, suggesting a significant additional potential for gravitational-waves burst production. A blind rather than a triggered search for bursts events in the local Universe seems to be appropriate by taking advantage of the all-sky monitoring capability of the gravitational-wave detectors in view of a beaming factor of gamma-ray bursts of $f_b<10$ $(\\theta>25$ deg) up to a few hundred ($\\theta\\sim 4$ deg, \\cite{fra01,van03,gue07}). A blind search is also expected to be competitive with current X/optical surveys for detecting the shock break-out associated with an emerging CC-SNe, as it lasts only a few dozens or minutes up to a few hours, and it naturally includes the possibility of long GRBs coming from merger events with no supernova, which may be exemplified by the long event GRB060614 of duraton 102 s discovered by {\\em Swift.} While the energy output in long gravitational wave bursts (GWBs) produced by rapidly rotating black holes should be large, searching for these bursts by matched filtering is challenging in view of anticipated phase-incoherence due to turbulent magnetohydrodynamical motions in the inner disk or torus. Here, we focus on the detection of a trajectory in the time frequency domain produced by long GWBs, satisfying phase-coherence on short up to intermediate timescales. This objective goes further than the detection of a burst signal, with the aim to extract reasonably accurate information on the burst evolution. Inevitably, the sensitivity distance for extracting trajectories is considerably more conservative than the sensitivity distance for a detection per se. We shall discuss a new matched filtering detection algorithm to detect trajectories in the time frequency domain for long GWBs with slowly varying frequencies lasting tens of seconds with intermittent phase coherence. For a burst lasting 50 s, for example, the algorithm searches by matched filtering using segmented templates on a time scale of, e.g., 1 s. This procedure gives a compromise between optimal matched filtering, applicable to phase-coherence extending over the entire burst duration as in binary coalescence of two black holes, and second order methods by correlation of independent detector signals in the time-domain. For our example, the compromise results in a sensitivity distance below that of optimal matched filtering by a factor of about $\\sqrt{50}\\sim 7$, and an improvement by a factor of $\\sqrt{1000}\\simeq$30 over second order methods for signals around 1 kHz. In Section 2, we discuss the astronomical origin of long GRBs from possible both CC-SNe and mergers. In Section 3, we introduce a model and template for long GWBs from rapidly rotating Kerr black holes. In Section 4, we describe the proposed time sliced matched filtering search algorithm and the evaluation of the sensitivity distance for a reasonably accurate extraction of trajectories in the time frequency domain. Our findings are summarized in Section 5. ", "conclusions": "The diversity in multi-wavelength phenomenology on long GRBs strongly suggests a common inner engine that is intrinsically long-lived representing the outcome of various astronomical scenarios. We here identify the inner engine with rapidly rotating Kerr black holes, whose lifetime is set by the secular time scale of spin in a process of spin down against surrounding high density matter. At present, the most quantitative observational evidence for black hole spin down is found by matched filtering analysis of 600 light curves of long GRBs in the BATSE catalogue \\citep{van09}. This mechanism points to major contemporaneous emissions in gravitational waves and MeV-neutron emissions, that is likely to be dominant over the presently observed electromagnetic radiation in GRB-afterglow emissions and kinetic energies in GRB-SNe. True calimetry on the inner engine requires observations in these non-electromagnetic windows. Most GRBs remain unobserved, due to beaming factors of ``a few'' to about 500 \\citep{fra01,van03}. For CC-SNe, there is a considerable uncertainty of up to several days for the time of onset based on extrapolating backwards in time the supernova optical light curve. Not all long GRBs appear to be associated with CC-SNe, i.e., GRB 060614 \\citep{del06,fyn06,geh06,gal06} appeared without a bright supernova. If GRB 060614 was not an anomalously faint (``failed\") supernova, it may have been a merger event \\citep{van08a,caito09}. Similar considerations apply to the halo event GRB070125, which appeared with no optically identified host galaxy \\citep{cen08,cha08}. A significant fraction of binaries exist in globular clusters as indicated by their population of luminous X-ray sources (\\cite{ver06} and references therein). As the number of globular clusters correlates with the luminosity of elliptical galaxies \\citep{har81,bur10}, the latter in particular may be preferred sites for GRB070125 type events. Therefore, blind, untriggered searches appear to be preferred in searches for local events, exploiting the all-sky survey capability of gravitational wave detectors. To explore the sensitivity distance of advanced gravitational-wave detectors to the anticipated long duration negative chirps from events in the local Universe, we introduce a time sliced matched filtering algorithm and apply it to the strain noise amplitude of the TAMA 300 detector with signal injection. A time sliced approach can circumvent the limitations posed by phase-incoherence on the time scale of the duration of tbe burst, which generally inhibits optimal matched filtering using complete wave form templates. By injecting our model template into the strain noise data of the TAMA 300 detector during a run with $h_n(t)\\simeq 1\\times 10^{-21}$ Hz$^{-\\frac{1}{2}}$, we compute correlations between the $i-$th time slice (of duration $\\tau$) and one frame (about 52 s) of detector output. For a typical black hole mass of $10M_\\odot$, the results indicate the importance of the middle and late time behavior of the bursts (Fig. 3) below 2000 Hz, but less so the initial spin down phase associated with higher frequencies when starting from a maximal rotation. This reduction results from the relatively large strain amplitude noise of the detector in the shot-noise region for frequencies of a few kHz. Effectively, the search is focused on the output post-maximum $(a/M<0.8)$, which obviates the need to consider templates over a wide range of initial spin. A complete search will include a scan over two parameters: a range of intermediate time scales, e.g., $\\tau=0.1-1$, and scaling of frequencies of the templates to account for a diversity in black hole masses, e.g., between 5-15 $M_\\odot$. The range $\\tau=0.1-1$ s represents tens to hundreds of orbital periods of the torus. This time scale in coherent evolution of the torus might correspond to that of sub-bursts in the light curves of long GRBs (e.g. \\cite{van99}), and quasi-steady evolution of the torus on a time scale of at least tens of orbital periods follows from an upper bound of about 10\\% on the electromagnetic field energy it can support relative to its kinetic energy \\citep{van03a}. Our estimated TAMA 300 sensitivity distance for extracting time frequency trajectories is summarized in Table II. The results show $D\\simeq 0.070-0.10$ Mpc for black holes in the mass range $M=10-12 M_\\odot$, which compares favorably with the sensitivity distance for neutron star-neutron star coalescence. Extrapolation points to a sensitivity distance $D\\simeq 35-50$ Mpc for a strain noise amplitude of $2\\times 10^{-24}$ Hz$^{-\\frac{1}{2}}$ in the planned advanced detectors LIGO-Virgo and the LCGT. This sensitivity distance serves as a conservative estimate for the sensitivity distance for a detection. In particular, the sums of the 15 SNRi's shown in Fig. 4 are 95 and 77 for $D=0.05$ Mpc and, respectively, $D=0.07$ Mpc. These sums point to a sensitivity distance for a detection (with no particular information of behavior in the time frequency domain) on the order of a few tenths of Mpc, corresponding to sensitivity distance of a well over 100 Mpc for the advanced detectors. Based on the observed event rate of long GRBs of about 1 per year within a distance of 100 Mpc, the observable event rate suitable for extracting time frequency trajectories will depend on the abundance of the parent population of aspherical, relativistic and radio-loud Type Ib/c that may be powered by irradiation of the stellar envelope from a long-lived black hole inner engine \\citep{van03b}. It may reach one per few years if their event rate is about one order of magnitude larger than the rate of successful GRB-supernovae. The observable event rate could be larger if a fraction of the supernovae of Type II is similarly powered by long-lived black hole inner engines. A much larger sensitivity distance is anticipated for the planned 10 km ET in Europe. Conceivably, all sky radio surveys, e.g., the LOw Frequency ARray \\citep{lof10}, will further provide us with a probe for long duration radio bursts from mergers \\citep{van09b} and, combined with gravitational wave surveys, provide a direct measurement of the relative event rate long GRBs from mergers to long GRBs from CC-SNe. The diversity in the origin of long GRBs in both CC-SNe and mergers \\citep{van09b} and the comparible sensitivity distances of their potential emissions in gravitational waves and those from binary coalescence suggests the need for extended searches over the complete frequency range of both. Apart from scaling by black hole mass and a diversity in initial spin, the proposed long GWBs are universal, and their progenitors are revealed only by the absence or presence of a precursor signal in gravitational waves in case of, respectively, a CC-SNe or merger event. {\\bf Acknowledgments.} The initial work for this research was partially supported by La R\\'egion Centre during a visit to Le STUDIUM Institute for Advanced Studies/CNRS-Orl\\'eans. The authors gratefully acknowledge the TAMA collaboration for providing the data. MVP thanks Lars Hernquist and Ramesh Narayan for stimulating discussions." }, "1101/1101.2044_arXiv.txt": { "abstract": "% {The data reported in \\Planck's Early Release Compact Source Catalogue (ERCSC) are exploited to measure the number counts ($dN/dS$) of extragalactic radio sources at 30, 44, 70, 100, 143 and 217 GHz. Due to the full-sky nature of the catalogue, this measurement extends to the rarest and brightest sources in the sky. At lower frequencies (30, 44, and 70\\,GHz) our counts are in very good agreement with estimates based on \\textit{WMAP} data, being somewhat deeper at 30 and 70\\,GHz, and somewhat shallower at 44\\,GHz. \\Planck's source counts at 143 and 217\\,GHz join smoothly with the fainter ones provided by the SPT and ACT surveys over small fractions of the sky. An analysis of source spectra, exploiting \\Planck's uniquely broad spectral coverage, finds clear evidence of a steepening of the mean spectral index above about 70\\,GHz. This implies that, at these frequencies, the contamination of the CMB power spectrum by radio sources below the detection limit is significantly lower than previously estimated. } ", "introduction": "\\Planck\\footnote{\\Planck\\ (http://www.esa.int/\\Planck ) is a project of the European Space Agency (ESA) with instruments provided by two scientific consortia funded by ESA member states (in particular the lead countries: France and Italy) with contributions from NASA (USA), and telescope reflectors provided in a collaboration between ESA and a scientific consortium led and funded by Denmark.} \\citep{tauber2010a, planck2011-1.1} is the third-generation space mission to measure the anisotropy of the cosmic microwave background (CMB). It observes the sky in nine frequency bands covering 30--857\\,GHz with high sensitivity and angular resolution from 31\\arcm\\ to 5\\arcm. The Low Frequency Instrument (LFI; \\citealt{Mandolesi2010, Bersanelli2010, planck2011-1.4}) covers the 30, 44, and 70\\,GHz bands with amplifiers cooled to 20\\,\\hbox{K}. The High Frequency Instrument (HFI; \\citealt{Lamarre2010, planck2011-1.5}) covers the 100, 143, 217, 353, 545, and 857\\,GHz bands with bolometers cooled to 0.1\\,\\hbox{K}. Polarization is measured in all but the highest two bands \\citep{Leahy2010, Rosset2010}. A combination of radiative cooling and three mechanical coolers produces the temperatures needed for the detectors and optics \\citep{planck2011-1.3}. Two data processing centres (DPCs) check and calibrate the data and make maps of the sky \\citep{planck2011-1.7, planck2011-1.6}. \\Planck's sensitivity, angular resolution, and frequency coverage make it a powerful instrument for galactic and extragalactic astrophysics as well as cosmology. Early astrophysics results are given in Planck Collaboration, 2011h--z. The \\Planck\\ Early Release Compact Source Catalogue (ERCSC, \\cite{planck2011-1.10}) reports data on sources detected during the first 1.6 full-sky surveys, and thus offers, among other things, the opportunity of studying the statistical properties of extragalactic sources over a broad frequency range never fully explored by blind surveys. We will focus here on counts of extragalactic radio sources and on their spectral properties in the 30--217\\,GHz range.\\footnote{In all our calculations we have used the effective central frequencies for the \\Planck\\ channels \\citep{planck2011-1.4,planck2011-1.5}, although we indicate their nominal values. The most relevant difference is at 30\\,GHz, with a central frequency of \\getsymbol{LFI:center:frequency:30GHz}.} Although knowledge of the statistical properties at high radio frequency for this population of extragalactic sources has greatly improved in the recent past -- thanks to many ground-based observational campaigns and to the Wilkinson Microwave Anisotropy Probe (\\textit{WMAP}) surveys from space -- above about 70\\,GHz these properties are still largely unknown or very uncertain. This is essentially due to the fact that very large area surveys at mm wavelengths are made difficult by the small fields of view of ground-based radio telescopes and by the long integration times required. The most recent estimates on source number counts up to $\\sim50-70$ \\,GHz, and the optical identifications of the corresponding bright point sources (see, e.g., \\cite{Massardi08,Massardi10}), show that these counts are dominated by radio sources whose average spectral index is ``flat'', i.e., $\\alpha\\simeq 0.0$ (with the usual convention $S_\\nu\\propto\\nu^\\alpha$). This result confirms that the underlying source population is essentially made of Flat Spectrum Radio Quasars (FSRQ) and BL Lac objects, collectively called blazars,\\footnote{Blazars are jet-dominated extragalactic objects characterized by a strongly variable and polarized emission of the non-thermal radiation, from low radio energies up to the high energy gamma rays \\citep{UrryPadovani95}. Detailed analyses of Spectral Energy Distributions (SEDs) of complete blazar samples built by using simultaneous \\Planck\\ , Swift and Fermi data are given in \\citep{planck2011-6.3a}.} with minor contributions coming from other source populations \\citep{Toffolatti98,DeZotti05}. At frequencies $> 100$\\,GHz, however, there is now new information for sources with flux densities below about $1\\,$Jy coming from the South Pole Telescope (SPT) collaboration \\citep{Vieira10}, with surveys over 87 deg$^2$ at 150 and 220\\,GHz, and from the Atacama Cosmology Telescope (ACT) survey over 455 deg$^2$ at 148\\,GHz \\citep{Marriage10}. \\begin{figure*} \\begin{center} \\includegraphics[width=0.45\\textwidth]{./figs/16471fg1a.eps} \\includegraphics[width=0.45\\textwidth]{./figs/16471fg1b.eps} \\caption{Comparison between the ERCSC flux densities at 30\\,GHz (left panel) and at 44\\,GHz (right panel) with the almost simultaneous ATCA measurements (PACO project) at 32.2 and 39.7\\,GHz, respectively. No correction for the slightly different frequencies has been applied. \\label{fig:paco}} \\end{center} \\end{figure*} The ``flat'' spectra of blazars are generally believed to result from the superposition of different components in the inner part of AGN relativistic jets, each with a different synchrotron self-absorption frequency \\citep{KellermannPauliny-Toth69}. At a given frequency, the observed flux density is thus dominated by the synchrotron-emitting component which becomes self-absorbed and, in the equipartition regime, the resulting spectrum is approximately flat. However, this ``flat'' spectrum cannot be maintained up to very high frequencies, because of electron energy losses in the dominant jet-emission component (i.e., electron ageing), or the transition to the optically-thin regime, with the onset of a ``steep'' spectrum with a standard spectral index $\\alpha= -0.7$ to $-0.8$. A slightly steepened spectrum may also be caused by the superposition of many jet components. The redshift moves the observed steepening to lower frequencies and, thus, a greater fraction of blazar sources are observed with a steep spectrum at sub-mm wavelengths. With current data it is not yet possible to decide among the different scenarios. However, given their sensitivity and full sky coverage, \\Planck\\ surveys are uniquely able to shed light on this transition from an almost ``flat'' to a ``steep'' regime in the spectra of blazar sources. The outline of this paper is as follows. In \\S\\,\\ref{sec:over} we briefly sketch the main properties of the ERCSC. In \\S\\,\\ref{sec:valid} we summarize the source validation. In \\S\\,\\ref{sec:sample} we describe the complete sample selected at 30\\,GHz, used for the analysis of spectral properties. In \\S\\,\\ref{sec:counts} we present the source counts over the frequency range 30--217\\,GHz. In \\S\\,\\ref{sec:spectra} we investigate the spectral index distributions in different frequency intervals. Finally, in \\S\\,\\ref{sec:concl} we summarize our main conclusions. ", "conclusions": "\\label{sec:concl} We have exploited the ERCSC to estimate the bright counts of extragalactic radio sources at 6 frequencies (30, 44, 70, 100, 143, and 217\\,GHz) and to investigate the spectral properties of sources in a complete sample selected at 30\\,GHz. The counts at 30, 44, and 70\\,GHz are in good agreement with those derived from \\textit{WMAP} data at nearby frequencies. The completeness limit of the ERCSC is somewhat deeper than that of \\textit{WMAP} at 30 and 70\\,GHz and somewhat shallower at 44\\,GHz. At higher frequencies the ERCSC has allowed us to obtain the first estimate of the differential number counts at bright flux density levels. At 30, 143 and 217\\,GHz, the present counts join smoothly to those from deeper surveys over small fractions of the sky. The \\cite{DeZotti05} model is consistent with the present counts at frequencies up to 100\\,GHz, but over-predicts the counts at higher frequencies by a factor of about 2.0 at 143\\,GHz and about 2.6 at 217\\,GHz. This implies that the contamination of the CMB power spectrum by radio sources below the 1\\,Jy detection limit is lower than previously estimated. No significant changes are found, however, if we consider fainter source detection limits, i.e., 100\\,mJy, given the convergence between predicted and observed number counts. The analysis of the spectral index distribution over different frequency intervals, within the uniquely broad range covered by \\Planck\\ in the mm and sub-mm domain, has highlighted an average {\\it steepening} of source spectra above about 70\\,GHz. The median values of spectral indices vary from $\\alpha_{30}^{70}= -0.18 \\pm 0.01$ ($\\sigma= 0.18$) to $\\alpha_{70}^{143}= -0.52 \\pm 0.01$ ($\\sigma = 0.22$). This steepening accounts for the discrepancy between the \\cite{DeZotti05} model predictions and the observed differential number counts at HFI frequencies. The current outcome is also in agreement with the findings of \\cite{planck2011-6.3a} on a complete sample of blazars selected at 37\\,GHz. The change detected in the spectral behaviour of extragalactic radio sources in the \\Planck\\ ERCSC at frequencies above 70-100\\,GHz can be tentatively explained by electron ageing or by the transition to the optically thin regime, predicted in current models for radio emission in blazar sources. However, with present data it is not yet possible to clarify the situation. In the near future, the data of the \\Planck\\ Legacy Survey will surely prove very useful in settling this open issue." }, "1101/1101.0954.txt": { "abstract": "{The detection of small mass planets with the radial-velocity technique is now confronted with the interference of stellar noise. HARPS can now reach a precision below the meter-per-second, which corresponds to the amplitudes of different stellar perturbations, such as oscillation, granulation, and activity. } {Solar spot groups induced by activity produce a radial-velocity noise of a few meter-per-second. The aim of this paper is to simulate this activity and calculate detection limits according to different observational strategies.} {Based on Sun observations, we reproduce the evolution of spot groups on the surface of a rotating star. We then calculate the radial-velocity effect induced by these spot groups as a function of time. Taking into account oscillation, granulation, activity, and a HARPS instrumental error of 80\\,cm\\,s$^{-1}$, we simulate the effect of different observational strategies in order to efficiently reduce all sources of noise.} {Applying three measurements per night of 10 minutes every three days, 10 nights a month seems the best tested strategy. Depending on the level of activity considered, from $\\log{R'_{HK}}= -5$ to $-4.75$, this strategy would allow us to find planets of 2.5 to 3.5\\,M$_{\\oplus}$ in the habitable zone of a K1V dwarf. Using Bern's model of planetary formation, we estimate that for the same range of activity level, 15 to 35\\,\\% of the planets between 1 and 5\\,M$_{\\oplus}$ and with a period between 100 and 200 days should be found with HARPS. A comparison between the performance of HARPS and ESPRESSO is also emphasized by our simulations. Using the same optimized strategy, ESPRESSO could find 1.3\\,M$_{\\oplus}$ planets in the habitable zone of K dwarfs. In addition, 80\\,\\% of planets with mass between 1 and 5\\,M$_{\\oplus}$ and with a period between 100 and 200 days could be detected.} {} %\\abstract{The detection of small mass planets, using the radial-velocity technique, is now getting confronted to stellar noise. HARPS can now reach a precision below the meter per second, which corresponds to the amplitudes of different stellar perturbations, such as oscillation, granulation and activity. % %In this paper, we present simulations, based on Sun observations, that reproduce the evolution of spot groups on the surface of a rotating star. The radial-velocity effect induced by these spot groups as a function of the time is then calculated. Taking into account oscillation, granulation, activity and an HARPS instrumental error of 80 cm\\,s$^{-1}$, we have tested different observational strategies in order to reduce at best all source of noise. %Applying 3 measurements per night of 10 minutes, every 3 days, 10 nights a month seems to be the best tested strategy considering a reasonable total observation time. This strategy would allow to find planets of 2.5 to 3.5\\,M$_{\\oplus}$ in the habitable zone of a K1V star such as $\\alpha$\\,Cen\\,B, depending on the level of activity considered ($-5<\\log(R'_{HK})<-4.75$). Using Bern's model of planetary formation, we estimate that for the same range of activity level, 15 to 35\\,\\% of the planets between 1 and 5\\,M$_{\\oplus}$ and with a period between 100 and 200 days should be found using HARPS. A comparison between the performance of HARPS and ESPRESSO is also emphasize by our simulations. Using the same optimized strategy, ESPRESSO could find planet of 1.3\\,M$_{\\oplus}$ in the habitable zone of K dwarfs and 80\\,\\% of the planets with a mass between 1 and 5\\,M$_{\\oplus}$ and with a period between 100 and 200 days. ", "introduction": "We have now discovered more than 400 exoplanets with the radial-velocity (RV) technique\\footnote{see The Extrasolar Planets Encyclopaedia, http://exoplanet.eu}. The majority of them are very massive and on very short period orbits, something that was quite unexpected from the theories of giant-planet formation. However, since a few years, giant planets much more similar to the solar system giants \\citep[e.g.][]{Wright-2008} as well as super-Earth planets have been detected (mass from 2 to 10 M$_{\\oplus}$) \\citep[e.g.][]{Mayor-2009a, Mayor-2009b, Udry-2007b}. This has become possible thanks to three important improvements of the RV technique. First, the precision of spectrographs improved considerably, reaching now a 100\\,cm\\,s$^{-1}$ precision level \\citep[typically 100 cm\\,s$^{-1}$ in 1 minute for a $V=7.5$ K0 dwarf, using the HARPS spectrograph on the ESO 3.6 meter telescope, see][]{Pepe-2005}. A second improvement has been achieved through an appropriate observational strategy, which allows us to average out perturbations caused by stellar oscillations \\citep[][]{Santos-2004a}. Finally, several years of follow-up helped us find long-period planets as well as very small mass ones. At the 100\\,cm\\,s$^{-1}$ level of accuracy, we start to be confronted with noises caused by the stars themselves. Stellar noise is the result of three types of perturbation produced by three different physical phenomena: oscillations, granulation, and magnetic activity. Oscillations of solar type stars, which can be seen as a dilatation and contraction of external envelopes over timescales of a few minutes \\citep[5 minutes for the Sun;][]{Schrijver-2000,Broomhall-2009}, is caused by pressure waves (p-modes) that propagate at the surface of solar type stars. The individual amplitudes of p-modes are typically from a few to tens of cm\\,s$^{-1}$, but the interference of tens of modes with close frequencies introduce RV variations of several cm\\,s$^{-1}$, depending on the star's spectral type and evolutionary stage \\citep[][]{Bedding-2007a,Bouchy-2003,Bedding-2003,Schrijver-2000}. The amplitude and period of the oscillation modes increase with mass along the main sequence. Theory and observations show that the frequencies of the p-modes rise with the square root of the mean density of the star and that their amplitudes are proportional to the ratio of the luminosity over the mass \\citep[][]{Christensen-Dalsgaard-2004,OToole-2008}. The convection in external layers of solar type stars drives different phenomena of granulation (granulation, mesogranulation, and supergranulation), which also affect the RV measurements on time scales going from several minutes to several hours. In order of timescale and size of convective pattern, we first have granulation, whose typical timescale is shorter than 25 minutes \\citep[][]{Title-1989,Del_Moro-2004b}. Then comes mesogranulation and finally supergranulation, with timescales of up to 33 hours \\citep[][]{Del_Moro-2004a}. When integrated over the entire stellar disk, all these convection perturbations present a noise level on the order of the meter-per-second. On longer timescales, similar to the star rotational period, the presence of activity related spots and plages perturbs precise RV measurements. Spots and plages on the surface of a star will break the flux balance between the red-shifted and the blue-shifted halves of the star. As the star rotates, a spot group, or a plage, moves across the stellar disk and produces an apparent Doppler shift \\citep[][]{Saar-1997,Queloz-2001,Huelamo-2008,Lagrange-2010}. This effect can be hard to distinguish from the signal caused by the presence of a planet. For the Sun at maximum activity of cycle 23, \\citet[][]{Meunier-2010a} find a noise related to spot groups and plages of 42\\,cm\\,s$^{-1}$. Because the temperature of spot groups and plages are different from the mean stellar surface, and because active regions contain both, the noise induced by them will usually be compensated, but not entirely, because the surface ratio between spots and plages varies \\citep[e.g.][]{Chapman-2001}. According to \\citet{Meunier-2010a}, the major perturbative effect of activity on RVs is not the one induced directly by spot groups and plages, but the one caused by the inhibition of convection in active regions \\citep[e.g.][]{Dravins-1982,Livingston-1982,Brandt-1990,Gray-1992}. This effect leads to a blueshift distortion of spectrum lines, which results in a noise varying from 40\\,cm\\,s$^{-1}$ at minimum activity to 140\\,cm\\,s$^{-1}$ at maximum. The Earth produces a radial-velocity perturbation of 9\\,cm\\,s$^{-1}$ on the Sun, which would be completely masked by the stellar noise. If we manage to understand the structure of these different kinds of noise, we could use appropriate observational strategies to reduce the stellar noise contribution as much as possible, and thus find very small mass planets far from their host star. This investigation of optimized observational strategies is essential for future and more accurate instruments, such as ESPRESSO@VLT (http://espresso.astro.up.pt/, precision expected: 10\\,cm\\,s$^{-1}$) or CODEX@E-ELT \\citep[e.g.][precision expected: 2\\,cm\\,s$^{-1}$]{Pasquini-2008}, in order to detect Earth twins (Earth-mass planets in habitable regions). The two first types of noise, which are caused by oscillations and granulation phenomena, have been discussed in a previous paper \\citep[][hereafter Paper I]{Dumusque-2010a}. Starting from HARPS asteroseismology measurements, we characterized these two kinds of noise using the power spectrum representation. Then we derived an optimized observational strategy, reducing at best the stellar noise contribution while keeping a reasonable total observational time. In the present paper, we continue this first study and include the noise coming from magnetic activity spot groups. Starting from observations of the Sun, which is the only star where we can resolve spot groups, we simulate the appearance of spot groups on its surface and calculate the RV contribution. Adding the noise coming from magnetic activity to the results obtained in Paper I, we develop an observational strategy that simultaneously reduces the three kinds of stellar noise. To finish, we calculate the corresponding detection limits, as well as the expected number of planets that could be found by comparing our results with Bern's model of planetary formation. ", "conclusions": "In Paper I we had proposed an efficient observational strategy to reduce the stellar noise generated by oscillation and granulation as much as possible. This strategy, requiring three measurements per night of 10 minutes over 10 consecutive days each month, improves the averaging out of stellar noise coming from oscillation and granulation by 30\\,\\%, with an observational cost only multiplied by a factor 2. In the present paper, we add the noise induced by activity-related stellar spot groups. To study the radial-velocity (RV) effect caused by stellar spot groups, we consider three different activity levels. Based on Sun observations, we simulate the minimum solar activity level ($\\log(R'_{HK})=-5$), the maximum one ($\\log(R'_{HK})=-4.75$), and an intermediate one ($\\log(R'_{HK})=-4.9$). Comparing the RV variation at maximum activity given by our simulation with the one calculated by real position and size of spot groups from cycle 23 \\citep{Meunier-2010a}, we find a very good agreement, 51\\,cm\\,s$^{-1}$ and 48\\,cm\\,s$^{-1}$, respectively. We note that the RV effect of activity is not fully simulated because the inhibition of convection in active regions is not yet implemented. According to \\citet{Meunier-2010a}, this effect could be important and consequently, the detection limits calculated here could be underestimated. Compared to the results of \\citet{Meunier-2010a}, our simulation is more general in the sense that we can predict the RV effect of activity related spots groups for other stars, knowing a few characteristics (mean spot number, distribution of spot groups lifetime, presence of active longitudes, differential rotation). This will be very important when a better knowledge of activity phenomena of other stars than the Sun will be known. The Kepler mission should give us some clues in the near future. Modeling the short-term activity with a three sine waves function with period $P_{fit}$, $P_{fit}/2$ and $P_{fit}/3$ can greatly reduce the spot-induced noise by approximately 70\\,\\%. However, we have to be very careful with this method because $P_{fit}$ can vary in a non negligible way from the rotational period of the star. Thus, signal of small mass planets with period similar to the rotational period of the star will be killed by this type of model. Only a study of other CCF parameters such as the FWHM or the BIS can give us clues on the true nature of the signal: short-term activity or planet. After simulating the effect of several observational strategies, the most efficient one to average out all kinds of noise is the 3N3. This strategy consists in measuring the star three times 10 minutes per night, with a spacing of two hours. Then the measurements are taken every third night, 10 days every month. With this strategy it would be possible to find planets of 2.5 to 3.5 Earth mass with HARPS in the habitable region of K dwarfs (200 days). The first mass value corresponds to a case without activity ($\\log(R'_{HK})=-5$) and the second one to a maximum activity level ($\\log(R'_{HK})=-4.75$). Even if the activity caused by spot groups introduces a non negligible noise, small mass planets in habitable regions could be detected with HARPS with an appropriate observational strategy. %To take into account this activity effect, based on Sun observations, we simulate the evolution of spot groups on the surface of a rotating star. Since the number of spot groups change during the 11-year activity cycle, we made 2 simulations, one for a $\\log(R'_{HK})$ equal to $-4.9$ and another one for a $\\log(R'_{HK})$ equal to $-4.75$. The RV effect induced by these spots was calculated and we found that the period of activity perturbations is between 10 and 30 days. Since this period is larger than the perturbation period for oscillation and granulation, the best strategy derived in Paper I is no more suitable when the activity rise up. Keeping the 3 measurements per night of 10 minutes strategy to average out oscillation and granulation, we found that the best strategy, with an equal total observation time, was to observe the star each 3 nights, 10 days a month. This new strategy manage to lower the detection limits up to 50\\,\\% for long periods. %The presence of activity is a problem for planets presenting a period similar to the typical time scales of activity effects, 10 to 50 days. In the case of shorter period planets, the RV signal will not be disturb by activity perturbations and in the case of longer period planets, we can use a huge binning, which will average out activity effect without reducing the RV amplitude induced by the planets. It is thus easier to average out activity effects for long period planets than for intermediate period ones (10 to 50 days) Using the population of low mass planets predicted by Bern's model, we calculate the proportion of planets that could be found with the 3N3 strategy. Using HARPS, it would be possible to find 35\\,\\% of the planets below $5\\,M_{\\oplus}$ with a period between 100 and 200 days and a low-activity level ($\\log(R'_{HK})=-5$). This value decreases to 15\\,\\% for a $\\log(R'_{HK})$ equal to $-4.75$. If we trust this model of formation, which is the most compatible with observations, HARPS could discover several planets below $5\\,M_{\\oplus}$ in the habitable region of early-K dwarfs. In our simulation, the 3N3 strategy appears to be the best one to average out activity noise. This is because of the rotational period of the Sun, which is fixed in our simulation at 26 days. Indeed, when choosing 10 observational nights a month, the best way to sample the entire rotational period is to observe the star every third night. For shorter rotational periods, increasing the measurement frequency would lead to a better averaging. For longer rotational periods, reducing the measurement frequency is not recommended because short period planets will be poorly sampled. Taking three measurements of 10 minutes per night gives us a very low photon noise when binning the data over the night. Thus, the photon-noise is not a limitation any longer, and 2 to 4-meters class telescopes can still be used to go down to the 10\\,cm\\,s$^{-1}$ level on bright stars. However, spectrographs must also intrinsically reach this level of precision and for the moment, only ESPRESSO@VLT (http://espresso.astro.up.pt/) is designed to reach this goal. With such a precision, improvements would be remarkable. The detection limits would go down in mass by a level of 1 $M_{\\oplus}$, reaching 1.3 $M_{\\oplus}$ in the habitable region of early-K dwarfs. In addition, ESPRESSO would find 80\\,\\% of the planets between 1 and 5 $\\,M_{\\oplus}$ and with a period between 100 and 200 days, which is more than twice as much as what HARPS could find. Moreover, CODEX@E-ELT \\citep[e.g][]{Pasquini-2008}, would improve the long term instrumental precision down to a few cm\\,s$^{-1}$ level and would detect $1\\,M_{\\oplus}$ planets in habitable regions." }, "1101/1101.4662_arXiv.txt": { "abstract": "Observations with the \\swift\\ satellite of \\x\\ afterglows of more than a hundred gamma ray bursts (GRBs) with known redshift reveal ubiquitous soft X-ray absorption. The directly measured optical depth $\\tau$ at a given observed energy is found to be constant on average at redshift $z > 2$, i.e., $\\langle \\tau (0.5~\\mathrm{keV}) \\rangle_{z > 2}\\, = 0.40\\pm 0.02$. Such an asymptotic optical depth is expected if the foreground diffuse intergalactic medium (IGM) dominates the absorption effect, and if the metallicity of the diffuse IGM reaches $\\sim 0.2 - 0.4$ solar at $z = 0$. To further test the IGM absorption hypothesis, we analyze the 12 highest S/N ($> 5000$ photon) $z > 2$ quasar spectra from the \\xmm\\ archive, which are all extremely radio loud (RLQs). The quasar optical depths are found to be consistent with the mean GRB value. The four lowest-$z$ quasars ($2 < z < 2.5$), however, do not show significant absorption. The best \\x\\ spectra of radio-quiet quasars (RQQs) at $z > 2$ provide only upper limits to the absorption, which are still consistent with the RLQs, albeit with much lower S/N ($\\lsim 1000$ photons at $z \\approx 4$). Lack of quasar absorption poses a challenge to the smooth IGM interpretation, and could allude to the opacity being rather due to the jets in RLQs and GRBs. However, the jet absorbing column would need to appear in RLQs only at $z~\\gsim\\ 2.5$, and in GRBs to strongly increase with $z$ in order to produce the observed tendency to a constant mean $\\tau$. High \\x\\ spectral resolution can differentiate between an absorber intrinsic to the source that produces discernible spectral lines, and the diffuse IGM that produces significant absorption, but no discrete features. ", "introduction": "\\label{into} \\swift\\ \\x\\ spectra of gamma ray bursts (GRB) afterglows reveal prevalent soft X-ray absorption, which is commonly assumed to originate in the host galaxy. The standard \\x\\ absorption measurement technique probes metal absorption, but quotes equivalent hydrogen column densities $N_H$ by assuming a neutral absorber at the host ($z$), and with solar abundances. Under these assumptions, $N_H$ towards GRBs shows a strong correlation with the host redshift $z$. The typical absorbing column rises from $N_{\\rm H}\\sim 10^{21}$~cm$^{-2}$ at $z<1$ up to $\\approx 10^{23}$~cm$^{-2}$ at the highest observed redshifts \\citep{Jakobsson2006, Campana2006, Watson2007, Campana2010, Rau2010}. Damped Lyman-$\\alpha$ absorption is also seen in some GRB afterglows. Unlike the \\x\\ absorber, the redshift of the Lyman-$\\alpha$ absorber is well constrained. The implied column is usually still well below the X-ray derived column \\citep[see, e.g.,][]{Watson2007}. If the abundances in the \\x\\ absorber are sub-solar, the \\x\\ derived $N_H$ values are even higher and the discrepancy with the Lyman-$\\alpha$ column grows accordingly. Although an appreciable ionization range in a single medium can possibly account for such a discrepancy \\citep[recently,][]{schady10}, there is also the possibility that the \\x\\ and Lyman-$\\alpha$ absorbers are physically distinct. High-$z$ {\\it quasars} also commonly reveal soft X-ray absorption \\citep[e.g.,][]{Fabian2001, Worsley2004a, Worsley2004b, Page2005, Yuan2005, Grupe2006, sambruna07}. The implied absoring column, if intrinsic, is also of the order of $10^{23}$~cm$^{-2}$. In the case of quasars as well, metal and hydrogen line absorption in the UV imply significantly lower $N_H$ columns. In the case of low-luminosity active galaxies, partially ionized outflows are known to have more \\x\\ column with only a trace of UV absorbing ions \\citep{crenshaw03}, a discrepancy that again can be partially reconciled with an ionization correction. The X-ray absorption profile in three different quasars at $z\\sim 4.3-4.7$ was noted by \\citet{Yuan2005} to be remarkably similar, while ionized quasar outflows are not necessarily expected to be so uniform. Indeed, partially ionized outflows have not been directly identified in such luminous quasars, and thus it is plausible that the \\x\\ and UV absorbers towards high-$z$ quasars are also physically distinct. Intrigued by the aforementioned puzzles associated with soft \\x\\ absorption of high-$z$ sources, we wish to explore the similarities of \\x\\ absorption of high-z GRBs and quasars and their possible origin. In Sec.~\\ref{GRB} we present the \\x\\ opacities towards GRB afterglows from the \\swift\\ sample \\citep{Evans2009, Campana2010}, but now without assuming neither a redshift, nor an ionization, nor a metallicity for the \\x\\ absorber. We then discuss in Sec.~\\ref{IGM} the viability of \\x\\ absorption by the diffuse intergalactic medium (IGM). In Sec.~\\ref{quasars}, we present a comparison sample of the highest signal to noise ratio (S/N) $z > 2$ quasar spectra drawn from the \\xmm\\ archive. Realizing that our selection criterion based on the number of detected photons allows only radio loud quasars into the sample, in Sec.~\\ref{RQQ} we also present the best S/N high-z radio quiet quasars from the works of \\citet{shemmer08,shemmer06,shemmer05}. In Sec.~\\ref{concl} we conclude and compare the expected absorption signature in the \\x\\ spectra of high-$z$ sources of a diffuse IGM versus that of jets and propose future tests that can reveal whether the observed absorption is indeed due to the IGM or intrinsic to the sources. ", "conclusions": "\\label{concl} The extragalactic soft \\x\\ opacity towards GRBs and quasars at low $z$ is dominated by absorption in the host galaxy. However, the contribution of a fixed host column to the optical depth $\\tau$ at a given observed energy sharply drops with $z$, and the extragalactic opacity observed in the soft X-ray spectra of high-$z$ GRBs may thus have a dominant contribution from the diffuse IGM. Due to the redshift and energy dependence of the photo-electric cross section, the diffuse IGM opacity saturates at $z~\\gsim\\ 2$ and is expected to be isotropic and relatively constant beyond that redshift. The redshift where saturation occurs depends on the IGM metallicity. The IGM opacity is not expected to directly correlate with the line absorption systems observed in the optical and UV that originate either in the host galaxy or in over-dense IGM clumps. A large sample of \\swift\\ GRBs and a high S/N sample of RLQs observed with \\xmm\\ appear to mostly be consistent with this picture of diffuse IGM absorption. However, a few RLQs at $z < 2.5$ and a few low S/N RQQs with very little absorption that nonetheless could represent a much larger quasar sample raise doubts. For the IGM to produce the \\x\\ opacities observed towards high-$z$ GRBs and RLQs, it needs to be relatively enriched with metals and not too highly ionized. If indeed it is the diffuse IGM responsible for this absorption, it can account for a significant fraction of the currently missing baryons, implied by big bang nucleosynthesis \\citep{Steigman2007}, the observed angular power spectrum of the cosmic microwave background radiation and the Thomson opacity inferred from its polarization \\citep{Komatsu2010}. Of these baryons, in the local universe only $\\sim 50\\%$ are present in the galaxies, galaxy clusters and UV-optical IGM line systems known to date \\citep[for a review see][]{bregman07}. Furthermore, if 90\\% of the baryons are in the IGM and if the diffuse low-$z$ IGM metallicity is indeed $\\sim 0.2 - 0.4$ solar as we postulate here, then the IGM contains the bulk of the metals in the present-day universe, compared to, say, solar metallicity in stars and galaxies that comprise only 10\\% of the baryons. The presence of absorption in GRBs and RLQs, that both harbor powerful jets, but less in RQQs, also raises the possibility that the absorption effect has something to do with the jet. It was discovered by \\rosat\\ \\citep{elvis94, fiore98} that high-$z$ RLQs are much more absorbed in the \\x s than RQQs, which seems to support a jet effect. On the other hand, these authors also realized the absorption (if intrinsic) increases with $z$ and not with luminosity, which argues against a jet-physics origin. The fact that local RLQs generally do not have the high column densities that are found in the high-$z$ sources has been confirmed recently by \\citet[][Table~4 therein]{galbiati05}. Indeed, ascribing photo-electric absorption to the jet is counter-intuitive, as the jet is not expected to comprise atomic material with bound electrons, and especially not metals, that produce the observed \\x\\ opacity. Moreover, the little scatter of the observed optical depth found in this paper, and even more the tendency to an asymptotic opacity at high $z$ that require a putative intrinsic column to scale approximately with $(1+z)^{2.5}$ to offset the decreasing cross section, calls into question the realistic role jets can play in determining the soft \\x\\ opacity. Turning to the trends of intergalactic line absorption towards GRBs and quasars in the optical band does not provide a clear-cut answer to the uniqueness of absorption towards jetted sources either. On one hand, the number of Mg\\,II intervening absorption systems towards GRBs \\citep{prochter06} and blazars \\citep{bergeron10} is a few times higher than that towards RQQs. As there is no obvious reason for line sights towards GRBs to be different from those towards quasars, it would require metals to be entrained in the jets of GRBs and RLQs. On the other hand, line sights towards high-$z$ RLQs seem to have a similar number density of intervening DLA systems as those towards optically-selected (again, mostly RQQs) samples \\citep{ellison01}. Prospectively, there is a clear way to distinguish a well confined absorber from a cosmologically diffuse one based on line absorption. The detection of absorption lines, or lack thereof at high spectral resolution (and high S/N) would provide a definitive characterization of the absorber. The CCD \\x\\ detectors used in this work do not have the spectral resolution required to discern lines, while the grating spectrometers on board \\xmm\\ do not have the sufficient effective area to provide adequate spectra, at least not with the modest exposures available to date. High S/N high-resolution spectra of the quasars either through a long \\xmm\\ exposure, or with future instruments should be able to unambiguously determine whether \\x\\ absorption of high-$z$ sources is due to their hosts or due to the diffuse IGM." }, "1101/1101.1337_arXiv.txt": { "abstract": "Nuclear regions of galaxies generally host a mixture of components with different exitation, composition, and kinematics. Derivation of emission line ratios and kinematics could then be misleading, if due correction is not made for the limited spatial and spectral resolutions of the observations. The aim of this paper is to demonstrate, with application to a long slit spectrum of the Seyfert~2 galaxy NGC~1358, how line intensities and velocities, together with modelling and knowledge of the point spread function, may be used to resolve the differing structures. In the situation outlined, the observed kinematics differs for different spectral lines. From the observed intensity and velocity distributions of a number of spectral lines and with some reasonable assumptions to diminish the number of free parameters, the true line ratios and velocity structures may be deduced. A preliminary solution for the nuclear structure of NGC~1358 is obtained, involving a nuclear point source and an emerging outflow of high excitation with a post shock cloud, as well as a nuclear emission line disk rotating in the potential of a stellar bulge and expressing a radial exitation gradient. The method results in a likely scenario for the nuclear structure of NGC~1358. For definitive results an extrapolation of the method to two dimensions combined with the use of integral field spectroscopy will generally be necessary. ", "introduction": "The nuclear and circumnuclear activity of galaxies generally involves the interplay between a number of different components and phenomena, e.g. a central active source surrounded by an absorbing torus, a rotating central bulge, outflowing jets, flowing streams due to the action of a bar, or even merging. To separate these different components and derive their respective line ratios and kinematic behaviour is generally difficult due to the limited spatial and spectroscopic resolution available. On the other hand, such a separation is crucial to the analysis of the structures and physical processes involved in the nuclear region and their roles in galaxy evolution. Then, as to be demonstrated here, such a separation could benefit by considering the differences of the velocities observed for different spectral lines and be eased by models of the activity, smoothed with the point spread function (PSF) and fitted to the observations. Evidently, to yield the best results this should require an integral field spectrum obtained at the best available spatial and spectral resolution. NGC~1358 is a barred Sa galaxy hosting an active galactic nucleus at a heliocentric velocity of $\\approx$4100~\\kms, included in the sample of Ulvestad \\& Wilson (1989) as a Seyfert~2. It attracted our attention because of its remarkable circumnuclear kinematics (Dumas et al. 2007, Lindblad et al. 2010). ", "conclusions": "Our model (Fig.~\\ref{fig:model}) resolves the nuclear structure of NGC~1358 into (1) a central unresolved emission line source (the Nucleus), (2) a 1\\farcs3 long jet emerging from the nucleus with a line of sight velocity of $-214$ \\kms, (3) a spherical nuclear stellar bulge containing a rotating emission line disk inclined to the stellar kinematic symmetry plane of the bulge, (4) an emission line region (the Cloud) outside the jet with positive velocity as the emission line disk. There is no room for and no need for a counter-jet in our model. Then the true nuclear redshift must be higher than the observed one. This best fit gives an observed redshift of the central source of 4100 \\kms, or 4089 \\kms\\ with heliocentric velocity correction. The total $M \\times \\sin^{2}(i)$ of the bulge is determined essentially by the velocities of H$\\alpha$ and \\NII\\ at the distance $-6$\\arcsec from the centre. The best fit gives $M \\times \\sin^{2}(i) = 14 \\times 10^{9}~\\rm M_{\\sun}$. Figure~\\ref{fig:model} shows that the velocities and intensities have been reproduced with fair accuracy within this nuclear region. The H$\\beta$ intensities are too weak and uncertain to give a unique solution in the present case. With the emission line fluxes for these structures we are able to set up the Starburst-AGN diagnostic diagrams introduced by Baldwin, Phillips \\& Terlevich (1981). Figure~\\ref{fig:lineratios} shows the diagnostic diagram \\OIII/H$\\beta$ versus \\NII/H$\\alpha$ where the \\OIII/H$\\beta$ ratio is derived from \\OIII/H$\\alpha$ assuming large optical depth, referred to as Case B described in chapter 4.2 of Osterbrock (1989). Accordingly, for an electron temperature of 10000 K, the H$\\alpha$/H$\\beta$ is 2.87, which we note is not much different from the case when the gas is assumed to be optically thin. As evident, the nuclear source and the jet fall in the region of highly excited AGNs, while the inner disk and the cloud fall close to the region of liners. The ratios \\NII/H$\\alpha$ places the outer part of the nuclear disk and the spiral arms in the HII region domain. The contribution from the underlying H$\\alpha$ and H$\\beta$ absorption lines only change the derived emission line ratios marginally. However, the effect of reddening has not been taken into account in the emission ratios presented in Fig.~\\ref{fig:lineratios}. As we have used the unreddened ratio H$\\alpha$/H$\\beta$, our derived \\OIII/H$\\beta$ ratios presented in the diagram are lower limits in the presence of reddening. This is an effect that favours our conclusions about the nature of the components (further details in Lindblad et al. 2010). \\begin{figure} \\begin{center} \\includegraphics[width=.49\\textwidth]{Fathi_NGC1358_Fig2.jpg} \\caption{Line ratio diagnostics derived from our ESO spectra. The solid curve shows the Starburst-AGN separation lines of Kewley et al. (2001), with the shaded region indicating their sample galaxies.} \\label{fig:lineratios} \\end{center} \\end{figure} We argue that, with this analysis, we have made a step towards a resolution of the nuclear region of NGC~1358 into a number of different components with different velocity behaviour and excitation, where a summary of the results are seen in Fig.~\\ref{fig:lineratios}. The differing velocities in the different lines give strong constraints for the line ratios. Experiments with our model, for example, shows that it is the weakness of \\OIII\\ in the Cloud and nuclear disk that causes the dip in velocity of the Eastern part of the rotation curve to be deeper in \\OIII\\ and the \\OIII\\ rotation curve to be flatter on both sides. The main ambiguity with the model presented here is the difference in systemic velocity of the bulge and outer disk relative to the nucleus and nuclear emission line disk, as well as the relation of the Cloud to the emission line disk and to the bulge. The Cloud, although introduced in a somewhat ad hoc fashion, could be a component that resides in the emission line disk. Of course, to get all necessary information about the velocities and spatial distribution of the different components from an isolated spectrum is futile, because i.a. of the unknown influence due to the point spread function from sources outside the slit. By adding more slits, there is still a difficulty to obtain an accurate and complete coverage (Lindblad et al. 1996). Obviously, the ideal is an integral field spectrometer and an extension of the method to a two-dimensional treatment covering the entire region of interest." }, "1101/1101.3935_arXiv.txt": { "abstract": "Phenomenological relations exist between the peak luminosity and other observables of type Ia supernovae (SNe~Ia), that allow one to standardize their peak luminosities. However, several issues are yet to be clarified: SNe~Ia show color variations after the standardization. Also, individual SNe~Ia can show residuals in their standardized peak absolute magnitude at the level of $\\sim 0.15$ mag. In this paper, we explore how the color and luminosity residual are related to the wavelength shift of nebular emission lines observed at $\\gsim 150$ days after maximum light. A sample of 11 SNe Ia which likely suffer from little host extinction indicates a correlation ($3.3\\sigma$) between the peak $B-V$ color and the late-time emission-line shift. Furthermore, a nearly identical relation applies for a larger sample in which only three SNe with $B-V \\gsim 0.2$ mag are excluded. Following the interpretation that the late-time emission-line shift is a tracer of the viewing direction from which an off-centre explosion is observed, we suggest that the viewing direction is a dominant factor controlling the SN color and that a large part of the color variations is intrinsic, rather than due to the host extinction. We also investigate a relation between the peak luminosity residuals and the wavelength shift in nebular emission lines in a sample of 20 SNe. We thereby found a hint of a correlation (at $\\sim 1.6 \\sigma$ level). The confirmation of this will require a future sample of SNe with more accurate distance estimates. Radiation transfer simulations for a toy explosion model where different viewing angles cause the late-time emission-line shift are presented, predicting a strong correlation between the color and shift, and a weaker one for the luminosity residual. ", "introduction": "Type~Ia supernovae (SNe Ia) are used to measure cosmological parameters and study the nature of the dark energy (see Leibundgut\\ 2008, and references therein). Thanks to the uniformity of their peak luminosities, once a phenomenological relation between the light-curve shape and the peak luminosity is applied (hereafter the light-curve correction, or the Phillips relation), they can be accurately used as cosmological standard candles (Phillips\\ 1993; Hamuy et al.\\ 1996; Phillips et al.\\ 1999). Their colors are also known to correlate with light-curve shape (Tripp 1998; Tripp \\& Branch 1999; Phillips et al.\\ 1999). In addition to light-curve shape and color, several other observables, mostly related to spectral features, have been shown to correlate with the SN peak luminosity (e.g., Nugent et al.\\ 1995; Mazzali et al.\\ 1998; Benetti et al.\\ 2005; Bongard et al.\\ 2006; Hachinger et al.\\ 2006; Foley et al.\\ 2008). Currently there are several issues yet to be clarified. One central issue is that the intrinsic color variations of SNe~Ia have not yet been fully understood. After application of the relation between color and light curve shape, there remain variations in the color excess of SNe Ia. So far, it has been practically impossible to discriminate between the contributions from a possible `residual' intrinsic color, which does not correlate with the light curve shape, and that from the extinction within the host or the environment around the SN. This issue could be related to the fact that when the dispersion of the Hubble diagram is minimized with $R_{V}$ being treated as a free parameter, one obtains low values of $R_{V}$ between $1 - 2$. However, as shown by Folatelli et al. (2010; hereafter F10), when one compares the colors or color excesses of normal SNe~Ia, a more typical Milky Way-like value of the reddening law is obtained, i.e. $R_{V} \\sim 3$. F10 argued that this apparent discrepancy suggests that there is an intrinsic color variation within SNe~Ia that correlates with luminosity, but is independent of the light curve decline-rate $\\Delta m_{15}$ (B).\\footnote[1]{$\\Delta m_{15} (B)$ is the magnitude difference in $B$-band between maximum brightness and 15 days later. } In the present study, we adopt an $R_{V}$ $= 1.72$ as derived from F10 (i.e. Calibration 7 of Table~9). Moreover, after application of the light-curve correction, Hubble diagram residuals at the level of $\\sim 0.15$ mag exist for individual SNe~Ia (e.g., Phillips et al.\\ 1999; Prieto et al.\\ 2006; Jha et al.\\ 2007; Hicken et al.\\ 2009ab). This is one of the issues that presently limit the precision in using SNe~Ia to constrain the value of the equation-of-state parameter of the dark energy (e.g., Hicken et al.\\ 2009b; see also Wood-Vasey et al. 2007, Kessler et al. 2009 for the current status of the precision in SN~Ia cosmology). Several suggestions have been made for a secondary parameter that may provide a more accurate luminosity calibration\\footnote[2]{Indeed, this is the `third' parameter, since light curve fitting methods usually use two parameters, i.e., the light curve shape and color. In this paper, we simply call the additional parameter the `second' parameter following the convention.} (or on parameters already including the effect of the second parameter). Suggestions include (i) metallicity (Gallagher et al.\\ 2005; Timmes et al. 2003; Mazzali \\& Podsiadlowski 2006; H\\\"oflich et al. 2010, but see also Howell et al.\\ 2009; Neill et al.\\ 2009; Yasuda \\& Fukugita\\ 2010), (ii) high-velocity spectral features (Wang et al.\\ 2009b), (iii) spectral flux ratios (Bailey et al.\\ 2009; Yu et al.\\ 2009), and (iv) the mass and/or the morphological type of the host galaxy (Kelly et al.\\ 2010; Lampeitl et al. 2010; Sullivan et al. 2010). An interesting possibility for the origin of the diverse properties of SNe Ia was recently suggested by Kasen et al. (2009) theoretically and by Maeda et al. (2010ab, hereafter M10a and M10b) observationally, namely an asymmetry in the SN explosion combined with the observer viewing angle. In particular, M10a identified potential signatures of asymmetry in a number of SNe Ia, based on the observed wavelength shift of late-time emission lines (see \\S 2 for more details). M10ab showed that the required configuration is qualitatively consistent with the expectation from a deflagration-to-detonation transition scenario\\footnote[3]{In the deflagration-to-detonation transition scenario, the thermonuclear sparks first trigger the deflagration flames which travel subsonically, and then the flames subsequently turn into a supersonic detonation flame (Khokhlov 1991). } if the first thermonuclear sparks are ignited offset from the centre of the progenitor white dwarf. M10b suggested that the viewing angle effect is a probable origin of the spectral evolution diversity of SNe Ia. Different SNe~Ia show different velocity gradients ($\\dot v_{\\rm Si}$), defined as the speed of the decrease in the Si II absorption velocity after maximum brightness (Benetti et al. 2005; see also Branch et al. 1988). SNe are divided into high-velocity-gradient (HVG) ($\\dot v_{\\rm Si} > 70$ km s$^{-1}$) and low-velocity-gradient (LVG) objects ($\\dot v_{\\rm Si} < 70$ km s$^{-1}$). M10b argued that different velocity gradients are a consequence of different viewing directions from which the SN is observed. It has been indicated that LVG and HVG SNe may show different properties in their intrinsic colors (e.g., Pignata et al.\\ 2008) and that their luminosities may have to be calibrated in a different manner (Wang et al. 2009b). Here we revisit this color issue in the context of our new interpretation of LVG and HVG SNe. In this paper, we explore whether the late-time emission line shift, and thereby the observer viewing angle on an asymmetric explosion, is related to the intrinsic color and the luminosity residuals of SNe~Ia after application of the Phillips relation. We find a correlation between the color at maximum brightness and the nebular emission line shift. We also investigate a possible relation between the luminosity residuals and the nebular line shifts, but since our sample is small the significance is not overwhelming. We then investigate the ramification of the viewing angle on the luminosity and color calibrations with the help of multi-dimensional radiation transfer calculations, and find that the predicted effect is qualitatively consistent with the trends seen in the data. The paper is organized as follows. In \\S 2, we summarize the findings of M10a regarding the asymmetry in SNe~Ia, which are then used throughout the present paper. In \\S 3, we present details of the sample of nearby SNe~Ia considered in this study. In \\S 4, we discuss how the viewing angle is related to the intrinsic color of SNe~Ia. In \\S 5, we discuss the procedures to estimate the intrinsic absolute magnitude and subsequent residuals. In \\S 6, we compare the late-time emission line shifts with the luminosity residuals. In \\S 7, we investigate the effect of the viewing angle on the peak brightness and color by simulating light curves for kinematic off-centre toy models. In \\S 8 the paper is closed with conclusions, discussion and future perspectives. ", "conclusions": "In this paper, we have investigated how the explosion geometry and viewing angle can influence the color and peak brightness of SNe Ia, and thereby lead to the residuals that remain in the peak magnitudes after application of the light-curve correction. We have used the wavelength shift of late-time emission lines to derive the viewing angle, and then compared this quantity with the color and the luminosity residual for a sample of 20 SNe Ia. \\subsection{Intrinsic Color Variation} We have found a correlation between the color at maximum and the wavelength shift seen in late-time emission lines. Using a sub-sample of 11 SNe which likely suffer from insignificant host reddening, selected based on the morphological type of the host galaxy and the position of the SN within the host, we have found that SNe which show a blueshift/redshift in nebular emission lines (i.e., viewed from the offset/anti-offset direction in the off-centre ignition scenario) are bluer/redder than expected from the color-$\\Delta m_{15} (B)$ relation. This indicates that the previously suggested color difference between the `high-velocity-gradient' (HVG) SNe and `low-velocity-gradient' (LVG) SNe is connected to the viewing direction to the observer. In a next step, we expanded our sample, including objects with potentially larger host reddening. Except for 3 very red SNe [($B_{\\rm max} - V_{\\rm max}) > 0.2$ mag: SNe 1998bu, 2002bo and 2006X], all other objects agreed well with the color vs. $v_{\\rm neb}$ relation derived from the ``low-extinction'' sample. This indicates that a significant part of the color excess previously attributed to host extinction is actually due to intrinsic color variations. This raises the question whether the often preferred value of $R_{V} \\lsim 2$ (much smaller than the typical Galactic value of $\\sim 3.1$) really reflects different properties of interstellar/circumstellar dust around SNe Ia. F10 argued that the optical color minus NIR color relation can be reproduced even with $R_{V} \\sim 3.2$ once heavily reddened SNe are omitted. On the other hand, they required $R_{V} \\sim 1 - 2$ to minimize the dispersion in the standardized luminosity calibration. They pointed out that an intrinsic color variation which does not correlate with the decline rate, but does correlate with the luminosity, may solve this apparent discrepancy. Indeed, the intrinsic variation related to the viewing angle that we have found in this paper does not correlate with the decline rate. If we exclude SNe which are clearly heavily reddened, we see that $R_{V}$ as large as the Galactic value could be acceptable. This issue requires more careful and systematic analysis based on a larger sample. Also, we note that even if this works for the majority of SNe Ia, some heavily reddened SNe appear to require $R_V \\sim 1 - 2$, as highlighted by SN 2006X (F10). A toy model constructed with the constraints so that it explains the late-time spectral variation (M10a) and its relation to the early-phase spectral diversity (M10b) predicts the bluer/redder color for smaller/larger (blueshift/redshift) $v_{\\rm neb}$, as is consistent with the observational data. For our model sequence, the predicted trend is not sensitive to $M$($^{56}$Ni) (or $\\Delta m_{15} (B)$). \\subsection{Second Parameter in the Luminosity Calibration?} We have also investigated a relation between the SN Ia luminosity residual after calibration with the Phillips relation, and the wavelength shift seen in the late-time nebular emission lines. For our small sample, the correlation is not strong, and we regard this result as tentative. Keeping this caveat in mind, we see a tendency that SNe Ia with a blueshift/redshift in the nebular lines show a higher/lower peak luminosity than expected by the Phillips relation. There is an average difference of $\\sim 0.25$ mag in the peak magnitudes between SNe Ia showing a blueshift and those showing a redshift. Calculating the light curves based on the geometry derived by M10a for the relatively faint SN Ia 2003hv, we have found that SNe Ia viewed from the offset direction should have a peak brightness larger than the mean, and those viewed from the opposite direction should be fainter. The difference between these two extreme observer directions is $\\sim 0.7$ mag in our fiducial model A0.3.\\footnote[13]{Including the effect that the light curve shape is different for different viewing directions.} This behavior is consistent with the observational data. Also, the averaged difference between SNe showing blueshifts and redshifts is $\\sim 0.4$ mag in this model, enough to explain the observed value ($\\sim 0.25$ mag). The comparison between SNe Ia with normal/large peak luminosity and the models is less straightforward, since the geometry of such explosions has not been directly constructed from observations. Assuming that the degree of the offset is similar to that of SN 2003hv, we expect that the viewing angle effect is less pronounced for brighter SNe Ia. This is consistent with the behavior we have found in the data. Further variation is expected since the degree of the offset could be different for SNe with different $\\Delta m_{15} (B)$ (e.g., Kasen et al. 2009; M10c). We therefore suggest that the residual could be explained by a combination of the configuration [e.g., the relative contribution between the ECAP/HD and LD zones, which may well be expressed by one parameter, i.e., $\\Delta m_{15} (B)$] and the viewing angle. Thus, we do not expect a single straight line to give a perfect fit to the residual vs. late-time velocity plot. This may be one reason why the correlation in Fig.~8 is not too strong. \\subsection{Future Perspectives and Cosmological Applications} Our results could be further tested by polarization measurements. SNe Ia generally show small continuum polarization around maximum brightness, indicating that they are more or less spherical (e.g., Wang et al.\\ 1996). However, early-phase polarization measurements mainly probe a region near the surface of the SN Ia ejecta, which we also suggest to be nearly spherical. Therefore, the small polarization is likely not a strong argument against our present interpretation. The low continuum polarization is sometimes translated into a small deviation of the photosphere from spherical symmetry, and hence a small dependence of the brightness on the viewing angle (e.g., H\\\"oflich et al.\\ 2010). However, this statement depends strongly on the assumed geometry. For example, a continuum polarization of $\\lsim 0.3\\%$, as is usually found in SNe Ia, implies an axis ratio of less than 10\\% for an ellipsoidal photosphere (H\\\"oflich\\ 1991). However, for a one-sided distribution of $^{56}$Ni, as suggested in the present work, the expected continuum polarization is generally much smaller than for an ellipsoid (Kasen \\& Plewa\\ 2007) despite a larger expected variation in the angle-dependent brightness than for the ellipsoidal case (e.g., Sim et al.\\ 2007; Kasen et al.\\ 2009; this work). There is a correlation between the velocity gradient and the Si II line polarization (Leonard et al. 2005; Chornock \\& Filippenko 2008; Maund et al. 2010). This may indicate that the viewing direction is indeed controlling the line polarization level, if the velocity gradient is determined by the viewing angle effect (M10b). Further study on the polarization of SNe Ia both in observations (e.g., Wang et al.\\ 2007) and in theory (e.g., H\\\"oflich\\ 1991; Kasen \\& Plewa\\ 2007) should provide a good test for the geometry of SNe Ia, and thus for the results in the present work. A problem in our analysis, especially for the residual issue, is the small sample size. The uncertainty in the distance measurement is also critical in our investigation of the viewing-angle effect on the luminosity residual. We suggest the following strategies to decrease the uncertainty in the distance estimate: (1) obtain late-time spectra for SNe~Ia at redshift $z \\gsim 0.02$. At $z \\sim 0.02$, the $V$-band magnitude of typical SNe Ia at $\\sim 200$ days after maximum brightness is $\\sim 21 - 22$ mag (depending on the extinction). Spectroscopy is thus possible with $6 - 8$m class telescopes, and (2) obtain comprehensive photometry during both the peak and tail phases. The residual arising from the viewing angle can in principle be seen in the peak-to-tail luminosity ratio, since the effect of the viewing angle vanishes at late-phases (e.g., Maeda et al.\\ 2006b). This may provide a distance-independent measurement of the residuals caused by the viewing angle effect. Such observations are less demanding than spectroscopy, and thus can reach to larger redshift. One of the current limitations in SN Ia cosmology is the fact that a dispersion at the level of $\\sim 0.15$ mag remains after the standardization of the peak luminosity with existing relations, and that the physical origin of the residual has not been identified. The viewing angle effect may account for part of the dispersion of the SN Ia luminosity calibration, although the presently available sample does not allow us to quantify how much improvement can be achieved by taking this effect into account. However, the effect of the random viewing angle enters in the SN Ia luminosity calibration as a source of a statistic error. Thus, increasing the number of SNe Ia for cosmology should effectively reduce this error in estimating cosmological parameters. In addition, it may be worthwhile looking into the frequency distribution of the residuals. Although we expect that the statistical error is introduced by random viewing angles, there is no reason to expect it to obey a Gaussian distribution. Investigating a non-Gaussian component in the scatter of the Hubble diagram may provide additional insight. For example, a larger SN Ia sample may allow us to estimate the non-Gaussian contribution to the statistical error, which will then provide a quantitative estimate of the viewing angle effect independent of uncertainties in the distance measurements for individual SNe. The result of the present work may shed light on how to develop a more accurate SN Ia cosmology than currently employed. The correlation between the color near maximum and $v_{\\rm neb}$ enables us to discriminate the intrinsic color and the host extinction when late-time spectra are available. This opens up a possibility of studying the intrinsic color and the host extinction, including $R_{V}$, in detail for nearby SNe up to a redshift of $\\sim 0.02$ for which late-time spectroscopy is possible. This will hopefully provide information about how to distinguish the intrinsic color and host extinction, the information applicable to high-redshift SNe. The relation between the velocity gradient and the nebular emission line shift (M10b) suggests that one could use the velocity gradient (accessible to higher redshift SNe) instead of the latter, in the color and luminosity calibrations (see also Foley \\& Kasen 2010). Indeed, it has been argued that LVG and HVG SNe should be treated differently in the luminosity calibration (e.g., Wang et al. 2009b). Further investigating relations among these observable (e.g., Fig.~5), as well as finding other observables which correlate with $v_{\\rm neb}$, could be useful in improving the color and luminosity estimates also for high-$z$ SNe~Ia." }, "1101/1101.4809_arXiv.txt": { "abstract": "We present a covariant formalism for general multi-field system which enables us to obtain higher order action of cosmological perturbations easily and systematically. The effects of the field space geometry, described by the Riemann curvature tensor of the field space, are naturally incorporated. We explicitly calculate up to the cubic order action which is necessary to estimate non-Gaussianity and present those geometric terms which have not yet known before. ", "introduction": "\\label{sec:intro} Inflation~\\cite{inflation} is currently the leading candidate to lay down the necessary initial conditions for the successful hot big bang evolution of the universe~\\cite{book}. The most recent observations from the cosmic microwave background (CMB) are consistent with the predictions of the inflationary paradigm~\\cite{Komatsu:2010fb}: the universe is homogeneous and isotropic with vanishing spatial curvature, and the primordial scalar perturbation is dominantly adiabatic and follows almost perfect Gaussian statistics with a nearly scale invariant power spectrum. Thus, any small deviation from these predictions would provide crucial information for us to distinguish different models of inflation. Especially, the non-linearities in the primordial perturbation have received an extensive interest nowadays in the light of upcoming precise cosmological observations. For example, while the current bound on the non-linear parameter $\\fnl$~\\cite{Komatsu:2001rj} is constrained to be $|\\fnl| \\lesssim \\mathcal{O}(100)$ from the Wilkinson Microwave Anisotropy Probe observation on the CMB~\\cite{Komatsu:2010fb}, the Planck satellite can probe with better precision to detect $|\\fnl| = \\mathcal{O}(5)$~\\cite{:2006uk}. The sensitivity may be even further improved from the observations on large scale structure~\\cite{LSSfNL}. The absence of the relevant scalar field which can support inflation in the standard model (SM) of particle physics\\footnote{ It was recently suggested that the SM Higgs field can play the role of the inflaton provided that it is non-minimally coupled to gravity~\\cite{Bezrukov:2007ep}. However the unitarity of the simplest Higgs inflation appears to be controversial. See e.g. Refs.~\\cite{unitX} and \\cite{unitO} and references therein for different points of view on this issue. } demands that inflation be described in the context of the theories beyond the SM. Typically there are plenty of scalar fields which can contribute to the inflationary dynamics~\\cite{Lyth:1998xn}. Further, in multi-field system we can obtain interesting observational signatures which deviate from the predictions of the single field models of inflation and can be detected in near future, such as isocurvature perturbation~\\cite{Choi:2008et} or non-Gaussianity~\\cite{nGreviews}. Thus we have both theoretical and phenomenological motivations to develop a complete formulation of general multi-field inflation. An important point in multi-field system is that in the field space which generally has non-trivial field space metric, the scalar fields play the role of the coordinate. Naturally, as we do in general relativity, it is preferable to formulate the dynamics in the field space in the coordinate independent manner. That is, we need a covariant formulation of multi-field inflation which allows us to describe the inflationary dynamics with arbitrary field space. However, most studies on multi-field inflation, especially regarding non-linear perturbations, are based on trivial field space~\\cite{Seery:2005gb} or non-covariant description~\\cite{Langlois:2008wt,Langlois:2008qf,Arroja:2008yy}. The existing studies with covariant approach to general field space metric are mostly on linear perturbation theory~\\cite{linearcov}. In this note, we develop a fully covariant formulation of non-linear perturbations in general multi-field inflation. Along with the covariance for general field space, it allows us to obtain arbitrary higher order action of cosmological perturbations easily and systematically. We consider the matter Lagrangian which is a generic function of the field space metric $G_{IJ}$ with $I$ and $J$ being generic field space indices, kinetic function $\\partial^\\mu\\phi^I\\partial_\\mu\\phi^J$ and the fields~\\cite{kinflation}. This form includes not only the matter Lagrangian with the standard canonical kinetic term but also more generic ones motivated from high energy theories, such as the Dirac-Born-Infeld (DBI) type~\\cite{DBI}. This note is outlined as follows. In Section~\\ref{sec:mapping} we set up the geodesic equation to describe the field fluctuation around the background trajectory. In Section~\\ref{sec:matterL}, we consider pure matter Lagrangian and present a covariant formulation to describe the field fluctuations up to arbitrary order. The extension to include gravity follows in Section~\\ref{sec:gravity} and we explicitly compute the perturbed action up to cubic order. We also discuss the genuine multi-field effects briefly. We conclude in Section~\\ref{sec:conclusions}. Technical details to compare with the previously known non-covariant description are presented in the Appendix. ", "conclusions": "\\label{sec:conclusions} In this note, we have studied a covariant formulation of general multi-field inflation. Starting from the geodesic equation parametrized by $\\lambda$ which connects a point on the background trajectory to the corresponding point with field perturbations, we have found the non-linear relation between the real physical field fluctuation $\\delta\\phi^I$ and the vector $Q^I$ living on the tangent space. Using this relation, we have expanded the general matter Lagrangian $P(G_{IJ},X^{IJ},\\phi^I)$ in terms of $\\lambda$ up to cubic order in $Q^I$. The resulting expression is fully covariant with the Riemann curvature tensor $R_{IJKL}$ describing the geometry of the field space. Including gravity, we have chosen the flat gauge where metric perturbations are given by the solutions of the constraint equations in terms of $Q^I$. For an explicit calculation up to cubic order, which is necessary to find the leading contribution to the bispectrum of the curvature perturbation, we need only the linear solutions of the metric perturbation which could be found from the second order action. With these solutions, we have explicitly computed the cubic order action in a fully covariant manner. Although we have presented up to cubic order action, our formulation can be straightforwardly extended to find arbitrary higher order action. We have also discussed briefly the genuine effects in multi-field inflation generated by the isocurvature perturbations. \\subsection*{Acknowledgement} JG thanks Ana Ach\\'ucarro for important conversations. JG is grateful to the Yukawa Institute for Theoretical Physics at Kyoto University for hospitality during the long-term workshop ``Gravity and Cosmology 2010 (GC2010)'' (YITP-T-10-01) and the YKIS symposium ``Cosmology -- The Next Generation --'' (YKIS2010), where this work was initiated, and the 20th Workshop on General Relativity and Gravitation in Japan (YITP-W-10-10) where this work was under progress. This work was supported in part by a Korean-CERN fellowship, the Japanese Society for Promotion of Science Grants N. 21244033, the Global COE Program ``The Next Generation of Physics, Spun from Universality and Emergence'', and the Grant-in-Aid for Scientific Research on Innovative Areas (Ns. 21111006 and 22111507) from the MEXT. \\appendix \\renewcommand{\\theequation}{\\thesection.\\arabic{equation}} \\setcounter{equation}{0}" }, "1101/1101.1227_arXiv.txt": { "abstract": "We show that the recently observed elemental abundance pattern of the carbon-rich metal-poor Damped Lyman $\\alpha$ (DLA) system is in excellent agreement with the nucleosynthesis yields of faint core-collapse supernovae of primordial stars. The observed abundance pattern is not consistent with the nucleosynthesis yields of pair-instability supernovae. The DLA abundance pattern is very similar to that of carbon-rich extremely metal-poor (EMP) stars, and the contributions from low-mass stars and/or binary effects should be very small in DLAs. This suggests that chemical enrichment by the first stars in the first galaxies is driven by core-collapse supernovae from $\\sim 20-50 M_\\odot$ stars, and also supports the supernova scenario as the enrichment source of EMP stars in the Milky Way Galaxy. ", "introduction": "At the end of the dark age of the Universe, the cosmic dawn was heralded by the birth of the first stars and galaxies. The nature of these first objects is still far from being well understood. From the theory of star formation from primordial gas, the first stars are believed to be very massive, with masses of the order of $100M_\\odot$, given the limited cooling of molecular hydrogen \\citep[e.g.,][]{bro04}. This depends on fragmentation in a cosmological minihalo, ionization prior to the onset of gravitational collapse, and the accretion rate from the cloud envelope \\citep{mck04,ohk09}, and it seems possible to form lower-mass stars ($\\sim10-40M_\\odot$) from primordial gas in recent numerical simulations \\citep{yos08,sta09}. The primordial stars with initial masses of $\\sim 140-270M_\\odot$ explode as pair instability supernovae (PISNe) \\citep{barkat,ume02,heg05}. In terms of chemical abundances, however, no observational signature for the existence of PISNe has been detected. Chemical enrichment from the first generation of stars has mainly been studied with extremely metal-poor (EMP) stars \\citep[e.g.,][]{bee05}. $10-25$\\% of EMP stars with [Fe/H] $\\ltsim -2$ \\citep{aok10} show carbon enhancement relative to iron ([C/Fe] $\\gtsim1$); this is also the case for the three hyper/ultra metal-poor stars known, with [Fe/H] $<-4.5$. Such stars are collectively known as carbon-enhanced metal-poor (CEMP) stars, and are further subdivided as CEMP-s and CEMP-no stars depending on whether they exhibit or not an enhancement in the abundances of slow neutron capture elements such as Barium. The abundance patterns of EMP stars are explained with one of, or the combination of, four enrichment sources: (1) core-collapse supernovae \\citep{ume03,iwa05,tom07}, (2) rotating massive stars \\citep{mey06}, (3) asymptotic giant blanch (AGB) stars in binary systems \\citep{sud04}, or (4) interstellar accretion \\citep{yos81,ibe83}. These enrichment sources have distinct signatures, so that it should be possible to distinguish between them by comparing the observed elemental abundances with theoretical calculations of nucleosynthetic yields. The high [C/Fe] can be explained with the mass transfer from AGB stars in binary systems, the mass loss from rotating massive stars, and a single core-collapse supernova forming a black hole. The Ba enhancement can be explained by AGB stars and possibly by rotating massive stars \\citep{pig08}. However, iron-peak elements have to come from supernovae. The effect of interstellar accretion has not been studied with hydrodynamical simulations, but is observationally estimated to be negligible \\citep{fre09}. The binarity of CEMP stars has been studied by monitoring radial velocity variations. A signature of binarity is seen statistically for possibly all CEMP-s stars \\citep{luc05}, but the binary fraction of CEMP-no stars seems to be much lower \\citep{aok10}. In particular, the binarity is not seen for the three stars with [Fe/H] $<-4.5$. \\citet{ume03} were the first to show that the observed abundance pattern from C to Zn can be well reproduced with the enrichment from a single core-collapse supernova that leaves behind a relatively massive black hole, under the assumption of inhomogeneous chemical enrichment \\citep{aud95}. \\citet{iwa05} showed that the N abundance can be as large as observed with enhanced mixing between H and He layers during the hydrostatic stellar evolution. The observations of very metal-poor Damped Lyman $\\alpha$ (DLA) systems have opened a new window to study the chemical enrichment of the Universe by the first generations of stars. DLAs are quasar absorbers defined by their high column density of neutral hydrogen, $\\log N$({H}\\,{\\sc i})/cm$^{-2} \\ge 20.3$. They appear to sample a range of galaxy types, from the extended H\\,{\\sc i} disks of galaxies, to smaller subgalactic size haloes, as well as smaller H\\,{\\sc i} clouds within larger galaxies \\citep{wol05}. Large scale surveys, such as the Sloan Digital Sky Survey, have increased ten-fold the number of known DLAs, which now number in excess of $\\sim 1000$ \\citep{not09,pro09}. Follow-up high resolution spectroscopy of the most metal-poor DLAs is of particular interest, since the gas they trace may have been enriched by very few generations of stars \\citep{pet08,pen10}. Moreover, measuring elemental abundances in DLAs is straightforward; the only potential complications are line saturation and dust depletion and both effects are of much reduced importance for metallicities $Z \\ltsim 1/100 Z_\\odot$. In addition, the chemical evolution of such systems is rather simple, whereas in more chemically evolved systems, there are uncertainties in the star formation history, gas inflow and outflow, large contributions from AGB stars and Type Ia supernovae, with the result that the signatures of the first stars can be easily washed out. Thus, if the most metal-poor DLAs hold the key to unravelling the chemical enrichment from the first stars, it is of great interest to compare their observed abundance patterns with the nucleosynthesis yields of metal-free stars. \\citet{coo10b} recently reported such a DLA with [Fe/H] $\\simeq -3$, which exhibits a strong carbon enhancement relative to all other available elements, including [C/Fe] $\\simeq +1.53$. This reminds us of the CEMP stars in the solar neighborhood. In this letter, we compare the elemental abundance pattern of the C-rich DLA with the nucleosynthesis yields of both core-collapse and pair-instability supernovae (\\S2). In \\S3 we give a more general discussion of the chemical enrichment of the Universe by the first generation of stars. We summarize our main conclusions in \\S4. ", "conclusions": "In the early stages of chemical enrichment, the interstellar medium is supposed to be highly inhomogeneous, so that the properties of the first objects can be directly extracted from the comparison between the observed elemental abundances and nucleosynthesis yields. We have shown that the observed abundance pattern of the very metal-poor C-rich DLA is in excellent agreement with the nucleosynthesis yields of a primordial star that explodes as a faint core-collapse supernova owing to the efficient mixing and fallback. The nucleosynthesis yields of PISNe are not consistent with the observation. The contribution from rotating massive stars seems to be small because of the lack of N enhancement. The contribution from AGB stars should be very small because of the N abundance and of the enrichment timescale. Since the DLA abundances reflect the chemical enrichment in gas-phase, the binary or accretion scenarios of the EMP stars do not work. Thus, we conclude that enrichment by primordial supernovae is the best solution to explain the abundance pattern of the C-rich DLA. The abundance pattern of the C-rich DLA is similar to those of EMP stars such as the ultra metal-poor star HE0557-4840. Some of EMP stars in dSphs and the Galactic outer halo also show similar carbon enhancement at [Fe/H] $\\ltsim -3$. Chemical enrichment by the first stars in the first galaxies is likely to be driven by core-collapse supernovae." }, "1101/1101.3708_arXiv.txt": { "abstract": "We studied the formation process of star clusters using high-resolution $N$-body/smoothed particle hydrodynamcs simulations of colliding galaxies. The total number of particles is $1.2\\times10^8$ for our high resolution run. The gravitational softening is $5~{\\rm pc}$ and we allow gas to cool down to $\\sim 10~{\\rm K}$. During the first encounter of the collision, a giant filament consists of cold and dense gas found between the progenitors by shock compression. A vigorous starburst took place in the filament, resulting in the formation of star clusters. The mass of these star clusters ranges from $10^{5-8}~{M_{\\odot}}$. These star clusters formed hierarchically: at first small star clusters formed, and then they merged via gravity, resulting in larger star clusters. ", "introduction": "Merging galaxies contain many young star clusters (e.g., \\cite[Whitmore 2003]{Whitmore2003}). These star clusters could potentially evolve into the present day metal-rich globular clusters and so they are widely accepted to be a good candidate for globular cluster progenitors. Galaxy-galaxy merger is therefore considered as one of the formation channels of globular clusters (\\cite[Schweizer 1987]{Schweizer1987}). There have been a number of numerical studies of merging galaxies. There have been, however, only a few numerical studies of star cluster formation in merging galaxies even though it has been shown that resolving a cloudy/multiphase interstellar medium (ISM) and/or clustered star formation can have important consequences for the formation history of early-type galaxies (\\cite[e.g. Bois et al. 2010]{Bois+2010}). Some of the existing studies adopted sub-grid models of star cluster formation (\\cite[e.g., Bekki \\& Couch 2001; Li et al. 2004]{BekkiCouch2001, Li+2004}), while more recently their formation has been captured directly (e.g. using the sticky particle method in \\cite[Baurnaud et al. 2008]{Baurnaud+2008}). Here we report the result of merger simulations that capture the multiphase nature of the ISM and include realistic models of star formation and feedback. ", "conclusions": "We performed high resolution simulations of merging galaxies that capture the multiphase ISM and the realistic modelings of star formation and of SN feedback. In the simulations, we found that a number of star clusters formed during the starburst at the first encounter. These star clusters grew via hierarchical mergers." }, "1101/1101.1820_arXiv.txt": { "abstract": "{Star formation occurs via fragmentation of molecular clouds, which means that the majority of stars born are members of binary systems. There is growing evidence that planets might form in circumprimary disks of medium-separation ($\\lesssim\\,50\\,\\mathrm{AU}$) binaries. The tidal forces caused by the secondary generally act to distort the originally circular circumprimary disk to an eccentric one. Since the disk eccentricity might play a major role in planet formation, it is of great importance to understand how it evolves.} {We investigate disk eccentricity evolution to reveal its dependence on the physical parameters of the binary system and the protoplanetary disk. To infer the disk eccentricity from high-resolution near-IR spectroscopy, we calculate the fundamental band ($4.7\\,\\mathrm{\\mu m}$) emission lines of the CO molecule emerging from the atmosphere of the eccentric disk.} {We model circumprimary disk evolution under the gravitational perturbation of the orbiting secondary using a 2D grid-based hydrodynamical code, assuming $\\alpha$-type viscosity. The hydrodynamical results are combined with our semianalytical spectral code to calculate the CO molecular line profiles. Our thermal disk model is based on the double-layer disk model approximation. We assume LTE and canonical dust and gas properties for the circumprimary disk.} {We find that the orbital velocity distribution of the gas parcels differs significantly from the circular Keplerian fashion. The line profiles are double-peaked and asymmetric in shape. The magnitude of asymmetry is insensitive to the binary mass ratio, the magnitude of viscosity ($\\alpha$), and the disk mass. In contrast, the disk eccentricity, thus the magnitude of the line profile asymmetry, is influenced significantly by the binary eccentricity and the disk geometrical thickness.} {We demonstrate that the disk eccentricity profile in the planet-forming region can be determined by fitting the high-resolution CO line profile asymmetry using a simple 2D spectral model that accounts for the velocity distortions caused by the disk eccentricity. Thus, with our novel approach the disk eccentricity can be inferred from high-resolution near-IR spectroscopy data acquired prior to the era of high angular resolution optical (ELT) or radio (ALMA, E-VLA) direct-imaging. By determining the disk eccentricity in medium-separation young binaries, we might be able to constrain the planet formation theories.} ", "introduction": "Star formation occurs via fragmentation of molecular clouds causing about 60\\% of stars to be born as a member of a binary system \\citep{DuquennoyMayor1991}. % Large initial specific angular momentum results in circumbinary disk formation around the protobinary, while circumstellar disks (a circumprimary and a circumsecondary) are formed around the protostars for lower initial specific angular momentum \\citep{BateBonnell1997}. \\citet{BonavitaDesidera2007} demonstrated that the overall frequency of giant planets in binaries and single stars does not statistically differ among planets discovered by radial velocity surveys. On the basis of a comprehensive survey for companions of 454 nearby Sun-like stars, \\citet{Raghavanetal2010} revealed that both single and multiple stars are equally likely to harbor planets. Most of the planet-bearing binaries have large separations where the planet formation might be unaltered by the companion's gravitational perturbation. However, stellar multiplicity might play a major role in planet formation in medium-separation ($\\lesssim 50\\,\\mathrm{AU}$) binary systems. Based on Doppler surveys, \\citet{EggenbergerUdry2010} showed that about 17\\% of circumstellar exoplanets are associated with binaries. Among this, five circumstellar exoplanets are known to date in $\\lesssim50\\,\\mathrm{AU}$ separation binary system \\citep{Quelozetal2000,Hatzesetal2003,Zuckeretal2004, Lagrangeetal2006,Chauvinetal2006,Correiaetal2008}. Two exceptional cases are also known, HW Virginis \\citep{Leeetal2009} and CM Draconis \\citep{Deegetal2008}, in which planets have been detected in circumbinary orbits. Thus, planet formation theories, such as core-accretion \\citep{BodenheimerPollack1986,Pollacketal1996} or gravitational instability \\citep{Boss2001}, must be able to explain the formation of planets in both circumprimary and circumbinary disk environments. The circumprimary disk is tidally truncated at $0.35-0.5$ times the binary separation, depending on the binary mass ratio, binary eccentricity, and magnitude of the disk viscosity \\citep{ArtymowiczLubow1994}. Owing to the angular momentum transfer between the disk and the companion \\citep{PapaloizouPringle1977}, the disk is truncated, resulting in a greatly reduced disk lifetime being available for planets to form in $\\lesssim50\\,\\mathrm{AU}$ binaries \\citep{Ciezaetal2009} that is a severe problem for core-accretion scenario. In contrast, the formation of gas giant planets by the relatively rapid-action of the gravitational instability might be induced by the secondary-generated shock waves if the gas cooling time is short. Nevertheless, the disk viscosity can heat the disk sufficiently to suppress the formation of clumps \\citep{Nelson2000}, but with small viscosity, gravitationally unstable clumps can still form \\citep{Boss2006}. The disk gas feels the companion's periodic perturbation leading to strong interaction at the location of the Lindblad resonances. Waves launched at Lindblad resonances carry energy and angular momentum from the binary. The disk experiences changes in its angular momentum where the waves dampen, resulting in the development of an eccentric disk \\citep{Lubow1991}. Several mechanism that may cause wave damping have been proposed, e.g. shocks that could be effective in colder disks, or turbulent disk viscosity acting as a dissipation source and radiative damping. Both SPH \\citep{ArtymowiczLubow1994} and grid-based \\citep{Kleyetal2008} hydrodynamical simulations have confirmed the eccentricity development in binaries assuming spatially constant viscosity for the gas as a source of the wave damping. Since the orbit of bodies (dust particles or pebbles) is perturbed not only by the periodic gravitational potential of the binary but the gas drag as well, the development of disk eccentricity might influence the core-accretion processes. The maximum size of the building blocks of planetesimals is affected by the impact velocity of sub-micron-sized grains in the dust coagulation process. According to Zsom et al. (2010), the aggregate sizes are lower in eccentric disks than in axisymmetric disk environments owing to the increase in the relative velocity between the dust particles. We note, however, that an investigation of the SED slopes of medium-separation T\\,Tauri binaries by \\citet{Pascuccietal2008} showed that the extent of dust processing in the disk surface layer and the degree of dust settling in binary disks do not differ significantly from those in disks around single stars. The planetesimal accretion phase, leading to between km-sized planetesimals and several 100\\,km-sized planetary embryos, should proceed in an environment where the mutual encounter velocities of planetesimals are on the order of planetesimal surface escape velocities. In this environment, the planetary embryos could grow quickly by runaway growth mode \\citep{WetherillStewart1989}. Since the runaway growth mode is sensitive to the encounter velocities, for increased encounter velocities, e.g., due to the stirring up of the planetesimal swarm, runaway growth might be stopped. \\citet{Thebaultetal2006} showed that the impact velocity of different-sized planetesimals tends to increase owing to the interaction between the companion and the gaseous friction in binaries with separations of $10\\leq a_\\mathrm{bin}\\leq50\\,\\mathrm{AU}$. \\citet{Paardekooperetal2008} found that the planetesimal encounter velocities with different sizes could be larger by an order of magnitude in eccentric disks than in the axisymmetric case. Consequently, planetesimal accretion might be inhibited in highly eccentric disks. In accretion disks, the presence of the double-peaked emission lines are the natural consequence of the gas parcels moving in Keplerian orbits around the host star. \\citet{Huang1972} and \\citet{Smak1981} presented this in connection with the emission lines of Be and cataclysmic variable stars. \\citet{HorneMarsh1986} investigated the emerging line profiles in accretion disks. Since the Keplerian angular velocity of gas parcels is highly supersonic in accretion disks, the Doppler shift of the line emitted by individual gas parcels exceeds the local line-profile width. Summing up the line profiles emitted by individual rings of gas parcels, and taking into account the radial dependence of the line surface brightness, the result is the well-known double-peaked broad symmetric line shapes \\citep{HorneMarsh1986}. Azimuthal asymmetries in disk surface brightness (e.g. density perturbations in optically thin lines, or supersonic anisotropic turbulence in saturated lines), will break the line profile symmetry \\citep{Horne1995}. Gas parcels orbiting non-circularly (i.e., in elliptic orbit) might also produce the asymmetric line profiles presented by \\citet{Foulkesetal2004} for cataclysmic variables, and Reg\\'aly et al., (2010) for protoplanetary disks. In this paper, we investigate the eccentricity evolution of a circumprimary disk in a young binary system. We perform an extensive parameter study to reveal the dependence of the disk eccentricity on several parameters, such as binary and disk geometry, and gas viscosity. The hydrodynamical simulations were done in 2D by a grid-based parallel hydrodynamic code FARGO \\citep{Masset2000}. We model the circumprimary disk evolution under the gravitational perturbation of the orbiting secondary assuming $\\alpha$-type viscosity \\citep{ShakuraSunyaev1973}. We calculate the fundamental band ($4.7\\,\\mathrm{\\mu m}$) ro-vibrational emission lines of the molecule $\\mathrm{^{12}C^{16}O}$ emerging from the disk atmosphere, providing a tool to determine the disk eccentricity from high-resolution near-IR spectroscopy by means of line profile distortions. Our thermal disk model is based on the double-layer disk model of \\citet{ChiangGoldreich1997}. Since the velocity distribution of the gas parcels show supersonic deviations from the circular Keplerian one, owing to the eccentric disk state, asymmetric molecular line profiles emerge from the optically thin disk atmosphere. The paper is structured as follows. In the next section, we present our hydrodynamical simulation modeling of the evolution in general of disk eccentricity. The calculation of fundamental-band CO ro-vibrational emission lines emerging from an eccentric circumprimary disk are presented in Sect. 3. In Sect. 4, we present an extensive parameter study to investigate the evolution of the disk eccentricity for a wide range of binary and disk parameters. Section 5 deals with the comparison of our results to other recent simulations, and the observability of eccentric signatures. The paper closes with conclusions. ", "conclusions": "Our study of both the hydrodynamic and eccentricity evolution of the circumprimary disk in medium-separation young binary systems has revealed the following findings: \\begin{enumerate} \\item{The quasi-steady eccentric disk state always develops in circumprimary disks of young medium-separation ($20-40\\,\\mathrm{AU}$) binaries within the average disk lifetime, if the viscosity is between widely accepted values ($0.01\\leq\\alpha\\leq0.1$).} \\item{The CO line profiles are asymmetric ($A_\\mathrm{pp}\\simeq20\\%$) as the average disk eccentricity is $\\bar{e}_\\mathrm{disk}\\simeq0.2$ inside 2-3\\,AU, where the CO is excited by the primary's irradiation.} \\item{The orbital eccentricity of binary systems $e_\\mathrm{bin}\\geq0.2$ or their high/low disk geometrical thickness ($h\\leq0.01$/$h\\geq0.1$) might inhibit the development of the quasi-static disk eccentric state.} \\item{The inner ($R\\leq2-3\\,\\mathrm{AU}$) disk eccentricity profile can be determined by fitting the observed high-resolution near-IR CO line profile asymmetry using a simple 2D spectral model.} \\end{enumerate} \\noindent Consequently, taking into account that the eccentricity of protoplanetary disks might strongly influence planet formation, by measuring it we might further constrain the planet formation theories in medium-separation binaries." }, "1101/1101.3478_arXiv.txt": { "abstract": "We report the detection of the 6.4 keV Iron K$\\alpha$ emission line in two infrared-luminous, massive, star-forming \\bzk~ galaxies at $z=2.578$ and $z=2.90$ in the CDF-S. The \\emph{Chandra} 4 Ms spectra of \\bzk4892 and \\bzk8608 show a reflection dominated continuum with strong Iron lines, with rest-frame equivalent widths EW$\\sim$2.3 keV and 1.2 keV, respectively, demonstrating Compton thick obscuration of the central AGN. For \\bzk8608 the line identification closely matches the existing photometric redshift derived from the stellar emission. We use the observed luminosities of the Iron K$\\alpha$ line, of the rest-frame mid-IR continuum and of the UV rest-frame narrow emission lines to infer intrinsic $L_{\\rm 2-10~keV}\\gtrsim 10^{44}$~erg~s$^{-1}$, about 1.0--2.5 dex larger than the observed ones, hence confirming the presence of an absorber with $N_{\\rm H} > 10^{24}$~cm$^{-2}$. The two \\bzk~ galaxies have stellar masses of $5\\times10^{10}$~M$_\\odot$ and, based on VLA 1.4~GHz and submm 870$\\mu$m observations, they appear to host vigorous starburst activity with $SFR\\sim 300$-700~M$_\\odot$~yr$^{-1}$ that is also optically thick. We estimate that the AGN might also conceivably account for an important fraction of the bolometric far-IR emission of the galaxies. The implied volume density of Compton thick (CT) AGN with $L_{\\rm 2-10~keV}>10^{44}$~erg~s$^{-1}$ is in agreement with predictions from X-ray background synthesis models. These sources provide one of the first clearcut observations of the long-sought phase of simultaneous, heavily obscured quasar and star formation activity, predicted by models of massive galaxy evolution at high redshifts. ", "introduction": "Highly obscured, deeply embedded AGN might contribute significantly to the total accretion power in the Universe (Marconi et al. 2004) and are required to account for the spectrum of the X-ray background (XRB, Gilli et al. 2007). In particular, AGN obscured in the X-ray band by column densities larger than N$_H=10^{24}$ cm$^{-2}$ (Compton thick AGN) represent 20--25\\% of the AGN detected by Integral and SWIFT/BAT in the local universe (Malizia et al. 2009; Burlon et al. 2010). The identification of highly obscured AGN is particularly challenging since, because of the high column densities, the vast majority of their X-ray emission below rest frame 10~keV is absorbed, so that they are largely missed even in the deepest hard X-ray surveys available today. At high redshift the quest for such objects is mainly pursued by indirect evidence, by selecting galaxies with high ratios of mid-infrared to optical, or mid-infrared to X-ray fluxes, and by deriving average X-ray properties by stacking techniques (e.g., Daddi et al. 2007b; D07 hereafter; Fiore et al. 2008, 2009; Alexander et al. 2005; 2008; Treister et al. 2010, Georgantopoulos et al. 2009, Eckart et al. 2010, Donley et al. 2010). These studies have suggested that a substantial fraction of massive galaxies at $z>1$ host highly obscured, intrinsically luminous AGN. Only a few tens of secure Compton thick AGN are known so far in the local Universe (Comastri et a. 2004, Dalla Ceca et al. 2008 and references therein), and only for a handful of them a quasar-like intrinsic hard X-ray luminosity has been inferred (L$_{2-10~keV}> 10^{44}$erg~s$^{-1}$; Braito et al. 2004, Piconcelli et al. 2010). Their X-ray spectra are characterized by the presence of a strong Iron K$\\alpha$ 6.4 keV fluorescent emission line, with large equivalent width EW~$\\gtrsim1$~keV, on a flat reflection-dominated continuum (Matt et al. 2000). Not much evidence currently exists in the distant Universe of AGNs showing prominent, high EW Iron K$\\alpha$ line: an IRAS selected hyperluminous galaxy at $z=0.93$ (Iwasawa et al. 2005) and two possible detections in the CDFS reported by Norman et al. (2002; a type 2 QSO at $z=3.7$) and by Tozzi et al. (2006; a galaxy at $z=1.53$); both confirmed by Comastri et al. (2011). In the following we present the significant \\emph{Chandra} detection of the Iron K$\\alpha$ emission line in two massive \\bzk\\ galaxies at redshift z$=2.578$ and z$=2.90\\pm0.10$ in GOODS-South, using the recently acquired 4~Ms dataset. We used for this purpose the entire set of 52 observations for a total exposure time of 4 Ms, available in the Chandra Data Archive (http://cxc.harvard.edu/cda/Contrib/CDFS.html). We adopt a $\\Lambda$CDM cosmology ($H_0=70$ km s$^{-1}$ Mpc$^{-1}$; $\\Omega_M$=0.3; $\\Omega_{\\Lambda}=0.7$) and a Chabrier stellar initial mass function (IMF). \\begin{figure*}[t] \\centering \\includegraphics[angle=-90,scale=.3]{4892_spe_bgd_15.eps} \\includegraphics[angle=-90,scale=.3]{4892_euff.eps}\\\\ \\includegraphics[angle=-90,scale=.3]{8608_spe_bgd.eps} \\includegraphics[angle=-90,scale=.3]{8608_eeuff.eps} \\caption{The spectra of \\bzk4892 (upper panels) and \\bzk8608 (lower panels), fitted with a pure reflection model plus an Iron K$\\alpha$ 6.4~keV line. The spectra are shown both in observed counts units (left panels) and in physical units (keV$^2$; right panels). Each bin corresponds to a 2$\\sigma$ measurement. The red dotted line represents the average background extracted from the whole field, excluding sources (Fiore et al. 2011, in preparation).} \\end{figure*} \\begin{figure*} \\centering \\includegraphics[scale=0.7]{test_both2.eps} \\caption{Upper panels: \\emph{Chandra} images of \\bzk4892 and \\bzk8608 in three energy bands. The middle images are centered at the energy of the Iron K line. The circles show the extraction regions. Lower panels: confidence levels (1, 2 and 3 $\\sigma$) obtained from the fit of the Iron line. The vertical blue lines show the optical spectroscopic/photometric redshift constraints.} \\label{fig:Xzp} \\end{figure*} ", "conclusions": "\\subsection{Intrinsic luminosities and obscuration} The observed (obscured) hard X-ray luminosities are $L_{\\rm 2-10~ keV}=10^{42.5}$ and $10^{42.8}$~erg~s$^{-1}$ for \\bzk4892 and \\bzk8608, respectively. Given the large $N_{\\rm H}$ implied by the detection of high EW Iron K$\\alpha$ lines, the X-ray emission is optically thick and we cannot directly estimate the intrinsic $L_{\\rm 2-10~ keV}$ using the X-ray data (e.g., La Massa et al. 2011). In order to estimate the intrinsic luminosities of the AGN we used a variety of independent methods. Following Iwasawa et al. (2005), a lower limit to the X-ray luminosity can be obtained in the expectation that in typical conditions the luminosity of the Iron K$\\alpha$ line is at most 3\\% of the 2-10 keV luminosity. For \\bzk4892 L$_{K\\alpha}=1.25\\times10^{42}$~erg s$^{-1}$ would imply a intrinsic 2-10 keV luminosity of 4$\\times10^{43}$~erg~s$^{-1}$. The relation L$_{K\\alpha}$/L$_{\\rm 2-10~keV}$ of Levenson et al. (2006) implies an higher luminosity, L$_{2-10~keV}\\sim10^{44.8}$~erg~s$^{-1}$. An alternative estimate is obtained from the 6.0\\micron~ luminosity. The 24\\micron~ flux of 591$\\mu$Jy implies L$_{6.0\\micron}=10^{45.5}$~erg~s$^{-1}$(the mid-IR SED of Mrk~231 is used for the small K-correction). While 24$\\mu$m at $z=2.578$ might be contaminated by PAH features, we find similar results if we use the 16$\\mu$m flux of 209$\\mu$Jy (Teplitz et al. 2010) for which no contamination is expected (4.5$\\mu$m rest-frame). This would suggest L$_{\\rm 2-10~ keV}=10^{45\\pm0.5}$~erg~s$^{-1}$ using the Lutz et al. (2004) correlation and a lower L$_{\\rm 2-10~ keV}=10^{44.4\\pm0.5}$~erg~s$^{-1}$ using Lanzuisi et al. (2009). A third estimate can be derived from the luminosity of the UV-optical emission lines. The VLT/FORS1-2 optical spectrum of \\bzk4892 shows prominent Ly$\\alpha$, CIV, HeII and CIII] narrow emission lines (Szokoly et al. 2004, Vanzella et al. 2006). Using the observed ratios between the line and the intrinsic 2-10~keV emission from Mulchaey et al. (1994), we infer $L_{\\rm 2-10~keV} = 10^{44.5\\pm0.5}$~erg~s$^{-1}$. Using the Netzer et al. (2006) conversion for UV-optical line luminosities would give larger X-ray luminosities by $\\sim$0.5 dex, or L$_{\\rm 2-10~ keV} = 10^{45\\pm0.5}$~erg~s$^{-1}$. All in all, it appears that the 2-10~keV luminosity of \\bzk4892 is in the QSO range, well in excess of $10^{44}$~erg~s$^{-1}$, and more likely of $10^{44.5-45}$~erg~s$^{-1}$. The implied column density is at the level of $N_{\\rm H}\\sim10^{24.5-25}$~cm$^{-2}$. For \\bzk8608, we measure L(K$\\alpha$)$=2.5\\times 10^{42}$~erg~s$^{-1}$. This converts to L$_{\\rm 2-10~ keV}\\sim10^{43.9}$~erg~s$^{-1}$ (Iwasawa et al. 2005), or $\\sim10^{45}$~erg~s$^{-1}$ adopting Levenson et al. (2006). This galaxy is about 10 times fainter in the mid-IR than \\bzk4892, which implies a lower L$_{6.0~\\mu m}$. Using the Lutz et al. (2004) relation, this in turn suggests L$_{\\rm 2-10~ keV}=10^{44}$~erg~s$^{-1}$. We cannot obtain an independent estimate from the UV-optical emission lines given that no lines were detected in the VIMOS spectrum. The intrinsic 2-10~keV luminosity of \\bzk8608 is likely $\\approx10^{44}$~erg~s$^{-1}$. The implied column density is likely in the range $N_{\\rm H}\\sim 1$--$5\\times10^{24}$~cm$^{-2}$, hence this object might be a {\\em borderline} CT AGN. We notice that both sources have estimated intrinsic 2-10 keV luminosities in the QSO regime. \\begin{figure*} \\centering \\includegraphics[angle=0,scale=.65]{FeK_SEDs-3.eps} \\caption{Optical to mid-infrared SEDs of \\bzk4892 (left panel, black symbols) and \\bzk8608 (right panel). The blue continuous line shows the fitted model SED of a pure star forming galaxy. The photometry from the GOODS datasets is taken from D07. HST+ACS color images are also shown as thumbnails. } \\label{fig:SED} \\end{figure*} \\vspace{0.5cm} \\begin{table}[b] \\caption{Galaxy properties} {\\small \\begin{tabular}{lcc} \\hline & \\bzk4892 &\\bzk8608 \\\\ \\hline Ra & 3:32:35.71 & 3:32:20.95 \\\\ Dec & -27:49:16.1 & -27:55:46.3 \\\\ Redshift & 2.578 & 2.90$\\pm0.10$ \\\\ HR & 0.58 & $>$0.54 ($3\\sigma$) \\\\ SFR$_{\\rm UV}$ [M$_{\\odot}$ yr$^{-1}]$ & 70 & 70 \\\\ SFR$_{\\rm Radio}$[M$_{\\odot}$ yr$^{-1}]$ & 1100 & 300 \\\\ L(K$\\alpha$) [erg~s$^{-1}$] & 10$^{42.1}$ & 10$^{42.4}$ \\\\ L$_{\\rm 2-10~keV,~abs}$ [erg~s$^{-1}$] & 10$^{42.5}$ & 10$^{42.8}$\\\\ L$_{\\rm 2-10~keV,~K\\alpha}$ [erg~s$^{-1}$] & 10$^{43.6-44.8} $ & 10$^{43.9-45}$ \\\\ L$_{\\rm 2-10~keV,~6 \\mu m}$ [erg~s$^{-1}$] & 10$^{44.4-45}$ & 10$^{44}$ \\\\ L$_{\\rm 2-10~keV,~UV~ lines}$ [erg~s$^{-1}$] & 10$^{44-45.5}$ & - \\\\ i [AB mag] & 24.75 & 24.72 \\\\ IRAC 3.6 $\\mu$m [AB mag] & 20.81 & 21.29\\\\ IRAC 4.5 $\\mu$m [AB mag] & 20.54 & 21.16 \\\\ IRAC 5.8 $\\mu$m [AB mag] & 20.1 & 20.84 \\\\ IRAC 8.0 $\\mu$m [AB mag] & 19.64 & 20.65 \\\\ MIPS 24 \\micron~ [mJy] & 0.591$\\pm0.07$ & 0.045$\\pm0.003$\\\\ MIPS 70 \\micron~ [mJy] & 3.3$\\pm1.8$ & - \\\\ 850 \\micron~ [mJy]& 3.3$\\pm$1.1 & - \\\\ 1.4 GHz [$\\mu$Jy] & 80$\\pm$14 & 17.5$\\pm6.5$\\\\ \\hline \\hline \\end{tabular} } \\label{tab:data} \\end{table} \\subsection{Bolometric luminosities and star formation rates} The ratio of the bolometric to the 2-10 keV luminosity, L$_{bol}$/L$_{\\rm 2-10~ keV}$, is typically of the order 30-50 for quasars (Marconi et al. 2004). Assuming such a bolometric correction, we estimate L$_{bol}\\sim 10^{46\\pm0.5}$~erg~s$^{-1}$ for \\bzk4892 and $\\sim10^{45.5\\pm0.5}$~erg~s$^{-1}$ for \\bzk8608. These very high luminosities, if expressed in solar units, correspond to L$_{bol}\\sim10^{12.4\\pm0.5}L_\\odot$ and $\\sim10^{11.9\\pm0.5}L_\\odot$. When powered by star formation, similar luminosities require star formation rates (SFR) in the host galaxies at the level of 100-1000~M$_\\odot$~yr$^{-1}$. It is interesting thus to compare these estimates to the inferred SFR for the two \\bzk\\ galaxies, in order to compare the relative AGN and SFR contributions. From the stellar SED (UV to the near-IR rest frame; Fig.~3), we estimate $SFR\\sim70$~M$_\\odot$~yr$^{-1}$ for \\bzk4892, corrected for dust reddening based on the observed UV slope (see D07 for more details). It is quite possible though that also the stellar UV emission might be optically thick. If we interpret its mid-IR emission (591 $\\mu$Jy at 24$\\mu$m) as due to star formation this would imply a whopping SFR$\\sim10^4$~M$_\\odot$~yr$^{-1}$, demonstrating that the mid-IR is likely completely dominated by the AGN (this object is among the most extreme mid-IR excess galaxies in the sample of D07). The SED fit with a star forming galaxy template (Fig.~3) shows an excess emission already at 5.8 \\micron~ (observed frame) probably due to the AGN contribution. The SED in the IRAC bands is indeed showing a steep power law, extending all the way to 24 \\micron. Due to its faint optical counterpart (i$\\sim$25 AB), \\bzk4892 has high mid-infrared to optical flux ratio, F(24)/F(R)$\\sim$2000 and therefore is also classified as a dust-obscured galaxy (DOG, Dey et al. 2008). It is an extreme object also in the Fiore et al. (2008) sample. Alonso-Herrero et al. (2006) reports a 70 \\micron~ flux of $3.3\\pm1.8$ mJy for this galaxy, which is also likely affected by the AGN emission. On the other hand, the galaxy is seen at 1.4 GHz with a flux of 80~$\\mu$Jy in the VLA data of Miller et al. (2008). If due to star formation, using the radio-IR correlation, this would imply SFR$\\sim1100$~M$_\\odot$~yr$^{-1}$. Inspecting the publicly available Apex+LABOCA 870\\micron~ map of GOODS-S (Weiss et al. 2009), we find a 3$\\sigma$ signal at the position of \\bzk4892 of 3.3~mJy, which would also convert into a similar SFR, $\\sim500$~M$_\\odot$~yr$^{-1}$. Even for the most luminous quasars it is generally found that the far-IR emission in the sub-mm bands is due to star formation. If this is the case also for \\bzk4892, our results suggest that the galaxy is witnessing also very powerful star formation activity at the level of 500-1000~M$_\\odot$~yr$^{-1}$. Only 5-10\\% of this is seen directly in the UV, implying that also the UV-emission from stars is optically thick, similarly to local ULIRGs (e.g., Goldader et al. 2002; da Cunha et al. 2010) and to what is expected to be found in major mergers. However, we notice that the radio and 870$\\mu$m derived SFRs would formally imply L$_{bol}\\sim 10^{46-46.5}$~erg~s$^{-1}$, comparable to what inferred for the obscured AGN. Hence, the AGN is likely contributing an important fraction of the total L$_{bol}$. The morphology of the galaxy from HST+ACS imaging is suggestive of a merger, showing two distinct clumps in the UV in the rest frame. When smoothing, after excluding the two bright knots, we detect significant faint low-level emission in the summed $i+z$ band, extended over a diameter of about 0.75\\arcsec (6~kpc). Overall, this is consistent with the size of a massive galaxy at z$\\sim 2$, and we cannot exclude that the two UV knots might be just luminous HII regions inside a big disk galaxy. The photometric information is of lower quality for \\bzk8608 due to its overall faintness. The SFR inferred from the UV is also 70~M$_{\\odot}/yr$. This object is also a mid-IR excess source in D07 sample (by a less extreme factor of 6), with S$_{24} = 45 ~\\mu$Jy ( $\\sim$1/10 of \\bzk4892), although it is not a DOG. The SED (Fig.~3) might also be consistent with pure stellar emission, but a deviation from the star formation template is observed at the $8~\\mu$m band. This galaxy is not detected at 70 \\micron~ and 870 \\micron, suggesting a lower AGN luminosity and/or SFR. Inspecting the VLA 1.4 GHz data (Miller et al. 2008) we detect a faint $3\\sigma$ source at its position with a flux of $17.5\\pm6.5~\\mu$Jy. This corresponds to a luminosity $3\\times10^{12}$ L$_{\\odot}$ and SFR$\\sim300$~M$_\\odot$~yr$^{-1}$. If the radio signal is real and not due to AGN, also in this case the stellar emission appears to be heavily obscured, i.e. optically thick in the UV, as expected for very active starbursts and mergers. The HST+ACS imaging of this galaxy (Fig.~3) is not particularly telling. For both galaxies we derive similar estimates of the stellar masses, at the level of $5\\times10^{10}M_\\odot$. Our results thus support the picture that the luminous, Compton thick AGN that we discovered are hosted by fairly massive galaxies at $z\\sim2.5$-3, that at the same time also host vigorous star formation activity heavily obscured by dust. This is relevant in the AGN-host galaxy co-evolution scenario, in which a phase of rapid, heavily obscured BH growth accompanied by intense obscured star-formation is predicted (Silk \\& Rees 1998, Fabian 1999, Granato et al. 2004). \\subsection{Implications for Compton thick nuclear activity at high redshift} We derive a crude value of the volume density of Compton thick AGN with high EW Iron K$\\alpha$ emission, based on our 2 detections. We conservatively use the full redshift range explored by the \\bzk~ selection (z$=$1.4-3), finding a space density of $3\\times 10^{-6}$ Mpc$^{-3}$. This is consistent with the predictions of the Gilli et al. (2007) X-ray background synthesis model, for Compton thick AGN with L$_{2-10~ keV}>10^{44}$~erg~s$^{-1}$. However, we notice that our sampling of L$_{2-10~keV}>10^{44}$~erg~s$^{-1}$ AGN with Compton thick absorption at $1.410^{43}$~erg~s$^{-1}$. Detecting Iron lines in galaxies with such luminosities might be challenging with current facilities, even if they have similar EW as in \\bzk4892 and \\bzk8608." }, "1101/1101.5859_arXiv.txt": { "abstract": "Spectroscopic and spectropolarimetric observations of the pre-main sequence early-G star HD 141943 were obtained at four observing epochs (in 2006, 2007, 2009 and 2010). The observations were undertaken at the 3.9-m Anglo-Australian Telescope using the UCLES echelle spectrograph and the SEMPOL spectropolarimeter visitor instrument. Brightness and surface magnetic field topologies were reconstructed for the star using the technique of least-squares deconvolution to increase the signal-to-noise of the data. The reconstructed brightness maps show that HD 141943 had a weak polar spot and a significant amount of low latitude features, with little change in the latitude distribution of the spots over the 4 years of observations. The surface magnetic field was reconstructed at three of the epochs from a high order ($l$ $\\le$ 30) spherical harmonic expansion of the spectropolarimetric observations. The reconstructed magnetic topologies show that in 2007 and 2010 the surface magnetic field was reasonably balanced between poloidal and toroidal components. However we find tentative evidence of a change in the poloidal/toroidal ratio in 2009 with the poloidal component becoming more dominant. At all epochs the radial magnetic field is predominantly non-axisymmetric while the azimuthal field is predominantly axisymmetric with a ring of positive azimuthal field around the pole similar to that seen on other active stars. ", "introduction": "\\label{Sec_int} The generation of magnetic fields is arguably one of the most important process operating inside a star affecting everything from the angular momentum evolution of the star through to the habitability of any planets around the star. In the solar case the magnetic dynamo is believed to operate in an interface layer between the differentially rotating convective zone and the radiative zone which rotates as a solid-body \\citep{ParkerEN:1993}. However, for young, rapidly-rotating solar-type stars evidence is growing that such stars may in fact have a different dynamo mechanism. It has been suggested \\citep{DonatiJF:2003} that such stars may in fact house distributed dynamos, i.e.\\ dynamos which operate across the entire convective zone, rather than being restricted to the interface-layer as in the solar case. The strongest evidence for this are large regions of near-surface azimuthal magnetic field seen on such stars \\citep[i.e.][]{DonatiJF:2003, MarsdenSC:2006a}. Such regions are believed to be the near-surface toroidal components of the large-scale dynamo field. In a solar-like interface-layer dynamo such fields should not be seen near the stellar surface and thus it is postulated that a distributed dynamo is operating in young, rapidly-rotating solar-type stars. Most current dynamo models are based on our understanding of the Sun and are tailored to reproduce solar observations \\citep[see][]{ParkerEN:1993, CharbonneauP:2005}. Such models usually involve an interface layer dynamo which, as discussed, may not apply to young solar-type stars, but there are some different dynamo models that have been examined. For example, the dynamo operating in fully convective stars \\citep[i.e.][]{BrowningMK:2008}, or near-surface dynamos \\citep[i.e.][]{BrandenburgA:2005} or rapidly-rotating solar models with no interface layer \\citep[i.e.][]{BrownBP:2010}. Still the operation of stellar magnetic dynamos is not well understood. One of the most direct ways of observing the stellar dynamo is through the observation of the global magnetic field on the surface of a star. Zeeman Doppler imaging \\citep[ZDI,][]{SemelM:1989, DonatiJF:1997b} has been used for a number of years now to observe the magnetic field configurations of solar-type stars \\citep[i.e.][]{DonatiJF:1992, DonatiJF:1999a, DonatiJF:1997a, DonatiJF:1999b, PetitP:2004a, PetitP:2004b, PetitP:2008, MarsdenSC:2006a, JeffersSV:2008, DunstoneNJ:2008} and there is now a growing body of observations of the surface topologies of other stars that any stellar dynamo models needs to be able to reproduce. \\citet{DonatiJF:2009} have summarised many of the observations of the magnetic fields of non-degenerate stars and for young solar-type stars (i.e.\\ those with Rossby numbers $\\la$ 1 and more massive than $\\sim$0.5 M\\subs{\\odot}) they find that these stars produce substantial toroidal fields and have mostly non-axisymmetric poloidal fields. For more mature stars, \\citet{PetitP:2008} have shown that for solar-type stars with rotation periods faster than $\\sim$12 days the global magnetic field of the star appears to be dominated by toroidal field with poloidal field being by far the most dominant field configuration for stars with slower rotation rates. This could indicate a change in the dynamo mechanism for solar-type stars rotating faster than $\\sim$12 days. Such observations have yet to be modelled in stellar dynamo theory. Doppler images of the surface spot topology of young solar-type stars show that they almost universally have large polar spot features \\citep[i.e.][]{BarnesJR:2000, BarnesJR:2001a, BarnesJR:2001b, DonatiJF:2003, MarsdenSC:2005b, MarsdenSC:2006a}. This is in contrast to the predictions of current dynamo theory \\citep[i.e.][]{SchusslerM:1996, GranzerT:2000, GranzerT:2004} and it has been suggested that a strong meridional flow may be responsible for transporting these spots from their emergence latitudes to the polar regions. This may also explain the mixed polarity of magnetic fields seen on the polar regions of many stars \\citep[see][]{MackayDH:2004}. Although great progress has been made, our knowledge of stellar magnetic dynamos is still in its infancy. How the operation of the stellar magnetic dynamo depends on basic stellar parameters such as age, mass and rotation rate is still unknown. One of the most obvious effects of the solar dynamo is the reversal of the Sun's global magnetic topology every $\\sim$11 years, but so far only two other solar-type stars have shown evidence of global magnetic polarity reversals, HD 190771 \\citep{PetitP:2009} and Tau Boo \\citep{DonatiJF:2008, FaresR:2009}. No strong evidence for a global polarity reversal has yet been seen on a young solar-type star, even though some, for example AB Dor \\citep{DonatiJF:1997a, DonatiJF:1999a, DonatiJF:2003} have been observed for a number of years. Thus it is still unknown if young solar-type stars undergo regular magnetic cycles like the Sun or have chaotic cycles, as indicated by the Calcium HK emission of young stars \\citep{BaliunasSL:1995}. For young solar-type stars the available spectropolarimetric observations are mainly of K-stars \\citep[i.e.][]{DonatiJF:2003}. At present there is only one young early-G star for which spectropolarimetric observations have been published in a refereed journal, HD 171488 \\citep{MarsdenSC:2006a, JeffersSV:2008, JeffersSV:2010}, with preliminary results for the late-F star HR 1817 \\citep{MengelM:2006, MarsdenSC:2006b, MarsdenSC:2010a} being published as conference proceedings. As part of a study into the magnetic topology and cycles of young late-F/early-G stars this paper along with that of \\citet[][Paper II]{MarsdenSC:2010b} and \\citet{WaiteIA:2010}, presents spectropolarimetric observations of two early-G pre-main sequence (PMS) stars. This paper and Paper II deals with the observations of the young early-G star HD 141943, while the \\citet{WaiteIA:2010} paper deals with observations of the similarly aged but more massive HD 106506. HD 141943 was first identified by \\citet{WaiteIA:2005} as a potential target for spectropolarimetric observations. It is active \\citep[Log(L\\subs{X}) = 30.7 erg s\\sups{-1},][]{CutispotoG:2002} and bright \\citep[V = 7.9,][]{CutispotoG:2002}. It also has a very strong Lithium line \\citep[A\\subs{Li} = 3.3,][]{CutispotoG:2003} indicating its youth \\citep[$\\sim$15 Myrs,][]{CutispotoG:2003} and it is what we class as a moderately rapid rotator with a \\vsinis $\\sim$38 \\kmss \\citep{CutispotoG:2003} and a period of 2.20 $\\pm$ 0.03 days \\citep{CutispotoG:1999}. According to \\citet{HillenbrandLA:2008} HD 141943 is also possibly host to a $\\sim$85 K debri disk. This paper (Paper I) describes the evolution in both the brightness (reconstructed at 4 epochs) and magnetic (reconstructed at 3 epochs) topologies of HD 141943 taken at the Anglo-Australian telescope (AAT). A further paper on HD 141943 (Paper II) discusses the differential rotation, H$\\alpha$ emission and coronal magnetic field maps reconstructed from these observations. ", "conclusions": "\\label{Sec_con} HD 141943 is one of the youngest and most massive stars for which we have analysed the surface magnetic topologies for. For young stars it has the shallowest convective zone that we have yet studied. We have presented reconstructed brightness (at 4 epochs) and magnetic (at 3 epochs) topologies of HD 141943. During the four year timebase of the brightness images the spot distribution on HD 141943 has changed very little with a smallish polar spot and a number of lower-latitude features at around 0\\sups{\\circ} to +30\\sups{\\circ} latitude, seen at all epochs. Over three years of observations the pattern of the magnetic topology of HD 141943 also looks at first glance to have experienced relatively little change, with positive and negative radial magnetic field seen at all latitudes and a ring of positive azimuthal magnetic field seen at all epochs. At all epochs the large-scale poloidal magnetic field of HD 141943 is mostly non-axisymmetric while the toroidal field is predominantly axisymmetric. The reconstructed magnetic topologies are rather complex with over 50 per cent of the magnetic energy in components higher than an octupole. When the magnetic images were analysed in more detail we found tentative evidence for a change in the magnetic field in 2009. The ratio of poloidal to toroidal field on HD 141943 goes from almost balanced in 2007 to being heavily dominated by poloidal magnetic field in 2009 and back to balanced in 2010. If real, this variation would indicate magnetic evolution on HD 141943." }, "1101/1101.2362_arXiv.txt": { "abstract": "Based on the observations of the Chinese Small Telescope ARray (CSTAR), the $i$ band observing conditions at Antarctic Dome A have been investigated. The over all variations of sky brightness and transparency are calculated and subsequently cloud cover, contributions to the sky background from various factors including aurorae are derived. The median sky brightness of moonless clear nights is about 20.5 mag arcsec$^{-2}$ in the SDSS $i$ band at the South Celestial Pole, which contains the diffused Galactic light of about 0.06 mag. There are no thick clouds in the year of 2008. Relatively strong aurorae are detected by their brightening the normal sky, which contribute up to about 2\\% of the observed images. ", "introduction": "In site selections for ground-based optical and IR astronomy, some of the most important parameters are the night-sky brightness, seeing, atmospheric transparency, clear nights and humidity. The number of clear nights provides the usable observing time. The sky background, seeing and atmospheric transparency have effect on the observation quality and depth. Humidity affects the telescope and infrared transparency. The Antarctic plateau offers some attractive advantages for ground-based astronomical observations. It is a unique continent where there is no contamination from the artifacts. The average altitude in Antarctic plateau is more than 3000 m and the temperature is very low. Site testing over the past decade has revealed that Antarctica, relative to temperate latitude observatories, has lower infrared sky brightness, better free-atmosphere seeing, greater transparency, a lower turbulent boundary layer, and much lower water vapor content (see, e.g., reviews by Ref.~\\citenum{ari05a,bur05,sto07}). Dome A is located at the highest peak of the continent. It has an elevation of about 4093 m. The year-round average temperature is about $-50^\\circ \\mathrm{C}$, dropping to as low as $-80^\\circ\\mathrm{C}$ on occasion. Such high altitude, low temperature and special geographical position make it be very dry, cold and windless. It might be reasonably predicted that Dome A could be as good as or even a better astronomical site than Dome C (altitude of about 3250 m), with better seeing, higher transparency, and thinner surface layer. Ref.~\\citenum{sau09} compared the sites Dome A, B, C, F and Ridge A and B in their cloud cover, free-atmosphere seeing, precipitable water vapor, temperature, and auroral emission, and concluded that, overall, Dome A might be the best of the existing bases for astronomical observations. It is necessary that some facilities for site tests should be installed at Dome A to assess the site quality. In 2008 January, two Chinese astronomers deployed several instruments at Dome A including a small telescope array named CSTAR. With the observation data of 2008, we leaned the observing conditions such as the sky brightness, transparency and cloud covers. ", "conclusions": "" }, "1101/1101.5297_arXiv.txt": { "abstract": "{ A series of stellar models of spectral type G is computed to study the rotation laws resulting from mean-field equations. The rotation laws of the slowly rotating Sun, the fast rotating MOST stars $\\epsilon$ Eri and $\\kappa^1$ Cet and the rapid rotators R58 and LQ Lup can easily be reproduced. We also find that differences in the depth of the convection zone cause large differences in the surface rotation law and that the extreme surface shear of HD 171488 can only be explained with a artificially shallow convection layer.\\\\ We also check the thermal wind equilibrium in fast-rotating G dwarfs and find that the polar subrotation (${\\rm d} \\Omega/{\\rm d} z<0$) is due to the barocline effect and that the equatorial superrotation (${\\rm d} \\Omega/{\\rm d} r>0$) is due to the $\\Lambda$ effect as part of the Reynolds stresses. In the bulk of the convection zones where the meridional flow is slow and smooth the thermal wind equilibrium actually holds between the centrifugal and the pressure forces. It does not hold, however, in the bounding shear layers including the equatorial region where the Reynolds stresses dominate.} ", "introduction": "Many stars show signs of differential rotation. The solar equator rotates with a shorter period than the polar caps. The difference of 132 nHz between the rotation periods found by \\cite{ulrich88} corresponds to a difference of 0.07 rad/day between the respective angular velocities or a lapping time of 88 days. Helioseismology has found that this pattern persists throughout the whole convection zone but not in the radiative zone below (\\cite{thompson03}). Stellar differential rotation can be inferred from the light curves of rotating spotted stars (see, e.g. Henry et al.~1995; Messina \\& Guinan 2003), from monitoring magnetic activity (Donahue et al.~1996), spectroscopically (\\cite{reiners03}), or by Doppler imaging (\\cite{barnes2000}; \\cite{donati2000}). While several studies have found a systematic dependence of the surface differential rotation on the rotation rate, no such dependence is found when the samples are combined (Hall 1991; \\cite{barnes05}). Moreover, measurements of surface differential rotation of the rapidly rotating K dwarfs PZ Tel and AB Dor with the Doppler imaging technique show that stars rotating much faster than the Sun show very similar surface shear values -- as first predicted by Kitchatinov \\& R\\\"udiger (1999). PZ Tel is a young K dwarf with a rotation period of 0.95 days. Its surface differential rotation $\\delta \\Omega = 0.075$ rad/day is remarkably close to that of the Sun (\\cite{barnes2000}). AB Dor is a rapidly rotating K0 dwarf with a rotation period of 0.51 days. Its surface differential rotation was found to vary with time between 0.09 and 0.05 rad/day (\\cite{cameron02}). Barnes et al. (2005) proposed a dependence on the effective temperature (albeit with large scatter) and found the power law \\bege \\delta \\Omega = \\Omega_{\\rm eq}-\\Omega_{\\rm pole} \\propto T_{\\rm eff}^{8.92\\pm0.31} \\label{barneslaw} \\ende where $\\delta \\Omega$ is the difference between the angular velocities at the equator and the polar caps \\footnote{the strength of the differential rotation is also expressed in terms of the lapping time between the equator and the poles, $ P_{\\rm lap} = {2 \\pi}/{\\delta \\Omega} $}. The power law (\\ref{barneslaw}) was confirmed by the findings using spectroscopic methods (Reiners 2006). The light curves of the stars $\\epsilon$ Eri and $\\kappa^1$ Cet recorded by the MOST satellite (\\cite{croll06}; \\cite{walker07}) both indicate with $\\delta\\Omega\\simeq 0.062$ and $\\delta\\Omega\\simeq 0.064$ a very similar equator-pole difference of the rotation rate as the Sun. Also the value of 0.11 for the young G star CoRoT-2a (\\cite{froehlich09}) well fits the common picture of a rotation-independent surface shear for G stars. It is challenged, however, by recent observations of the young G dwarfs LQ Lup, R58 and HD171488. Marsden et al.~(2005) report a surface differential rotation of $0.025 \\pm 0.015$ rad/day for R58 while \\cite{jeffers09b} find the much larger value of $0.138 \\pm 0.011$ rad/day. Donati et al. (2000) find a surface differential rotation $\\delta \\Omega = 0.12 \\pm 0.02$ rad/day for LQ Lup. \\cite{jeffers09a} determined the surface rotation of HD 171488 using the Zeeman-Doppler imaging technique and found a very strong surface differential rotation of 0.5 rad/day. Marsden et al. (2006) report a smaller but still large value of $0.402\\pm0.044$ rad/day. Huber et al.~(2009) found no evidence for differential rotation at all but could not rule it out either. Jeffers et al.~(2010) confirmed the findings of Marsden et al.~(2006). The values found for LQ Lup and R58 are in line with the Barnes et al. picture but the large values found for HD 171488 are not. The three G dwarfs are similar in their ages, effective temperatures and radii yet HD 171488 shows a much stronger differential rotation than the other two stars. The studies mentioned have focused on rotation rate and effective temperature as the properties determining the surface differential rotation. As the stars are very similar by their stellar parameters such as age and effective temperature, we ask if there could be a difference in their internal structure that would cause the observed difference in their surface rotation. All three stars have just reached the zero age main sequence or are approaching it. Given the rapid retreat of the convection zone in the final part of the pre-main sequence evolution and the uncertainty of the stellar age we ask if the strong differential rotation of HD 1714888 can be explained by a different depth of its outer convection zone. In the following we compute model convection zones of different depth and their large-scale gas motions, i.e.~rotation and meridional flow. The models are based on the mean field formulation of fluid dynamics which has been very successful for the Sun, where the models reproduce the surface rotation, the internal in the convection zone, and the surface meridional flow very well (Kitchatinov \\& R\\\"udiger (1999, KR99), K\\\"uker \\& Stix (2001, KS01)). Models have been constructed for a variety of stars with spectral types from M to F (K\\\"uker \\& R\\\"udiger 2008). A new scheme allows the computation of stellar rotation laws and meridional flow patterns based on a mean-field model of the large-scale flows in stellar convection zones also for fast rotation rates when narrow boundary layers exist. It assumes strict spherical symmetry for the basic stratification, ignoring any flattening that might occur for very fast rotation. However, the impact of rotation on the thermal structure can be taken into account by including a gravity darkening term in the heat transport equation so that the model remains applicable even for moderately flattened stars. ", "conclusions": "Stellar differential rotation is driven by Reynolds stress, by the centrifugal-induced (`Biermann-Kippenhahn') flow and (for stratified density) by the baroclinic flow. In our model both meridional circulations have opposite directions. The barocline flow becomes important for faster rotation as the convective heat-flux deviates from the radial direction. This deviation can be interpreted as a tilt of the heat-flux vector towards the rotation axis and causes the poles to be slightly warmer than the lower latitudes. The baroclinic flow is responsible for the strong negative gradient ${\\rm d}\\Omega/{\\rm d} z$ along the rotation axis, and the Reynolds stress produces the typical positive ${\\rm d}\\Omega/{\\rm d} r$ along the equatorial midplane both known from helioseismology. As a result of both impacts the isolines of $\\Omega$ known for the Sun are almost radial in mid-latitudes. Studying real stellar models implies varying a bunch of parameters as mass, radius, effective temperature and the depth of the convection zone are interdependent. Kitchatinov \\& Olemskoy (2010) studied differential rotation along the lower ZAMS for fixed rotation rate and found that the surface shear is a function of the effective temperature alone. Using some artificial models, we find that for the same effective temperature the depth of the convection zone has a big impact on stellar rotation. Shallow convection zones produce stronger surface shear. The presented mean-field theory for G stars naturally explains the rotation laws of the Sun, of the MOST stars $\\epsilon$ Eri and $\\kappa^1$ Cet, and of such fast-rotating stars like R58 and LQ Lup. The discrepancy between these stars and the much stronger differential rotation observed for HD 171488 -- if real -- is hard to explain for stars of similar age and spectral type. If, however, HD 171488 had a shallower convection zone for some reason, its strong surface differential rotation follows immediately. During the pre-MS evolution the convection zone retreats from the central region and the originally fully convective star forms a radiative zone around the core. This change in the depth of the convection zone should be reflected by the surface differential rotation." }, "1101/1101.0684_arXiv.txt": { "abstract": " ", "introduction": "Both professional and amateur astronomers have been studying the skies for centuries. Their respective r\\^{o}les, however, have changed considerably. In the 18$^{\\rm th}$ and 19$^{\\rm th}$ century, for instance, one of the main tasks of professional astronomers was to calculate astronomical data to be used by the merchant fleet and the military, and to provide the society at large with important data such as the times of sunrise and sunset and, in fact, time itself. Amateur astronomers meanwhile entertained themselves with what most astronomers might now consider more interesting activities, such as discovering planets and comets, and observing nebulae. Well-known amateurs include Caroline Herschel (1750-1848) who discovered several comets, and her brother William (1738-1822) who discovered Uranus, several moons of that planet and of Saturn, created a catalogue of nebulae (a term used at the time to describe any extended object), and observed double stars. He constructed hundreds of telescopes, and made a number of important discoveries related to light and radiation. Needless to say, amateur astronomers at the time were wealthy individuals. While in the 19$^{\\rm th}$ century the interests of professional and amateur astronomers started to overlap more and more, they diverged again during the 20$^{\\rm th}$ century. One of the main reasons for this is that professionals started to use expensive and exclusive telescopes, such as the ones used by Hale and Hubble in California, which were out of reach of all amateurs. We are now at the start of the 21$^{\\rm st}$ century, and one of the characteristics of our modern times is the almost ubiquitous availability of high-quality yet relatively cheap technology. So whereas professional astronomers now use very advanced instruments, including very large optical and radio telescopes, and massive supercomputer power, interested amateurs can start to use technologically advanced telescopes, cameras, and computers at a cost which is accessible to many millions of citizens around the world. In the next Section of this paper, I will summarise some of the areas in which amateurs can, and do, use these tools to help advance professional astronomy. ", "conclusions": "" }, "1101/1101.4825_arXiv.txt": { "abstract": "We carry out the first time-dependent numerical MagnetoHydroDynamic modeling of an extrasolar planetary system to study the interaction of the stellar magnetic field and wind with the planetary magnetosphere and outflow. We base our model on the parameters of the HD 189733 system, which harbors a close-in giant planet. Our simulation reveals a highly structured stellar corona characterized by sectors with different plasma properties. The star-planet interaction varies in magnitude and complexity, depending on the planetary phase, planetary magnetic field strength, and the relative orientation of the stellar and planetary fields. It also reveals a long, comet-like tail which is a result of the wrapping of the planetary magnetospheric tail by its fast orbital motion. A reconnection event occurs at a specific orbital phase, causing mass loss from the planetary magnetosphere that can generate a hot spot on the stellar surface. The simulation also shows that the system has sufficient energy to produce hot-spots observed in Ca~II lines in giant planet hosting stars. However, the short duration of the reconnection event suggests that such SPI cannot be observed persistently. ", "introduction": "\\label{sec:Intro} The newly discovered exoplanets are not only newly discovered worlds literally, but are also new worlds in terms of the physical system they introduce us to. Many exoplanets have now been observed since their first discovery \\citep{mayor95, exoplanet95, exoplanets03}. In particular, many Jupiter-like giant planets have been observed at a distance of less than 0.1~AU from their parent star, and some even within $10R_\\star$, where they are essentially located inside the extended stellar corona \\citep{exoplanet95, exoplanets03}. It is not unreasonable to suppose that these giant planets have a substantial internal magnetic field \\citep{sanches-lavega04,durand-Manterola09} and a magnetosphere, which could encompass a significant fraction of the extended coronal volume. Indeed, the Jovian magnetosphere is the largest entity in our solar system after the solar one \\citep{Bagenal04}. Perhaps the most fundamental difference between the close-in planet scenario and magnetospheres in our own solar system is that close-in planets and their magnetospheres can be located within the stellar Alfv\\'en radius. In analogy to the hydrodynamic sonic point (where the flow speed equals the sound speed), the Alfv\\'en radius (or the Alfv\\'enic point) is the distance at which the accelerating stellar wind equals the Alfv\\'en speed, $u_A=B/\\sqrt{4\\pi\\rho}$. Here $B$ is the magnetic field strength and $\\rho$ is the plasma mass density. Beyond the Alfv\\'en point, the flow is super-Alfv\\'enic and magnetic information (or energy) cannot propagate back towards the star and affect the corona. Close-in planets could then have a substantial effect on the structure of the stellar corona, and the electromagnetic interaction between the planetary magnetic field and the corona might generate some observational signatures. In recent years, some signatures of star-planet interaction (SPI) have been observed. In a series of papers, \\cite{shkolnik03, shkolnik05a,shkolnik05b,shkolnik08} have described the observed modulations in the Ca II K emission line, which is an indicator of chromospheric activity. They found a line intensity increase corresponding to chromospheric ``hot spots'' with a period that is correlated with the planetary orbital motion on HD~189733, HD~179949, $\\tau$ Boo, and $\\nu$ And. In some cases the location of these hot spots was aligned with the star-planet radial vector, while in other cases, the location of the intensity peak was shifted from this vector by 70-170 degrees. Some of these observations also revealed an on/off nature \\citep{shkolnik08}. At higher energies, \\cite{kashyap08} performed a statistical survey of systems with close-in giant planets, which revealed that the X-ray flux of these systems is about 30-400\\% higher than the typical fluxes from similar stars with planets located further out in the stellar system. \\cite{Poppenhaeger10} did not find a similar effect for a sample of nearby stars, though they did uncover weak evidence for a combined effect of planetary mass and orbital distance. \\cite{saar08} observed the HD~179949 system in X-rays and found that emission associated with a hot spot likely engendered by SPI contributed $\\approx$30\\% to the emission at a mean plasma temperature of $\\approx$ 1 keV. A similar trend in the UV band has been recently presented by \\citep{shkolnik10}. Stellar coronae are dominated by the stellar magnetic field, which also dictates the structure of the stellar wind \\citep{parker58}. The presence of a close-in planet with a significant internal magnetic field could perturb this stellar field and might affect the large-scale structure of the stellar corona. The planet can affect the corona in different ways. First, a purely hydrodynamic affect is the planet acting as an obstacle, deflecting the flow of plasma (i.e., the stellar wind) around it. Second, an electrostatic effect is the large-scale potential field topology being modified by the superposition of the stellar and planetary magnetic fields. Third, a MagnetoHydroDynamic (MHD) effect is expected due to the orbital motion of the 'external' field (the planetary field) inside the background conducting plasma of the stellar corona and the stellar wind. In any case, the process involves a modification of the coronal magnetized medium via energy transfer between the planetary magnetic field and the corona. Few theoretical SPI models have been developed to explain the observed hot spots, as well as the overall increase in observed X-ray flux \\citep{cuntz00,cranmer07}. In particular, \\cite{lanza08,lanza09} proposed a mechanism to explain the observed phase shift in the location of the spots. This work also demonstrated that energy transfer from the magnetic field to the plasma via magnetic reconnection between the coronal and planetary field (as the planet is moving through the corona) can provide the observed overall energy increase (about $10^{21}~W$). This amount of energy requires rather strong stellar and planetary magnetic fields or alternatively, this energy can be provided assuming the planet triggers magnetic reconnection in a stellar corona that is already energized due to accumulation of magnetic helicity provided by photospheric motions. Performing numerical simulations of SPI is challenging. First, one needs to provide a dynamic model for the ambient stellar corona and stellar wind. Second, it is necessary to introduce the planet, which is an additional boundary condition in the model. Third, for capturing the dynamics due to the planetary motion, one needs to consider the planet as a {\\it time-dependent} boundary condition. \\cite{Ip04}, \\cite{preusse06}, and \\cite{preusse06} have studied the structure of Alfv\\'en wings in the close-in planet around HD~179949 and made the analogy of the magnetic interaction between Io and the Jovian magnetosphere. \\cite{lipatov05} and \\cite{Johansson09} have studied the structure of the planetary magnetosphere using a hybrid code driven by an approximation to the stellar wind as a boundary condition. In both cases, the simulations included only the planetary magnetosphere. In a recent paper, \\cite{cohen09} (hereafter Paper I) have presented two simplified MHD simulations of SPI. In the first case (``Case A''), both the stellar and planetary magnetic fields were represented by magnetic dipoles, and the planet was fixed in the inertial frame of reference. A relative motion due to stellar rotation of the coronal plasma with respect to the planetary magnetic field was obtained. In the second case (``Case B''), the planetary magnetic field was dipolar again, while the stellar magnetic field was obtained using high-resolution {\\sl solar} magnetic field data. This provided a realistic stellar field, in contrast to the idealized dipole field used in Case A. In Case B, the planet was fixed the whole time and the simulation was calculated in the frame of reference that is rotating with the star in a tidally-locked manner \\citep[like the $\\tau$ Boo system,][]{Butler97}. The synthetic X-ray light curves of Case A revealed that there is a drop in the X-ray flux when the planet is behind the star. In Case B, the intensity drop was shifted by $\\approx\\;60^\\circ$ from the star-planet plane. The simulation also showed an increase in the total X-ray flux compare to a reference case without a planet. In paper I, we concluded that the increase in X-ray flux in systems with close-in planets is due to the planet preventing the corona from expanding, so that field lines that would be open remain closed. As a result, the plasma cannot escape and the overall coronal density is higher than in the case without a planet, or with a planet located further from the star. Since the X-ray flux is essentially a line of sight integration of the square of the (electron) density, the observed X-ray flux is higher in such systems. In addition, in paper I we proposed that the observed hot-spots are associated with the magnetic connectivity between the star and the planet, and that the shift in the location of the hot-spots from the star-planet vector is probably due to the complexity of the stellar magnetic field. In a follow-up study to Paper I, \\citet[][hereafter Paper~II]{Cohen10c} investigated the influence of changing the semi-major axis of the co-rotating planet on the stellar mass and angular momentum loss. The disruption of the corona and wind occurs when the separation is sufficiently small that the planetary and stellar Alfv\\'en surfaces start to interact. We found that the spin-down of stars harboring close-in planets is reduced compared to that of stars with distant planets or no planets at all. While the results of the simulations presented in Papers I and II are significant, both cases were stationary and did not capture all of the dynamical effects that are predicted in the theoretical descriptions of SPI (see above). In particular, magnetic reconnection between the stellar field and the coronal field was not well captured, even in Paper~I Case A where some relative motion between the magnetic bodies was introduced. In this paper, we present a dynamical MHD simulation, where we include {\\it time-dependent} circular orbital motion of the planet around the star. The model is based on the observed properties of the HD 189733 system. However, we emphasize that the main goal of this study is to investigate the interaction of a close-in giant planet with the stellar corona of its host star as a fundamental plasma physics problem, rather than to model HD~189733 specifically. The structure of the paper is as follows. We present the numerical model in Section~\\ref{sec:Model}. The results are presented in Section~\\ref{sec:Results}, and our main findings are discussed in Section~\\ref{sec:Discussion}. The results are summarized in Section~\\ref{sec:Conclusions}. ", "conclusions": "\\label{sec:Conclusions} We have performed a realistic MHD simulation of SPI using the parameters of the HD 189733 planetary system. The simulation is based on the observed stellar surface magnetic field and includes the stellar wind, time-dependent planetary orbital motion and planetary outflow. We choose to study two cases. One with strong planetary magnetic field and high base density and one with weak planetary field and low base density. Based on the simulation results, our main findings are: \\begin{enumerate} \\item The use of a realistic stellar field results in a highly structured stellar corona characterized by sectors with different plasma properties. \\item As the planet orbits the star, it interacts with the different sectors. Each sector crossing, as well as the transitions between the sectors, leads to changes in the SPI. \\item The planet is followed by a long, comet-like tail of plasma that represents its magnetospheric tail wrapped around due to the rapid planetary orbital motion. This tail is almost perpendicular to the stellar wind direction and can lead to magnetospheric dynamics which differ to those in the case in which the magnetospheric tail is aligned with the wind direction. \\item A large reconnection event occurs when the planetary and stellar field are opposite in polarity, where the reconnection site varies depends on the planetary field strength. This reconnection event causes a release of cold plasma from the planetary magnetosphere with a mass flux of $\\dot{M}_p=2.5\\cdot 10^{-12}\\;M_J\\;yr^{-1}$. However, this escape of material can be defined as a magnetospheric escape rather than atmospheric escape. The latter cannot be define by the simulation presented here since it does not include a self-consistent atmospheric escape mechanism and is only controlled by a boundary value for the planetary density; \\item The short duration of the reconnection event suggests that such SPI cannot be observed persistently, and in particular when such signatures are being searched for in radio bands. \\item The simulation indicates that a sufficient amount of energy can be dissipated by the star-planet interaction to explain the hot-spots observed in the Ca~II lines of stars hosting close-in planets. \\item The total mass flux carried by the stellar wind is modulated by the planetary orbit and the particular value at any time depends on the magnetic topology. \\end{enumerate} In this work, we used the parameters of HD 189733, with planetary semi-major axis of $a=8.8R_\\star$. It is reasonable to believe that with a smaller value of $a$, some of the SPI effects in the simulation would be enhanced. Further studies of SPI involving a physical atmospheric escape from the planet, as well as the impact of stellar coronal mass ejections on the planetary magnetosphere, would be highly motivated." }, "1101/1101.1692.txt": { "abstract": "We report on the first search for atmospheric and for diffuse astrophysical neutrino-induced showers (cascades) in the IceCube detector using 257 days of data collected in the year 2007-2008 with 22 strings active. A total of 14 events with energies above 16\\,TeV remained after event selections in the diffuse analysis, with an expected total background contribution of $8.3\\pm 3.6$. At 90\\% confidence we set an upper limit of $E^2\\Phi_{90\\%CL}<3.6\\times10^{-7}\\,\\mathrm{GeV \\cdot cm^{-2} \\cdot s^{-1}\\cdot sr^{-1}}$ on the diffuse flux of neutrinos of all flavors in the energy range between $24$\\,TeV and $6.6$\\,PeV assuming that $\\Phi \\propto E^{-2}$ and that the flavor composition of the $\\nu_e : \\nu_\\mu : \\nu_\\tau$ flux is $1 : 1 : 1$ at the Earth. The atmospheric neutrino analysis was optimized for lower energies. A total of 12 events were observed with energies above 5 TeV. The observed number of events is consistent with the expected background, within the uncertainties. ", "introduction": "The origin of high-energy cosmic-rays is an area of active research in astrophysics. The sites where cosmic rays are accelerated are expected to produce high energy neutrinos. Many types of objects, ranging from supernovae and gamma-ray bursters to active galactic nuclei~\\cite{theory}, have been proposed as point sources of high energy neutrinos and many searches for such sources have been made~\\cite{ICSearches}, yielding results consistent with background only assumptions. If there are many point sources, each with an unobservably low flux, then the aggregate flux may still be observable as a diffuse flux. Diffuse searches rely on the energy spectrum of the detected events to separate an extraterrestrial signal from atmospheric neutrinos produced in the interaction of cosmic rays with atomic nuclei in the Earth's atmosphere. Low energy (below $\\sim10$\\,GeV) atmospheric muon and electron neutrinos have been observed in underground detectors~\\cite{lowenu}. At higher energies, from 100\\,GeV to 400\\,TeV, neutrino telescopes have measured the spectrum of atmospheric $\\nu_\\mu$~\\cite{IceCube2010}. In this energy range, the flux of $\\nu_e$ is expected to be lower by about a factor of 20~\\cite{beacom} and has not been observed. The main component of the atmospheric neutrino spectrum is produced by the decays of $\\pi^\\pm$ and $K^\\pm$. Asymptotically it can be parametrized by $dN_\\nu/dE_\\nu \\propto E_\\nu^{-3.7}$, where $E_\\nu$ is the neutrino energy~\\cite{ICCRspectrum}. Decays of hadrons containing charm and bottom quarks form an additional component that is expected to be close to the primary cosmic-ray spectrum, $dN_\\nu/dE_\\nu \\propto E_\\nu^{-2.7}$~\\cite{naumov,enberg,martin}, and produces nearly equal numbers of $\\nu_\\mu$ and $\\nu_e$. These prompt neutrinos are expected to dominate the $\\nu_e$ spectrum at energies above $\\sim30\\,\\mathrm{TeV}$~\\cite{beacom}. The production of $\\nu_\\tau$ is expected to be negligible. Fermi acceleration of charged particles in magnetic shocks followed by collisions with matter or radiation between the source and the Earth naturally leads to an energy spectrum for extraterrestrial neutrinos that is harder than that for atmospheric neutrinos, typically close to $dN_\\nu/dE_\\nu \\propto E_\\nu^{-2}$. This allows diffuse extraterrestrial neutrinos to be visible as a hard component to the observed spectrum. The ratio of the $\\nu_e:\\nu_\\mu:\\nu_\\tau$ flux in a single astrophysical source depends on the neutrino energy~\\cite{kashti05}. At moderate (high) energies, the neutrino flavor flux ratio behaves like the one from a pure pion (muon-damped) source, leading to an observed $1:1:1$ (1:1.8:1.8) ratio at the Earth after taking into account neutrino oscillations. The energy at which a flavor ratio transition occurs thus depends on the properties of the source~\\cite{kashti05}. The neutrino flux is not known, although it is expected to be below the Waxman-Bahcall bound~\\cite{WB}. Previous searches for a diffuse flux have been performed with muon neutrinos \\cite{ICCRspectrum}, and with cascades~\\cite{ICCascades,OxanaPaper}. Cascades are the particle showers (electromagnetic and hadronic) initiated by charged current interactions of $\\nu_e$ and $\\nu_\\tau$ and the neutral-current neutrino interactions of all three flavors in a medium. In the charged-current interactions, an average of $80$\\% of the (high energy) neutrino energy goes into the produced lepton~\\cite{gandhi}. For $\\nu_e$, this leads to an electromagnetic shower, while for $\\nu_\\tau$ the character of the lepton-induced shower depends on the $\\tau$ decay mode. The remainder of the energy is transferred to the target nucleon, producing a hadronic cascade. In the neutral-current interactions, the neutrino transfers a fraction of its energy to the target nucleon producing only a hadronic cascade. A typical cascade deposits its electromagnetic energy in a thin cylinder about 30 cm in radius and 5 m in length. Hadronic energy is deposited over a larger volume, about 11 m long and 75 cm in diameter. IceCube observes the Cherenkov radiation produced by the charged secondary particles from neutrino-nucleon interactions through an optical sensor array. While a neutrino-induced muon has a track-like signature in IceCube, a cascade event looks effectively like a point source of Cherenkov light in the detector. For diffuse searches, cascades from all flavor $\\nu$ interactions have two advantages over tracks from $\\nu_\\mu$ interactions, despite their inherently poor angular resolution compared to muon tracks. The first is that the background from atmospheric neutrinos is lower than for $\\nu_\\mu$. Second, because of their short shower length, the cascades are well-contained in the detector, with a Cherenkov light output proportional to the shower energy, so the shower energy is well measured. The detector acts as a calorimeter. Since the energy spectrum of extraterrestrial neutrinos is expected to be harder than the atmospheric neutrino spectrum, searching for a break in the energy spectrum with cascades is easier than with muons, both due to the expected break being at a lower energy in the cascade channel than the muon channel (a consequence of lower fluxes of atmospheric $\\nu_{e}$ than $\\nu_{\\mu}$), and better intrinsic energy resolution of cascades over muons. This paper reports on searches for diffuse extraterrestrial and for atmospheric neutrino-induced cascades using $257$ days (livetime) of data collected in the year 2007-2008 with a partially completed IceCube detector consisting of 22 of the planned 86 strings. The IceCube detector and data sample are described in section~\\ref{sec:detector}. Section~\\ref{sec:analysis} describes the analysis. Results are given in section~\\ref{sec:results}, and a summary follows in section~\\ref{sec:summary}. \\begin{figure} \\centering \\hspace*{-0.35cm}\\includegraphics[width=0.3\\textwidth]{fig1a} \\includegraphics[width=0.34\\textwidth]{fig1b} \\caption{(Color online) a) The filled circles show the positions of the strings in the $x-y$ (horizontal) plane for the $22$-string detector configuration. The lines show the boundaries of the fiducial volume and are described in the text. b) The reconstructed center-of-gravity position $x$ (COGX) versus $y$ (COGY) for events passing the cascade online filter. The right axis shows the rate [Hz].} \\label{fig:fiducial} \\end{figure} %% DETECTOR AND DATA ", "conclusions": "\\label{sec:summary} In summary, we report the first search for cascades induced by atmospheric and by diffuse astrophysical neutrinos with the IceCube detector. The data, obtained in 2007--2008 with a configuration of 22 active strings, amount to 257 days of livetime and was searched for charged current interactions of $\\nu_e$ and $\\nu_\\tau$, and for neutral current interactions of neutrinos of all flavors. The atmospheric neutrino analysis used neural-network based event selections and resulted in a total of $12$ candidate events with energies above $5$ TeV after event selections. Within the large uncertainties, the observed number of events is consistent with the expected background. The astrophysical neutrino analysis used one and two dimensional selection criteria and was optimized for higher energies than the atmospheric neutrino analysis. A total of 14 events with energies above 16\\,TeV remained after event selections, with an expected total background contribution of $8.3 \\pm 3.6$ events. We derive an upper limit at 90\\% confidence of $E^2 \\Phi_{90\\%CL} < 3.6\\times10^{-7}\\,\\mathrm{GeV \\cdot cm^{-2} \\cdot s^{-1} \\cdot sr^{-1}}$ on the diffuse flux of astrophysical neutrinos with the assumption that the energy spectrum $\\Phi \\propto E^{-2}$ and that the flavor composition of the $\\nu_e : \\nu_\\mu : \\nu_\\tau$ flux is $1 : 1 : 1$ at the Earth. In this limit, 90\\% of the expected signal events have energies between $24$\\,TeV and $6.6$\\,PeV. This is below the limit that was recently reported from final AMANDA data, corresponding to 1001 days of livetime~\\cite{OxanaPaper}. Once construction is completed in 2011, IceCube will consist of 86 strings covering a volume of $1\\,\\mathrm{km}^3$. Future IceCube searches will thus benefit from a considerably larger size and are expected to have significantly improved detection sensitivity." }, "1101/1101.0547_arXiv.txt": { "abstract": "Multiple recent investigations of solar magnetic field measurements have raised claims that the scale-free (fractal) or multiscale (multifractal) parameters inferred from the studied magnetograms may help assess the eruptive potential of solar active regions, or may even help predict major flaring activity stemming from these regions. We investigate these claims here, by testing three widely used scale-free and multiscale parameters, namely, the fractal dimension, the multifractal structure function and its inertial-range exponent, and the turbulent power spectrum and its power-law index, on a comprehensive data set of 370 timeseries of active-region magnetograms ($17,733$ magnetograms in total) observed by SOHO's {\\it Michelson Doppler Imager} (MDI) over the entire Solar Cycle 23. We find that both flaring and non-flaring active regions exhibit significant fractality, multifractality, and non-Kolmogorov turbulence but none of the three tested parameters manages to distinguish active regions with major flares from flare-quiet ones. We also find that the multiscale parameters, but not the scale-free fractal dimension, depend sensitively on the spatial resolution and perhaps the observational characteristics of the studied magnetograms. Extending previous works, we attribute the flare-forecasting inability of fractal and multifractal parameters to {\\it i)} a widespread multiscale complexity caused by a possible underlying self-organization in turbulent solar magnetic structures, flaring and non-flaring alike, and {\\it ii)} a lack of correlation between the fractal properties of the photosphere and overlying layers, where solar eruptions occur. However useful for understanding solar magnetism, therefore, scale-free and multiscale measures may not be optimal tools for active-region characterization in terms of eruptive ability or, ultimately, for major solar-flare prediction. ", "introduction": "\\label{S-Intro} The ever-increasing remote-sensing capabilities of modern solar magnetographs have led to the undisputed conclusion that solar (active region in particular) magnetic fields exhibit an intrinsic complexity. ``Complexity'' is a term commonly used to describe an array of properties with one underlying characteristic: a scale-invariant, self-similar (fractal) or multiscale (multifractal) behavior. The measured photospheric magnetic fields in active regions are indeed {\\it multifractal} ({\\it e.g.} \\opencite{Lawrence_etal93}; \\opencite{Abramenko_05}), that is, consisting of a number of fractal subsets. As such, they are also {\\it fractal}, with a fractal dimension equal to the maximum fractal dimension of the ensemble of fractal subsets. Fractality is a mathematical property but with important physical implications. Scale-free or multiscale manifestations are thought to stem from an underlying self-organized, or self-organized critical (SOC), evolution in active regions. Self-organization refers to the internal, intrinsic reduction of the various parameters (also called degrees of freedom) of a nonlinear dynamical system, such as a solar active region, into a small number of {\\it important} parameters that govern the system's evolution and, perhaps, its dynamical response \\cite{Nicolis_Prigogine89}. Assumptions on the nature of just these important parameters can lead to models of active-region emergence and evolution encapsulated in simplified {\\it cellular automata} models \\cite{Wentzel_Seiden92,Seiden_Wentzel96,Vlahos_etal02,Fragos_etal04}. Self-organized criticality, on the other hand, implies that the self-organized system evolves through a sequence of metastable states into a state of {\\it marginal} stability with respect to a critical threshold. Local excess of the threshold gives rise to spontaneous, intermittent instabilities lacking a characteristic size (\\opencite{Bak_etal87}; \\opencite{Bak_96}). The intrinsic self-organization in solar active regions may be attributed to the {\\it turbulence} dominating the emergence and evolution of solar magnetic fields. Tangled, fibril magnetic fields rising from the convection zone can be explained via Kolmogorov's theory of fluid turbulence ({\\it e.g.}, \\opencite{Brandenburg_etal90}; \\opencite{Longcope_etal98}; \\opencite{Cattaneo_etal03}, and others). Turbulence in the generation and ascension of solar magnetic fields leads to turbulent photospheric flows ({\\it e.g.}, \\opencite{Hurlburt_etal95}). Thus, the turbulent photosphere is viewed as a driver that gradually but constantly perturbs an emerged magnetic-flux system, such as an active region, dictating self-organization in it and possibly forcing it toward a marginally stable, SOC state ({\\it e.g.}, \\opencite{Vlahos_Georgoulis04}). Turbulent action does not cease in the photosphere, but it extends into the solar corona. However, coronal low-$\\beta$ turbulence may not be the Kolmogorov fluid turbulence applying to the high-$\\beta$ plasma of the convection zone and the photosphere. Instead, it might be an intermittent magnetohydrodynamic (MHD) turbulence \\cite{Kraichnan_65,Biskamp_Welter89}. Fractal, multifractal, and turbulent properties of photospheric active-region magnetic fields have been intensely studied in recent years. Fractality is traditionally investigated via the {\\it fractal dimension}, often inferred using box-counting techniques ({\\it e.g.}, \\opencite{Mandelbrot_83}). Box-counting is also used for multifractal studies in space and time ({\\it e.g.}, \\opencite{Evertsz_Mandelbrot92}), involving also {\\it generalized correlation dimensions} \\cite{Georgoulis_etal95,Kluiving_Pasmanter96}. A commonly used multifractal method that does not require box counting is the calculation of the {\\it multifractal structure function spectrum} \\cite{Frisch_95}. Moreover, a practical method for quantifying turbulence is the calculation of the turbulent {\\it power spectrum}, stemming from the original work of \\inlinecite{Kolmogorov_41}. If the power spectrum shows a power law over a range of scales, perceived as the turbulent inertial range, its slope determines whether the inferred turbulence is Kolmogorov-like (scaling index $\\approx - 5/3$) or Kraichnan-like (scaling index $\\approx - 3/2$) if either of these two applies. Multiple studies on fractality, multifractality, and turbulence in photospheric active-region magnetic fields have raised claims that flaring active regions exhibit distinct, distinguishable complexity. These works might lead to the impression that fractal, multifractal, or turbulent measures hold significant flare-predictive capability or, at least, they might be used to identify flaring active regions before they actually flare. To summarize some of these works, \\inlinecite{Abramenko_etal03} suggested that a ``peak in the correlation length might be a trace of an avalanche of coronal reconnection events''. \\inlinecite{McAteer_etal05} reported that ``solar flare productivity exhibits an increase in both the frequency and GOES X-ray magnitude of flares from [active] regions with higher fractal dimension''. Further, \\inlinecite{Abramenko_05} found that ``the magnitude of the power index at the stage of emergence of an active region ... reflects its future flare productivity when the magnetic configuration becomes well evolved'', while \\inlinecite{Georgoulis_05} reported that ``the temporal evolution of the [inertial-range] scaling exponents in flaring active regions probably shows a distinct behavior a few hours prior to a flare''. More recently, \\inlinecite{Conlon_etal08} worked on a sample of four active regions and reported evidence for a ``direct relationship between the multifractal properties of the flaring regions and their flaring rate'', while \\inlinecite{Hewett_etal08}, reporting on ``preliminary evidence of an inverse cascade in active region NOAA 10488'' found a ``potential relationship between energy [power-spectrum] scaling and flare productivity''. Many of these works are also reviewed by \\inlinecite{McAteer_etal10}. If the above findings are confirmed, they may well lead to notable improvements in our physical understanding of active regions and in highlighting possible differences between flaring (that is, hosting major flares) and non-flaring (that is, hosting only sub-flares) regions. In \\inlinecite{Georgoulis_05} we studied three different scale-free and multiscale parameters, namely, the fractal dimension, the spectrum of generalized correlation dimensions, and the structure-function spectrum and its inertial-range exponents over a limited magnetogram sample of six active regions, three of them hosting at least one major flare (M- or X-class in the GOES X-ray 1--8 \\AA$\\;$ flare classification scheme). In one case of a X-flaring active region -- NOAA active region (AR) 10030 with an X3 flare at the time of the observations -- we noticed a sharp preflare increase of the inertial-range exponent of the structure functions followed by a significant ($\\approx 20$\\% and much above uncertainties), permanent decrease after the flare. We suggested that this analysis should be repeated on a much larger sample of {\\it both} flaring and non-flaring regions to determine whether this behavior was incidental. In this study we analyze a comprehensive sample of 370 timeseries of active-region magnetograms, with each timeseries corresponding to a different active region. In this sample, 77 active regions hosted at least one M- or X-class flare during the observations and they are considered {\\it flaring} (17 X-class flaring, 60 M-class flaring), while the remaining 293 active regions were not linked to major flares and are hence considered {\\it non-flaring}. We calculate three of the most promising scale-free and multiscale measures on this data set, namely, the fractal dimension, the multifractal structure function spectrum, and the turbulent power spectrum. A detailed description of the data and techniques follows in Section~\\ref{S-data_methods}. In Section~\\ref{S-spres} we test the sensitivity of the calculated parameter values on the spatial resolution of the studied magnetogram. A statistical analysis of the active-region sample is performed in Section~\\ref{S-comp} while Section~\\ref{S-conc} summarizes the study, discusses the results, and outlines our conclusions. ", "conclusions": "\\label{S-conc} This study investigates previous claims on the efficiency of fractal and multifractal techniques as reliable predictors of major solar flares and/or parameters reflecting the overall flare productivity of solar active regions before they actually flare. From the array of parameters implemented in the literature, we select three of the reported most promising ones: the fractal dimension, the multifractal intermittency index, and the scaling index of the turbulent power spectrum. Our objective is not to judge the methods {\\it per se} but, rather, to test the notion of utilizing fractality and multifractality to gain predictive insight into major solar flares. Statistical analyses such as this one must guarantee that the assembled active-region sample is representative: the sample must contain numerous flaring {\\it and} non-flaring regions. Comprehensive statistics often help avoid the interpretation of incidental signals as statistically significant behavior. Section~\\ref{S-prepeak} (Figure \\ref{pre_max}d) includes examples of results that might have been interpreted in a misleading way had the statistics of our active-region sample been insufficient. We study 370 SOHO/MDI low-resolution ($1.98''$ per pixel) timeseries of active-region magnetograms, 293 of which correspond to active regions without major flares and 77 correspond to M- and X-class flaring regions. MDI line-of-sight fields are used for regions within $30^o$ of the central meridian in order to approximate the longitudinal-field component with the normal-field component and avoid any corrections or otherwise modifications of the original MDI data. We find that neither scale-free (fractal) nor multiscale (multifractal) techniques can be used to predict major flares, or for the {\\it a priori} assessment of the flaring productivity of active regions. In particular, we find that their diagnostic capability is not better than that of the unsigned magnetic flux of active regions, a traditional, but unreliable, activity predictor. Since the fractal and multifractal measures tested here are less effective than the unsigned flux (Figure \\ref{cprob}), they should not be used for flare prediction or for flaring productivity assessment. On the fundamental question of whether flaring active regions are more fractal, multifractal, or turbulent than other, non-flaring ones, the answer {\\it per} our results has to be negative: flaring regions tend to exhibit relatively large peak values of scale-free and multiscale parameters but these values, or even higher ones sometimes, are also exhibited by non-flaring regions. For all statistical distributions, the means and standard deviations are such that the different populations of flaring and non-flaring regions overlap considerably (Table \\ref{tab2}). At this point we emphasize our willingness to follow the guidelines of multiple previous studies in the inference of the above fractal and multifractal parameters. In particular, we followed \\inlinecite{McAteer_etal05} when inferring the fractal dimension $D_0$, \\inlinecite{Abramenko_05} when inferring the turbulent scaling index $\\alpha$ (despite the fact that Abramenko worked exclusively on high-resolution MDI magnetograms), and a previous work of this author \\cite{Georgoulis_05}, together with \\inlinecite{Abramenko_etal03}, when inferring the intermittency index $\\zeta (q)$. As a result, the findings of both \\inlinecite{McAteer_etal05} and \\inlinecite{Abramenko_05} were qualitatively reproduced in this analysis, while we showed that the distinct $\\zeta (3)$-behavior reported by \\inlinecite{Georgoulis_05} was just one incidental case and not part of a systematic trend. In addition, this work (Section \\ref{S-spres}) exposes a dependence of multiscale parameters $\\zeta (q)$ and $\\alpha$ on the spatial resolution of the studied magnetograms. In contrast, the scale-free $D_0$ appears fairly insensitive to varying spatial resolution. Therefore, results and comparisons for $\\zeta (3)$ and $\\alpha$ in Section \\ref{S-comp} are valid only for MDI low-resolution data and should not be generalized to data sets of other instruments. Possible susceptibility of the $D_0$-value should also be studied with respect to the threshold it requires, unlike $\\zeta (3)$ and $\\alpha$. This investigation has not been carried out here. In previous works, however, \\inlinecite{Meunier_99} reported a decreasing trend of $D_0$ with increasing threshold, while \\inlinecite{Janssen_etal03} reported a slighter decrease, or a near insensitivity, of $D_0$ for increasing thresholds, in case these thresholds are sufficiently above noise levels or the magnetic field data have been treated for noise, respectively. It is useful to mention here a very recent result by \\inlinecite{Abramenko_Yurchyshyn10} that the turbulent power-spectrum index $\\alpha$, either alone or coupled with the integral of the power-spectrum for all wavenumbers, correlates better than $\\Phi_{\\mathrm{tot}}$ with the flaring index in a large sample of 217 active regions recorded in high-resolution MDI magnetograms. While correlating some parameter with the flaring index is not identical to inferring the predictive capability of this parameter, these results appear in likely contrast with the results presented here. Further investigation is clearly needed, therefore. Nonetheless, some convergence of views appears in that multiscale parameters may not be ideal tools for solar flare prediction (Abramenko, 2010, private communication). Perhaps more instructive than pointing out the inability of scale-free and multiscale techniques to assess {\\it a priori} the flaring record of active regions is to explain {\\it why} this is the case. In this author's view, there are at least two distinct reasons that justify our findings: First, fractality and multifractality are extremely widespread in the solar atmosphere, eruptive and quiescent alike. This may well be due to the turbulence dominating the magnetic-flux generation and emergence process (see Introduction). For example, recall the fractality of white-light granules \\cite{Roudier_Muller87,Hirzberger_etal97}, the fractality and multifractality of active regions and the quiet-Sun magnetic field \\cite{Schrijver_etal92,Cadavid_etal94,Meunier_99,Janssen_etal03}, the fractality of flares and sub-flares in the EUV (\\opencite{Aschwanden_Parnell02}; \\opencite{Aschwanden_Aschwanden08a}; \\citeyear{Aschwanden_Aschwanden08b}), the fractality of the quiet network in the EUV \\cite{Gallagher_etal98}, that of Ellerman bombs in off-band H$\\alpha$ \\cite{Georgoulis_etal02}, and others. The fractal dimension in most, if not all, of these works varies between 1.4 and 1.8, practically indistinguishable from the fractal dimension of active regions found here. As a result, it appears unlikely that these same methods may reflect particular characteristics of active regions, let alone flare productivity. Second, there is a lack of correlations between the fractal dimension in the photosphere and that of the overlying chromosphere and corona, where major flares occur. \\inlinecite{Dimitropoulou_etal09} assumed nonlinear force-free magnetic fields extending above the photosphere and calculated volumes of enhanced electric currents and steep magnetic gradients from these extrapolated fields. They found no correlation between the three-dimensional fractal dimension of these volumes and that of the two-dimensional photospheric boundary. In other words, all photospheric ``memory'', in terms of fractality and multifractality, is erased above the photosphere due to the fact that these unstable volumes become nearly space-filling slightly above this boundary. Attempting to assess the fractality of layers higher than the photosphere -- where flares occur -- by using the photospheric fractality as a proxy will not yield meaningful results, similarly to the lack of correlation between photospheric electric currents and the coronal X-ray brightness \\cite{Metcalf_etal94}. In addition, it is possible that both flaring and non-flaring regions share a similar degree of self-organization in the distribution of their magnetic free energy, as reported by \\inlinecite{Vlahos_Georgoulis04}. Flaring regions have an ``opportunity'' to show their self-organization via flaring, with flares inheriting the statistics of their host active regions, while non-flaring regions retain this property without demonstrating it. In this sense {\\it i)} fractality alone cannot be responsible for flaring, and {\\it ii)} fractality, as a global characteristic of the active-region atmosphere, cannot be used to determine {\\it a priori} which active regions will flare. There are, of course, sophisticated multiscale techniques not treated in this work, such as wavelet methods used to extract the magnetic-energy spectrum in active regions \\cite{Hewett_etal08} or to distinguish active regions from quiet Sun for further treatment \\cite{Conlon_etal10}, or the flatness function and its intermittency index \\cite{Abramenko_etal08}. While we cannot comment on methods that we have not tested, per our conclusions it would seem rather surprising if a scale-free or multiscale technique delivered a notable improvement in our forecasting ability, as this would apparently contradict what scale-free and multiscale behavior caused by self-organization is meant to imply: spontaneity in the system's dynamical response to external forcing, both in timing and in amplitude, and hence a lack of certainty in predicting this response. Let us finally mention that alternative flare prediction approaches have been developed in recent years. Rather than fractality, multifractality, or intermittency and turbulence, these methods rely on parameters stemming from morphological and topological characteristics of active regions, such as those of the photospheric magnetic-polarity inversion lines or photospheric properties in general (\\opencite{Falconer_etal06}; \\opencite{Schrijver_07}; \\opencite{Georgoulis_Rust07}; \\opencite{Leka_Barnes07}; \\opencite{Mason_Hoeksema10}), or those of the subsurface kinetic helicity prior to active-region emergence \\cite{Reinard_etal10}, among others. It remains to be seen whether these parameters can lead to advances in the forecasting of major solar eruptions or whether forecasting will remain inherently probabilistic which, per our results, seems entirely possible. In any case, fractal and multifractal methods -- perhaps not extremely useful as eruption predictors -- will always be excellent tools for a fundamental understanding of the origins and nature of solar magnetism. \\begin{acks} This work is based on a talk given by the author during the Fourth Solar Image Processing (SIP) Workshop in Baltimore, MD, USA, 26-30 October 2008. Thanks are due to the organizers for an interesting and productive meeting. During the author's tenure at the Johns Hopkins University Applied Physics Laboratory (JHU/APL) in Laurel, MD, USA, this work received partial support from NASA's LWS TR\\&T Grant NNG05GM47G and Guest Investigator Grant NNX08AJ10G. The author gratefully acknowledges the Institute of Space Applications and Remote Sensing (ISARS) of the National Observatory of Athens for the availability of their computing cluster facility for massive runs related to this work. SOHO is a project of international cooperation between ESA and NASA. {\\it Hinode} is a Japanese mission developed and launched by ISAS/JAXA, with NAOJ as domestic partner and NASA and STFC (UK) as international partners. It is operated by these agencies in co-operation with ESA and NSC (Norway). Finally, the author thanks the two anonymous referees for contributing to the clarity, accuracy, and focus of this work. \\end{acks}" }, "1101/1101.2068_arXiv.txt": { "abstract": "Optical $UBVRI$ photometry and medium resolution spectroscopy of the type Ib supernova SN 2009jf, during the period $\\sim -15$ to $+250$ days with respect to the $B$ maximum are reported. The light curves are broad, with an extremely slow decline. The early post-maximum decline rate in the $V$ band is similar to SN 2008D, however, the late phase decline rate is slower than other studied type Ib supernovae. With an absolute magnitude of $M_{V} = -17.96\\pm0.19$ magnitude at peak, SN 2009jf is a normally bright supernova. The peak bolometric luminosity and the energy deposition rate via $^{56}$Ni $\\rightarrow$ $^{56}$Co chain indicate that $\\sim {0.17}^{+0.03}_{-0.03}$ M$_{\\odot}$ of $^{56}$Ni was ejected during the explosion. He\\,I 5876~\\AA\\ line is clearly identified in the first spectrum of day $\\sim -15$, at a velocity of $\\sim 16000$ km sec$^{-1}$. The [O\\,I] 6300-6364~\\AA\\ line seen in the nebular spectrum has a multi-peaked and asymmetric emission profile, with the blue peak being stronger. The estimated flux in this line implies $\\ga 1.34$ M$_\\odot$ oxygen was ejected. The slow evolution of the light curves of SN 2009jf indicates the presence of a massive ejecta. The high expansion velocity in the early phase and broader emission lines during the nebular phase suggest it to be an explosion with a large kinetic energy. A simple qualitative estimate leads to the ejecta mass of M$_{\\rm ej} = 4-9$ M$_\\odot$, and kinetic energy E$_{\\rm K} = 3-8 \\times 10^{51}$ erg. The ejected mass estimate is indicative of an initial main-sequence mass of $\\ga 20- 25$ M$_\\odot$. ", "introduction": "Type Ib supernovae (SNe Ib) are core-collapse supernovae, characterized by the presence of prominent helium lines and the absence of hydrogen lines. They are believed to be the results of violent explosions of massive stars, such as the Wolf-Rayet stars, which are stripped of most or all of their hydrogen envelope, either by mass transfer to a companion (e.g., \\citealt{nomoto94}, \\citealt{pods04}), or via strong winds \\citep[e.g.,][]{woosley93}, or by sudden eruptions. These supernovae are also termed as stripped-envelope supernovae. The presence of hydrogen in type Ib supernovae remains an open issue for investigation. There are some type Ib events which show a deep absorption at $\\sim 6200$~\\AA\\ in their early spectra, which could be attributed to H$\\alpha$ (\\citealt{branch02}, \\citealt{anupama05}, \\citealt{soderberg08}), whereas some others show a shoulder in the red wing of the [O\\,I] 6300-6364~\\AA\\ line in their nebular spectra, due to H$\\alpha$ (\\citealt{sollerman98}, \\citealt {strit09}). Using the SYNOW code, \\cite{elmhamdi06} have shown the presence of a thin layer of hydrogen ejected at high velocity in almost all the SNe Ib in their sample. \\cite{maurer10} have recently investigated various mechanisms that can produce strong H$\\alpha$ emission in the late phase, and shown that it can be explained well by radioactive energy deposition, if hydrogen and helium are mixed in suitable fractions and clumped strongly. Late phase observations of SNe Ib have gained special importance as these phases probe deeper into the core of the expanding stars. The nebular spectrum originating from an optically thin ejecta provides important clues to the nature of progenitor star and the explosion mechanism. Asphericity in the explosion of stripped envelope supernovae is confirmed by a higher degree of polarization through spectropolarimetric studies of these objects during early phases (\\citealt{wang03}, \\citealt{leonard06}). Independent indications of the asphericity in the explosion come from the narrower width of [O\\,I] 6300-6364~\\AA\\ line compared to the [Fe\\,II] features at $\\sim 5200$~\\AA\\ (\\citealt{mazzali01}, \\citealt{maeda02}) and/or from the asymmetric profile of the [O\\,I] 6300-6364~\\AA\\ line \\citep{mazzali05}. SN 2009jf was discovered by \\cite{li09} in the Seyfert 2, barred spiral galaxy NGC 7479 on September 27.33. This supernova was classified as a young type Ib supernova by \\cite{kasliwal09}, and \\cite{sahu09a}, based on early spectra obtained on September 29. \\cite{itagaki09} reported the detection of a dim object at an unfiltered magnitude of $\\sim 18.2$ in an image obtained on 2006 November 8.499 and at a magnitude of $\\sim 18.3$ in an image obtained on 2007 August 13.74. They also report the presence of the object at $\\sim 18$ magnitude in the DSSS images. They have estimated the absolute magnitude of the object as $-14.5$ and suggested that these may be recurring outbursts of a luminous blue variable. In this paper we report optical photometry and spectroscopy of SN 2009jf in the early and nebular phase and discuss the results based on the observations. ", "conclusions": "The light curve and spectral evolution of SN 2009jf show some peculiarities compared to other SNe Ib. The light curves indicate a post-maximum decline that is slower compared to other type Ib supernovae. This slow decline continues even during the late phases, making the light curve of SN 2009jf broad. The absolute $V$ magnitude at peak is comparable to the mean of the absolute magnitude distribution of type Ib supernovae. Using the bolometric light curve and the energy deposition rate via $^{56}$Ni $\\rightarrow$ $^{56}$Co, the mass of $^{56}$Ni synthesized during the explosion is estimated to be 0.17 M$_\\odot$. SN 2009jf shows a very early emergence of He\\,I lines in the spectrum. He\\,I 5876~\\AA\\ line is identified in the first spectrum obtained 15.3 days before $B$ maximum. Other lines due to He\\,I at 4471~\\AA, and 6678~\\AA \\ were identified in the $-13$ day spectrum. Further, the expansion velocity estimated using He\\,I line $\\sim 16,000$ km sec$^{-1}$, indicating that helium is excited at high velocity. In case of SN 2008D, He\\,I lines became apparent around 11.5 days before $B$ maximum, and were prominent only around 5 days before $B$ maximum \\citep{modjaz09}. The He\\,I lines seen in the spectra of type Ib supernovae require non-thermal excitation and ionization, as the temperature present in the ejecta is too low to cause any significant absorption \\citep{lucy91}. $\\gamma$-rays, emitted by newly synthesized $^{56}$Ni during the explosion, accelerate electrons that act as a source of non-thermal excitation for He (\\citealt{harkenss87}, \\citealt{lucy91}). For exciting helium at such a high velocity as seen in SN 2009jf, the $\\gamma$-rays need to be close to the helium layer, which can be possible either through the escape of $\\gamma$-rays from the $^{56}$Ni dominated region, or through some large scale instability causing substantial mixing of $^{56}$Ni to the outer layers. The slower decline of the light curves of SN 2009jf gives an indication that it has a massive ejecta and the probability of $\\gamma$-rays escaping will be low. Though substantial mixing of different inner layers appears to be the most probable way for an early excitation of He at high velocities, the possibilty of some $\\gamma$-rays reaching the He layer and exciting it cannot be ruled out, especially since SN 2009jf is a rather luminous supernova. The profile of [O\\,I] 6300-6364~\\AA\\ feature in the nebular spectrum is multi-peaked and asymmetric with a sharp, stronger blue peak. The peak of this feature is blueshifted by $\\sim 30$~\\AA\\ around day +86, which reduces to a blueshift of $\\sim 15$~\\AA\\ by day +99. Such observed blueshifts are explained as a result of residual opacity in the core of the ejecta \\citep{taubenberger09}. The asymmetric and multi-peaked profile seen at phases later than 200 days can be produced by additional components of arbitrary width and shift with respect to the main component. Such profiles are indicative of an ejecta with large-scale clumping, a single massive blob, or a unipolar jet. The asymmetric [O\\,I] line profile of SN 2009jf with a stronger blue peak is very similar to the line profiles of SNe 2000ew and 2004gt. \\cite{taubenberger09} have explored a possible configuration which can give rise to this asymmetric line profile, and interperted the profile as originating from the deblended 6300~\\AA\\ and 6364~\\AA\\ lines of a single narrow, blueshifted component. \\cite{maurer10} have shown that the profile of [O\\,I] 6300-6364~\\AA\\ doublet is likely to be influenced by H$\\alpha$ absorption. If hydrogen concentration is located around $\\sim$ 12000 km sec$^{-1}$, it causes a split of the [O\\,I] 6300-6364~\\AA\\ doublet, leading to a double-peaked oxygen profile. Neither scenarios account for the stronger blue peak of the 6300~\\AA\\ line. \\cite{taubenberger09} explain the stronger blue peak with a complex ejecta structure with additional blueshifted emission on top of an otherwise symmetric profile. Alternatively, the asymmetry in the profile is explained by a damping of the redshifted emission component in an originally toroidal distribution, caused by the optically thick inner ejecta. The light curve evolution of SN 2009jf indicates the presence of an ejecta more massive than other stripped core collapse supernovae. Hence, it is quite likely that in the case of SN 2009jf also the redward component is damped by an optically thick inner ejecta. It should however be noted that asymmetric and multi-component profiles cannot be reproduced within spherical symmetry \\citep{mazzali05, maeda07}. This needs further investigation with observations at phases later than presented here, as well as spectrophotometric observations and detailed modelling. The brightness and width of Type Ib light curves are determined by the interplay of nickel mass, opacity and $\\gamma$-ray deposition. In general, a greater amount of $^{56}$Ni will make the light curve brighter. A more massive ejecta will have a larger optical depth, and it will take longer for the trapped decay energy to diffuse through the envelope, which will broaden the light curve \\citep{ensman88}. The time taken for the bolometric light curve to decline from peak to the moment when the luminosity is equal to $1/e$ times the peak luminosity (which is equivalent to a decline of 1.1 mag from peak), is known as the effective diffusion time $\\tau_{\\rm m}$. The effective diffusion time is related to the mass of the ejecta M$_{\\rm ej}$ and the kinetic energy E$_{\\rm k}$ of the ejecta $\\tau_{\\rm m} \\propto \\kappa_{\\rm opt}^{1/2} M_{\\rm ej}^{3/4} K_{\\rm E}^{-1/4}$ \\citep{arnett82}, where $\\kappa_{\\rm opt}$ is optical opacity. The expansion velocity $v$ can be expressed in terms of M$_{\\rm ej}$ and E$_{\\rm k}$ as $v \\propto M_{\\rm ej}^{-1/2} K_{\\rm E}^{1/2}$. The broad peak and slower decline rates of the light curves of SN 2009jf in comparison to other supernovae indicate that SN 2009jf has a massive ejecta. Further, the broader emission lines at late phase indicates a larger explosion energy. There are several core-collapse supernovae for which the progenitor mass has been constrained using hydrodynamical modelling. With this approach, \\cite{nomoto03} and \\cite{nomoto04} constructed the $E_{\\rm K} - M_{\\rm MS}$ diagram and introduced a hypernova branch. Recent updates of this approach include SN 1998bw \\citep{maeda06}, SN 2008D \\citep{tanaka09}, and SN 2003bg \\citep{mazzali09}. For the well studied bright hypernova SN 1998bw ($M_V=-19.35$ \\cite{galama98}), the main sequence mass of the progenitor is constrained by \\cite{maeda06} as $\\sim$ 40 M$_\\odot$. Though SN 2009jf has a brighter peak compared to SN 2008D, the fact that the light curves of SN 2009jf around maximum and the initial decline rate $\\Delta m_{15}(V)$ are similar to those of SN 2008D can be used to estimate the mass of the ejecta M$_{ej}$ and kinetic energy E$_{\\rm k}$ of the ejecta, using SN 2008D as the reference, assuming the optical opacity $\\kappa_{\\rm opt}$ to be the same. The effective diffusion time $\\tau_{\\rm m}$ for SN 2009jf and SN 2008D are estimated to be 30 days and 26 days, respectively. The photospheric expansion velocity estimated using the Fe\\,II 5169~\\AA\\ line at maximum is $\\sim$ 10000 km sec$^{-1}$, similar for both the objects. For SN 2008D, the mass of the ejecta M$_{\\rm ej}$, the kinetic energy E$_{\\rm k}$ and the progenitor mass have been derived by \\cite{mazzali08}, \\cite{soderberg08} and \\cite{tanaka09}. \\cite{mazzali08} could reproduce the spectral evolution and light curve with a spherically symmetric explosion energy E$_{\\rm k} = 6.0\\times 10^{51}$ erg and ejected mass M$_{\\rm ej} \\sim$ 7 M$_\\odot$ with a progenitor of mass $\\sim$ 30 M$_\\odot$ while \\cite{soderberg08} have arrived at E$_{\\rm k} = 2-4 \\times 10^{51}$ erg and M$_{\\rm ej}$ = $3-5$ M$_\\odot$, by applying rescaling arguments. \\cite{tanaka09} have calculated the hydrodynamics of explosion and explosive nucleosynthesis for SN 2008D with varying mass for the He core of the progenitor and concluded that the progenitor star of SN 2008D had a He core mass $6-8$ M$_\\odot$ prior to explosion. This corresponds to a main sequence mass of $M_{\\rm MS} = 20-25$ M$_\\odot$. The explosion energy and mass of ejecta for SN 2008D were estimated to be E$_{\\rm k}=6.0\\pm 2.5\\times10^{51}$ erg and M$_{\\rm ej} = 5.3\\pm 1.0$ M$_\\odot$, respectively. Thus, for SN 2008D the mass of ejecta M$_{\\rm ej}$ and explosion energy E$_{\\rm k}$ range between $3-7$ M$_\\odot$ and $2-6 \\times 10^{51}$ erg, respectively. Using the observed photospheric velocity and the estimated diffusion time for SN 2009jf, and treating SN 2008D as a reference, we estimate M$_{\\rm ej} = 4-9$ M$_\\odot$ and E$_{\\rm k} = 3-8\\times 10^{51}$ erg for SN 2009jf. This indicates that SN 2009jf was an energetic explosion of a star having a mass similar, or somewhat more massive than the progenitor of SN 2008D ($M_{\\rm MS} \\ga 20-25$ M$_\\odot$). The physical parameters of SN 2009jf may also be compared with those of the type IIb supernova SN 2003bg, which had an absolute peak magnitude of $M_V=-17.5$ \\citep{hamuy09} and an oxygen mass estimate of 1.3 M$_\\odot$. \\cite{mazzali09} have estimated the physical parameters for SN 2003bg based on detailed light curve and spectral modelling. The best fit model gives an ejected mass of $\\sim 4.8$ M$_\\odot$, kinetic energy $\\sim 5\\times 10^{51}$ erg and mass of $^{56}$Ni $\\sim 0.15-0.17$ M$_\\odot$. The mass of the progenitor for SN 2003bg is estimated as $20-25$ M$_\\odot$. Our qualitative analysis of light curve and spectra of SN 2009jf hints towards a higher kinetic energy and a slightly more massive ejecta than SN 2003bg, and in turn a progenitor with $M_{\\rm MS} \\ga 20-25$ M$_\\odot$. The mass of oxygen $M$(O) in the ejecta of the core collapse SNe is very sensitive to the main-sequence mass $M_{\\rm MS}$ of the progenitor. For $M_{\\rm MS}$ = 15, 18, 20, 25, 30, and 40 M$_\\odot$, $M$(O) = 0.16, 0.77, 1.05, 2.35, 3.22, and 7.33 M$_\\odot$, respectively \\citep{nomoto06}. These values are obtained for E$_{\\rm k}=1.0\\times10^{51}$ erg and the metallicity $z=0.02$, but are not so sensitive to E$_{\\rm k}$ and $z$. In fact, for ($M_{\\rm MS}$/M$_\\odot$, E$_{\\rm k}/10^{51}$ erg) = (20, 10), (25, 10), and (30, 20), and (40, 30), $M$(O)/M$_\\odot$ = 0.98, 2.18, 2.74, and 7.05, respectively \\citep{nomoto06}. Therefore, the lower limit of the oxygen mass $M$(O) $\\ga$ 1.34 M$_\\odot$ estimated from the nebular spectra is quite consistent with the progenitor mass of $M_{\\rm MS} \\ga 20-25$ M$_\\odot$ estimated from the light curve shape and the photospheric velocities. The [Ca\\,II] 7291-7324/[O\\,I]6300-6364 line ratio is also a good diagnostic of $M_{\\rm MS}$, because the mass of the explosively synthesized Ca in the ejecta, $M$(Ca), is not sensitive to $M_{\\rm MS}$. For $M_{\\rm MS}$/M$_\\odot$ = 15, 18, 20, 25, 30, and 40, $M$(Ca)/10$^{-2}M_\\odot$ = 0.40, 0.45, 0.37, 0.66, 1.6, and 1.6, respectively \\citep{nomoto06}. Also, for ($M_{\\rm MS}$/M$_\\odot$, E$_{\\rm k}/10^{51}$ erg) = (20, 10), (25, 10), and (30, 20), and (40, 30), $M$(Ca)/10$^{-2}M_\\odot$ = 0.50, 0.57, 0.93, and 1.4, respectively \\citep{nomoto06}. This is in contrast to $M$(O), which sensitively increases with $M_{\\rm MS}$. Thus a smaller [Ca\\,II]/[O\\,I] ratio indicates a massive core. The [Ca\\,II] 7291-7324/[O\\,I]6300-6364 emission line ratio for SN 2009jf is estimated as 0.51 and 0.49 using the nebular spectrum observed on days $+229$ and $+251$, respectively. For SN 2007Y, SN 1996N, SN 1990I and SN 1998bw, this ratio was found to be 1.0, 0.9, 0.7 and 0.5, respectively, at similar epochs (\\citealt{elmhamdi04}, \\citealt{strit09} and references therein). \\cite{fransson89} have theoretically calculated the [Ca\\,II]/[O\\,I] line ratio for progenitor masses of 15 and 25 M$_{\\odot}$. The observed [Ca\\,II]/[O\\,I] ratio for SN 2009jf is very close to the ratio expected for a star with $M_{\\rm MS}$= 25 M$_\\odot$, as indicated by Model 1b in \\cite{fransson89}. The estimates of the mass of $^{56}$Ni synthesized during the explosion, the kinetic energy of explosion and the main sequence mass of the progenitor star places SN 2009jf between the normal core-collapse supernovae and the hypernovae branch in the $E_K - M_{MS}$ diagram of \\cite{tanaka09}, at the upper end of the normal core-collapse supernovae branch. It is however to be noted that the [O\\,I] line profile during the nebular phase indicates asymmetry of the explosion. This can have some effect in the kinetic energy estimate, as shown by \\cite {maeda06} and \\cite{tanaka07} for SN 1998bw. A detailed modelling is therefore required for a better estimate of the various parameters. \\cite{itagaki09} suggest the progenitor could have undergone luminous blue variable type mass loss events, based on their detection of a dim object at the location of the supernova on three occasions. Pre-supernova images of the host galaxy obtained in the ultraviolet by the {\\it Swift} satellite, and available in the {\\it Swift} data archives, clearly indicate the presence of a bright HII region at the supernova location. It is hence quite likely that the object detected by \\cite{itagaki09} corresponds to the underlying HII region." }, "1101/1101.2632_arXiv.txt": { "abstract": "{We present the results of the observations of the $(J,K)=(1,1)$ and the $(J,K)=(2,2)$ inversion transitions of the \\nh\\ molecule toward a large sample of 40 regions with molecular or optical outflows, using the 37 m radio telescope of the Haystack Observatory. We detected \\nh\\ emission in 27 of the observed regions, which we mapped in 25 of them. Additionally, we searched for the $6_{16}-5_{23}$ \\hho\\ maser line toward six regions, detecting \\hho\\ maser emission in two of them, HH265 and AFGL 5173. We estimate the physical parameters of the regions mapped in \\nh\\ and analyze for each particular region the distribution of high density gas and its relationship with the presence of young stellar objects. In particular, we identify the deflecting high-density clump of the HH270/110 jet. We were able to separate the \\nh\\ emission from the L1641-S3 region into two overlapping clouds, one with signs of strong perturbation, probably associated with the driving source of the CO outflow, and a second, unperturbed clump, which is probably not associated with star formation. We systematically found that the position of the best candidate for the exciting source of the molecular outflow in each region is very close to an \\nh\\ emission peak. From the global analysis of our data we find that in general the highest values of the line width are obtained for the regions with the highest values of mass and kinetic temperature. We also found a correlation between the nonthermal line width and the bolometric luminosity of the sources, and between the mass of the core and the bolometric luminosity. We confirm with a larger sample of regions the conclusion of Anglada et al.\\ (1997) that the \\nh\\ line emission is more intense toward molecular outflow sources than toward sources with optical outflow, suggesting a possible evolutionary scheme in which young stellar objects associated with molecular outflows progressively lose their neighboring high-density gas, weakening both the \\nh\\ emission and the molecular outflow in the process, and making optical jets more easily detectable as the total amount of gas decreases.} ", "introduction": "Over the last decades, a great effort has been made to study the processes that take place in the earliest stages of stellar evolution. It is now widely accepted that low-mass stars begin their lives in the densest cores of molecular clouds and that the earliest stages of stellar evolution are associated with processes involving a strong mass loss, traced by molecular outflows, Herbig-Haro objects, and optical jets, which emanate from the deeply embedded young stellar objects. These mass-loss processes have been proposed as a way to eliminate the excess of material and of angular momentum as well as to regulate the IMF (Shu, Adams \\& Lizano 1987). The molecular outflow phase is known as one of the earliest observable phases of the stellar evolution. Several studies indicate that most, if not all, of the Class 0 and Class I sources drive molecular outflows (e.g. Davis et al.\\ 1999) and that an important fraction of these sources are associated with Herbig-Haro objects, as well as with molecular outflows (Eiroa et al.\\ 1994; Persi et al.\\ 1994). These results suggest that both phenomena start and coexist in the early stages of the star-formation process. These outflows emanating from protostars collide with the remaining molecular cloud and disperse the surrounding material, and determine the evolution of the dense core where the star is born (Arce \\& Sargent 2006). In this sense, the study of these mass-loss processes and the molecular environment of the embedded objects from which they emanate has became an important tool in order to better understand the earliest stages of stellar evolution. Ammonia observations have proved to be a powerful tool for studying the dense cores where the stars are born. Since the first surveys of dense cores (Torrelles et al.\\ 1983, Benson \\& Myers 1989, Anglada et al.\\ 1989), a clear link was established between dense cores, star formation and outflows (see e.g., the review of Andr\\'e et al.\\ 2000). From these surveys, it was clearly established that the driving sources of outflows are usually embedded in the high-density gas, which is traced by the \\nh\\ emission, and are located very close to the emission peak (Torrelles et al.\\ 1983; Anglada et al.\\ 1989). Following these results, Anglada et al.\\ (1997) (hereafter Paper I) undertook a survey of dense cores to investigate the relationship between the type of outflow and the dense gas associated with their exciting sources. A statistical study of the sources observed in that survey reveals that the ammonia emission is more intense toward molecular outflow sources than toward sources with only optical outflows, indicating that molecular outflows are associated with a larger amount of high-density gas. From this result, a possible evolutionary scheme was suggested in which young objects associated with molecular outflows progressively lose their neighboring high-density gas, while both the \\nh\\ emission and the molecular outflow become weaker in the process, and the optical jets become more easily detectable as the total amount of gas and extinction decreases. In this sense, the observations of high-density tracers, such as the \\nh\\ molecule, confirm the decrease of high-density gas around the stars. We present here new ammonia observations. Additional regions allow us to obtain a sample of outflow regions observed in ammonia that doubles the number used in Paper I. We selected a sample of 40 star-forming regions, taking into account the presence of molecular outflows, optical outflows, or both, and mapped with the Haystack 37 m telescope the \\nh\\ emission around the suspected outflow exciting sources. In Sect.~\\ref{obs4} we describe the observational procedure, in Sect.~\\ref{sources4} we present the observational results (the discussion of individual sources is presented in Appendix A), in Sect.~\\ref{results} we discuss the global results, in Sect.~\\ref{evol} we describe the relationship between the high-density gas and the nature of the outflow based on the sample, and in Sect.~\\ref{conclus4} we give our conclusions. ", "conclusions": "We detected the ammonia emission in 27 sources of a sample of 40 sources associated with molecular and/or optical outflows, and we were able to map 25 of them. We also searched for \\hho\\ maser emission toward 6 sources, and detected new \\hho\\ masers in HH 265 and AFGL 5173. Our main conclusions can be summarized as follows: \\begin{enumerate} \\item In all molecular outflow regions mapped, the \\nh\\ emission peak falls very close to the position of a very good candidate for the outflow excitation (except in the case of IRAS 05490+2658 where we propose an alternative location for the exciting source). On the other hand, the sources associated with optical outflow (except HH 270 IRS and HH 290 IRS) are not associated with an ammonia emission peak. \\item Four regions (HH 265, L588, L673 and IRAS 20050+2720) could be in the process of gravitational collapse at scales $\\geq0.2$ pc, as their derived masses exceed the virial mass by a factor $>5$. The rest of ammonia condensations appear to be close to the virial equilibrium. \\item In several regions the ammonia structure presents more than one clump. In the M120.1+3.0-N and IRAS 20050+2720 regions, two clumps with different velocities, which are gravitationally bound though were identified. \\item We identified a high-density clump where the HH 270/110 jet can suffer the collision responsible for the deflection of the jet. \\item We were able to separate the \\nh\\ emission from the L1641-S3 region into two overlapping clouds, one with signs of strong perturbation, probably associated with the driving source of the CO outflow, and a second, quiescent clump, which probably is not associated with star formation. \\item In general, the observed \\nh\\ condensations are very cold, with line widths dominated by nonthermal (turbulent) motions. Among the observed sources, the more massive regions appear to produce a larger perturbation in their molecular high density environment. \\item We found that generally the more luminous objects are associated with broader ammonia lines. A correlation between the nonthermal component of the line width and the luminosity of the associated object, $\\log(L_{\\rm bol}/L_{\\sun})=(3.6\\pm0.9)\\,\\log(\\Delta V_{\\rm nth}/\\rm km~s^{-1})+(1.8\\pm0.2)$ was found for sources with $D\\leq$ 1 kpc. \\item We found that there is a significant correlation between the luminosity of the source and the mass of the core and that this correlation is not caused by any distance bias in the sample. Both parameters are related by $\\log(L_{\\rm bol}/L_{\\sun})=(0.8\\pm0.2)\\,\\log(M_{\\rm core}/M_{\\sun})+(0.2\\pm0.3)$. \\item The ammonia brightness temperature and column density of the sources decrease as the outflow activity becomes prominent in the optical. These results give an evolutive scheme in which young objects progressively lose their surrounding high-density gas. The ammonia emission and the observational appearance of outflows trace an evolutive sequence of sources. \\end{enumerate}" }, "1101/1101.2626_arXiv.txt": { "abstract": "{ We present $^{12}$CO(1--0) and $^{12}$CO(2--1) maps of the interacting barred LINER/Seyfert 2 galaxy \\nnn\\ obtained with the IRAM interferometer at resolutions of 2\\farcs1 $\\times$ 1\\farcs3 and 0\\farcs9 $\\times$ 0\\farcs6, respectively. We also present single-dish IRAM 30\\,m $^{12}$CO(1--0) and $^{12}$CO(2--1) observations used to compute short spacings and complete interferometric measurements. These observations are complemented by IRAM 30\\,m measurements of HCN(1--0) emission detected in the center of \\nnn. The molecular gas emission shows a nuclear peak, an elongated bar-like structure of $\\sim$18\\arcsec\\ ($\\sim$900\\,pc) diameter in both $^{12}$CO maps and, in $^{12}$CO(1--0), a two-arm spiral feature from $r$$\\sim$9\\arcsec\\ ($\\sim$450\\,pc) to $r\\sim$16\\arcsec\\ ($\\sim$800\\,pc). The inner $\\sim$18\\arcsec\\ bar-like structure, with a north/south orientation (PA = 14$^{\\circ}$), forms two peaks at the extremes of this elongated emission region. The kinematics of the inner molecular gas shows signatures of non-circular motions associated both with the 18\\arcsec\\ bar-like structure and the spiral feature detected beyond it. The 1.6\\,$\\mu$m $H$-band 2MASS image of \\nnn\\ shows a stellar bar with a PA = $-21^{\\circ}$, different from the PA (= $14^{\\circ}$) of the $^{12}$CO bar-like structure, indicating that the gas is leading the stellar bar. The far-infrared \\spitzer-MIPS 70 and 160\\,$\\mu$m images of \\nnn\\ show that the dust emission is intensified at the nucleus and at the ansae at the ends of the bar, coinciding with the $^{12}$CO peaks. The \\textit{GALEX} far-ultraviolet (FUV) morphology of \\nnn\\ displays an inner elongated (north/south) ring delimiting a hole around the nucleus, and the $^{12}$CO bar-like structure is contained in the hole observed in the FUV. The torques computed with the \\hst-NICMOS F160W image and our PdBI maps are negative down to the resolution limit of our images, $\\sim$60 pc in $^{12}$CO(2--1). If the bar ends at $\\sim$3\\,kpc, coincident with corotation (CR), the torques are negative between the CR of the bar and the nucleus, down to the resolution limit of our observations. This scenario is compatible with a recently-formed rapidly rotating bar which has had insufficient time to slow down because of secular evolution, and thus has not yet formed an inner Lindblad resonance (ILR). The presence of molecular gas inside the CR of the primary bar, where we expect that the ILR will form, makes \\nnn\\ a potential \\textit{smoking gun} of inner gas inflow. The gas is fueling the central region, and in a second step could fuel directly the active nucleus. ", "introduction": "The Nuclei of Galaxies (NUGA) project \\citep[][]{santi03} is an IRAM Plateau de Bure Interferometer (PdBI) and 30\\,m single-dish survey of nearby low-luminosity active galactic nuclei (LLAGN). The aim is to map, at high resolution ($\\sim$0\\farcs5-2\\arcsec) and high sensitivity ($\\sim$2-4\\,mJy\\,beam$^{-1}$), the distribution and dynamics of the molecular gas in the inner kpc of the galaxies of our sample, and to study the different mechanisms for gas fueling of LLAGN. NUGA galaxies analyzed so far show that there is no unique circumnuclear molecular gas feature linked with nuclear activity, but rather a variety of molecular gas morphologies which characterize the inner kpc of active galaxies. We have found one- and two-armed instabilities \\citep{santi03}, well-ordered rings and nuclear spirals \\citep{francoise04,vivi08a}, circumnuclear asymmetries \\citep{melanie05}, large-scale bars \\citep{fred07,leslie08}, and smooth disks \\citep{vivi10}. Among these morphologies, analyzing the torques exerted by the stellar gravitational potential on the molecular gas shows that only four NUGA galaxies: NGC\\,6574 \\citep{lindt-krieg08}, NGC\\,2782 \\citep[][]{leslie08}, NGC\\,3147 \\citep{vivi08a}, and NGC\\,4579 \\citep[][]{santi09} show evidence for gas inflow. These galaxies have several features in common: (1)~a large circumnuclear mass concentration (i.e., a dominant stellar bulge); (2)~a high circumnuclear molecular gas fraction ($\\ga$10\\%); and (3)~kinematically decoupled bars with overlapping dynamical resonances. The large amount of gas around the nucleus, combined with dynamical features that enable the gas to penetrate the inner Lindblad Resonance (ILR), seem to be necessary (and perhaps sufficient) ingredients for inducing gas inflow in circumnuclear scales. The existence of different nuclear molecular morphologies can be sought in the variety of timescales characterizing nuclear activity. Strong fueling only lasts for a time of $t_{\\rm fuel}\\sim0.002\\times t_{\\rm H}$, where $t_{\\rm H}\\sim1.4\\times10^{10}$\\,yr is the age of the Universe \\citep{heckman04}. Thus, the total time during which strong fueling can occur is around $t_{\\rm fuel}\\sim3\\times10^{7}$\\,yr; if there are $N$ fueling events per black hole per Hubble time, each event would have a duration of $t_{\\rm event}\\sim3\\times10^{7}/N$\\,yr. This implies that the strong accretion phase is a fraction $\\simeq 0.3/N$ of the characteristic galaxy dynamical time ($\\sim10^{8}$\\,yr). Although large-scale bars can produce gas inflow \\citep[e.g.,][]{francoise85,sakamoto99} and in some cases also drive powerful starbursts \\citep[e.g.,][]{knapen02,jogee05}, a correlation between large-scale bars and nuclear activity has not yet been verified \\citep[e.g.,][]{mulchaey97}. This lack of correlation is probably related to the different timescales for bar-induced gas inflow \\citep[$\\gtrsim$300 Myr,][]{jogee05}, AGN duty cycles ($\\sim$$10^{7}$\\,yr), and intermittent active accretion every $\\sim$$10^{8}$\\,yr \\citep[][]{ferrarese01,marecki03,janiuk04,hopkins06,king07}. The comparison of these different timescales suggests that most AGN are in an intermediate phase between active accretion episodes making the detection of galaxies with nuclear accretion somewhat difficult. Gravitational torques act on timescales of $\\sim$10$^{6-7}$\\,yr and are the most efficient mechanism in driving the gas from large spatial scales (some tens of kpc) to intermediate spatial scales (a few hundreds of pc). Dynamical friction and viscous torques are often invoked, in addition to gravitational torques, as possible mechanisms of AGN fueling. However, dynamical friction of giant molecular clouds in the stellar bulge of a galaxy tends to be a slow, inefficient process which, to first approximation, can be neglected relative to gravity torques \\citep{santi05}. Viscous torques can be more effective, and are favored in the presence of large density gradients and high galactic shear \\cite[see][for details]{santi05}. Nevertheless, they are relatively inefficient when there are strong (positive) gravity torques. This paper is dedicated to the galaxy \\nnn, the eleventh object of the core NUGA sample studied on a case-by-case basis. \\nnn\\ (Messier 66, $D$ = 10.2\\,Mpc, $H_{0}$ = 73 km s$^{-1}$ Mpc$^{-1}$) is an interacting \\citep[e.g.,][]{vivi04} and barred galaxy classified as SAB(s)b showing signatures of a LINER/Seyfert 2 type nuclear activity \\citep{ho97}. With NGC~3623 and NGC~3628, it forms the well-known Leo Triplet (M\\,66 Group, VV\\,308). Since the discovery of a long plume in H{\\scshape i} extending about 50$^{\\prime}$ to the east of NGC~3628 \\citep{zwicky56,haynes79}, evidence of past interactions between \\nnn\\ and NGC~3628 (the two largest spirals in the group), the Leo Triplet has been extensively studied from the radio to the optical, and in X-ray bandpasses. Optical broad-band images of \\nnn\\ reveal a pronounced and asymmetric spiral pattern with heavy dust lanes, indicating strong density wave action \\citep{ptak06}. While the western arm is accompanied by weak traces of star formation (SF) visible in H$\\alpha$, the eastern arm contains a star-forming segment in its inner part \\citep{smith94}. \\nnn\\ also possesses X-ray properties of a galaxy with a recent starburst \\citep{dahlem96}. Both the radio continuum (2.8\\,cm and 20\\,cm) and the $^{12}$CO(1--0) emissions show a nuclear peak, extend along the leading edges of the bar forming two broad maxima at the bar ends, and then the spiral arms trail off from the bar ends \\citep{haan08,paladino08,haan09}. On the contrary, the H{\\scshape i} emission exhibits a spiral morphology without signatures of a bar in the atomic gas \\citep{haan08,walter08,haan09}. \\begin{table} \\caption[]{Fundamental parameters for \\nnn.} \\begin{center} \\begin{tabular}{lll} \\hline \\hline Parameter & Value$^{\\mathrm{b}}$ & Reference$^{\\mathrm{c}}$ \\\\ \\hline $\\alpha_{\\rm J2000}$$^{\\mathrm{a}}$ & 11$^h$20$^m$15.02$^s$ & (1) \\\\ $\\delta_{\\rm J2000}$$^{\\mathrm{a}}$ & 12$^{\\circ}$59$^{\\prime}$29\\farcs50 & (1) \\\\ $V_{\\rm hel}$ & 744 km\\,s$^{-1}$ & (1) \\\\ RC3 Type & SAB(s)b & (2) \\\\ Nuclear Activity & LINER/Seyfert 2 & (3) \\\\ Inclination & 61\\fdg3 & (1) \\\\ Position Angle & 178$^{\\circ}$ $\\pm$ 1$^{\\circ}$ & (1) \\\\ Distance & 10.2\\,Mpc (1\\arcsec\\ = 49\\,pc) & (2) \\\\ L$_{B}$ & $4.2 \\times 10^{10}$\\,L$_{\\odot}$ & (4) \\\\ M$_{\\rm H\\,I}$ & $8.1 \\times 10^{8}$\\,M$_{\\odot}$ & (5) \\\\ M$_{\\rm H_{2}}$ & $4.1 \\times 10^{9}$\\,M$_{\\odot}$ & (6) \\\\ M$_{\\rm dust}$(60 and 100\\,$\\mu$m)& $4.5 \\times 10^{6}$\\,M$_{\\odot}$ & (4) \\\\ L$_{\\rm FIR}$ & $1.2 \\times 10^{10}$\\,L$_{\\odot}$ & (7) \\\\ \\hline \\hline \\end{tabular} \\label{table1} \\end{center} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] ($\\alpha_{\\rm J2000}$, $\\delta_{\\rm J2000}$) is the phase tracking center of our $^{12}$CO interferometric observations, assumed coincident with the dynamical center of \\nnn\\ (see Sect. \\ref{sec:dyncen}). \\item[$^{\\mathrm{b}}$] Luminosity and mass values extracted from the literature have been scaled to the distance of $D$ = 10.2\\,Mpc. \\item[$^{\\mathrm{c}}$] (1) This paper; (2) NASA/IPAC Extragalactic Database (NED, http://nedwww.ipac.caltech.edu/); (3) \\citet{ho97}; (4) \\citet{vivi04}; (5) \\citet{haan08}; (6) \\citet{kuno07}; (7) {\\it IRAS} Catalog. \\end{list} \\end{table} \\begin{table*} \\caption[]{$^{12}$CO(1--0) flux values, both obtained by our observations and extracted from the literature, for \\nnn.} \\begin{center} \\begin{tabular}{llllll} \\hline \\hline Reference & Telescope &\tDiameter & Primary beam\tor FOV$^{\\mathrm{a}}$ & Beam & Flux \\\\ & & [m] &\t[\\arcsec] & [\\arcsec $\\times$ \\arcsec] &\t[Jy km\\,s$^{-1}$] \\\\ \\hline \\citet{young95}\t & FCRAO\t & 14 & 45 &\t\t & 786\t\\\\ This paper\t & PdBI+30\\,m &\t& 42 & 2.1 $\\times$ 1.3\t& 668\t\\\\ This paper\t & PdBI+30\\,m & \t& 22$^{\\mathrm{b}}$ & 2.1 $\\times$ 1.3\t& 359 \\\\ This paper\t & PdBI\t &\t& 22$^{\\mathrm{b}}$ & 2.0 $\\times$ 1.3\t& 251\t\\\\ This paper\t& 30\\,m\t & 30\t& 22 (central position) & & 343$^{\\mathrm{c}}$\t\\\\ \\citet{helfer03} & NRAO & 12 & 55 (inner 50\\arcsec$\\times$50\\arcsec) & & 1100--1200\t\\\\ This paper & 30\\,m\t & 30\t& 22 (inner 50\\arcsec$\\times$50\\arcsec) & &1097$^{\\mathrm{d}}$ \\\\ \\hline \\hline \\end{tabular} \\label{table2} \\end{center} \\begin{list}{}{} \\item[$^{\\mathrm{a}}$] Primary beam is considered for single-dish observations, while field-of-view (FOV) for interferometric or combined (interferometric+single-dish) ones. \\item[$^{\\mathrm{b}}$] The photometry has been performed within 22\\arcsec, the $^{12}$CO(1--0) primary beam for the 30\\,m telescope. \\item[$^{\\mathrm{c}}$] The $^{12}$CO(1--0) recovered flux for the central position (0\\arcsec, 0\\arcsec). \\item[$^{\\mathrm{d}}$] The $^{12}$CO(1--0) recovered flux for inner $\\sim$50\\arcsec$\\times$50\\arcsec, $5 \\times 5$ mapping with 7\\arcsec spacing (see Sect. \\ref{sec:30m-obs}). \\end{list} \\end{table*} The most recent \\hi\\ mass determination for \\nnn\\ has been obtained by \\citet[][]{haan08}, M$_{\\rm H\\,I}=8.1\\times$10$^{8}$\\,M$_{\\odot}$ (reported to our adopted distance of $D=10.2$\\,Mpc), on average less than the typical value expected for interacting galaxies of the same Hubble type \\citep[][]{vivi04}. The H$_{2}$ mass content estimated by \\citet{kuno07} is $4.1\\times10^{9}\\,$M$_{\\odot}$ (scaled to our distance of $D=10.2$\\,Mpc for \\nnn). These H$_{2}$ and \\hi\\ mass values give a H$_{2}$/\\hi\\ mass ratio of 5.1, high compared to the average ratio expected for galaxies similar to \\nnn, M$_{\\rm H_{2}}/$M$_{\\rm HI} = 0.9$ \\citep[][]{vivi04}. The high H$_{2}$/\\hi\\ mass ratio in \\nnn\\ is probably due to the tidal interaction with NGC\\,3628, since this galaxy has ``captured'' much of the \\hi\\ in \\nnn\\ \\citep{zhang93}. Other molecular transitions have been detected in \\nnn, including HCN(1--0), HCN(2--1), HCN(3--2), HCO$^{+}$(1--0), and HCO$^{+}$(3--2), suggesting the presence of high density gas \\citep[][]{gao04,melanie08}. We list in Table \\ref{table1} the main observational parameters of \\nnn. The structure of this paper is as follows. In Sect. \\ref{sec:obs}, we describe our new observations of \\nnn\\ and the literature data with which we compare them. In Sects. \\ref{sec:30m} and \\ref{sec:pdbi}, we present the observational results, both single-dish and interferometric, describing morphology, excitation conditions, and kinematics of the molecular gas in the inner kpc of \\nnn. Comparisons between $^{12}$CO observations and those obtained at other wavelengths are given in Sect. \\ref{sec:comparison}. In Sect. \\ref{sec:torques}, we describe the computation of the gravity torques derived from the stellar potential in the inner region of \\nnn, and in Sect. \\ref{sec:interpretation}, we give an dynamical interpretation of the results. Finally, Sect. \\ref{sec:conclusions} summarizes our main results. We assume a distance to \\nnn\\ of $D=10.2$\\,Mpc, (HyperLeda DataBase\\footnote{\\citet[][]{paturel03}, http://leda.univ-lyon1.fr}) and a Hubble constant $H_0=73$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. This distance means that 1\\arcsec\\ corresponds to 49\\,pc. \\begin{figure*} \\centering \\begin{tabular}{c} \\includegraphics[height=0.8\\textwidth,angle=-90]{fg1a.ps}\\\\ \\includegraphics[height=0.8\\textwidth,angle=-90]{fg1b.ps}\\\\ \\end{tabular} \\caption{ Spectra maps of \\nnn\\ made with the IRAM 30\\,m with 7\\arcsec\\ spacing in $^{12}$CO(1--0) [top] and $^{12}$CO(2--1) [bottom]. The positions are arcsec offsets relative to the phase tracking center of our interferometric observations (see Table \\ref{table1}). Each spectrum has a velocity scale from $-300$ to $300\\,{\\rm km\\,s^{-1}}$, and a beam-averaged radiation temperature scale (T$_{\\rm mb}$) from $-0.10$ to $0.48\\,{\\rm K}$ for $^{12}$CO(1--0) and from $-0.25$ to 0.70\\,K for $^{12}$CO(2--1). } \\label{fig:n3627-30m} \\end{figure*} ", "conclusions": "} The molecular gas, traced by $^{12}$CO(1--0) and $^{12}$CO(2--1) transitions, in the interacting barred LINER/Seyfert 2 galaxy \\nnn\\ is distributed along a bar-like structure of $\\sim$18\\arcsec\\ ($\\sim$900\\,pc) diameter (PA = 14$^{\\circ}$) with two peaks at the extremes. The 1.6\\,$\\mu$m $H$-band 2MASS and 3.6\\,$\\mu$m \\spitzer-IRAC images of \\nnn\\ show a stellar bar in the nucleus with PA = $-21^{\\circ}$, different from the PA (= $14^{\\circ}$) of the molecular gas bar-like structure, suggesting that the gas is leading the stellar bar. Instead, the \\textit{GALEX} FUV emission of \\nnn\\ displays an inner elongated (north/south) ring delimiting a hole around the nucleus and containing the $^{12}$CO $\\sim$18\\arcsec\\ bar-like structure. This kind of anti-correlation between FUV, molecular gas, and the stellar bar is perhaps related to star formation efficiency and timescale variations in response to a spiral density wave. The gravity torques exerted by the stellar potential on the gas computed with the \\hst-NICMOS F160W image and our PdBI maps are negative within the inner 0.4\\,kpc, down to the resolution limit of our observations. The torques computed with the \\spitzer-IRAC 3.6\\,\\micron\\ image and BIMA $^{12}$CO map (with a resolution limit of $\\sim$670\\,pc) are also negative within the inner $\\sim$3\\,kpc. If the bar ends at $\\sim$3\\,kpc, with a pattern speed of the bar $\\Omega_{\\rm p}$$\\sim$65\\,km s$^{-1}$ kpc$^{-1}$ then the CR of the bar would be at $\\sim$3.3\\,kpc. There is no clear ring signature and we thus exclude the presence of an ILR. Without an ILR, the torques are negative between the CR of the bar and the AGN, down to the resolution limit of our observations. This scenario is compatible with a young/incipient bar which had no time to slow down from secular evolution, and has not yet formed any ILR. \\nnn\\ is a potential \\textit{smoking gun} of inner gas inflow: the gas is certainly fueling the central region, and in a second step could fuel directly the AGN. \\nnn\\ is the fifth \\textit{smoking gun} NUGA galaxy, together with NGC\\,6574 \\citep{lindt-krieg08}, NGC\\,2782 \\citep{leslie08}, NGC\\,3147 \\citep{vivi08a}, and NGC\\,4579 \\citep{santi09}. The common feature shared by these galaxies is a slowly rotating stellar bar (or oval) with overlapping dynamical resonances \\citep{vivi08b} associated with kinematically decoupled inner bars or ovals. Such resonances and kinematic decoupling are fostered by a large central mass concentration and high gas fraction. \\nnn\\ is the unique potential \\textit{smoking gun} NUGA galaxy with only one slowly rotating stellar bar. For \\nnn, this drives a molecular bar-like structure, apparently sufficient to transport the gas toward the AGN that, in a second step, could fuel directly the active nucleus." }, "1101/1101.5971_arXiv.txt": { "abstract": "{ Parallax measurements of pulsars allow for accurate measurements of the interstellar electron density and contribute to accurate tests of general relativity using binary systems. The Square Kilometre Array (SKA) will be an ideal instrument for measuring the parallax of pulsars, because it has a very high sensitivity, as well as baselines extending up to several thousands of kilometres. We performed simulations to estimate the number of pulsars for which the parallax can be measured with the SKA and the distance to which a parallax can be measured. We compare two different methods. The first method measures the parallax directly by utilising the long baselines of the SKA to form high angular resolution images. The second method uses the arrival times of the radio signals of pulsars to fit a transformation between time coordinates in the terrestrial frame and the comoving pulsar frame directly yielding the parallax. We find that with the first method a parallax with an accuracy of 20\\% or less can be measured up to a maximum distance of 13\\,kpc, which would include 9\\,000 pulsars. By timing pulsars with the most stable arrival times for the radio emission, parallaxes can be measured for about 3\\,600 millisecond pulsars up to a distance of 9\\,kpc with an accuracy of 20\\%. ", "introduction": "Distance measurements have always played a hugely important role in most aspects of astronomy, on both Galactic and extragalactic scales, but are often very difficult to access observationally. Arguably, the most reliable measurements are naturally obtained from parallax measurements when the annual movement of the Earth around the Sun can be used to detect a variation in the apparent source position. For nearly all sources and applications, this involves imaging the source against the background sky. This can be done at optical wavelengths or, usually more precisely, at radio frequencies using Very Long Baseline Interferometry (VLBI). In the case of radio pulsars however, a second type of parallax can be measured using timing measurements of the radio pulses. Here, distances are retrieved by detecting a variation in pulse arrival times at different positions of the Earth's orbit caused by the curvature of the incoming wavefront. In contrast to an imaging parallax, the precision of timing parallax measurements is highest for low ecliptic latitudes and lowest for the ecliptic pole. Pulsar timing of binary pulsars offers further possibilities to determine the distance from a secular or annual variation in some orbital parameters. In all cases, however, the methods are usually limited to relatively nearby sources, but improvement in telescope baseline lengths and sensitivity promise improvements in precision and hence the accurate measurement of distances. The most significant advance will be achieved by the Square Kilometre Array (SKA), which is being planned as a multi-purpose radio telescope constructed from many elements, both dishes and aperture arrays, leading to a total collecting area approaching 1 million square metres \\citep[e.g.][]{Dewdney09}. In the frequency range of 200\\,MHz to 3\\,GHz, the sensitivity of the SKA will be around 10\\,000\\,m$^2$/K, but depending on the design that will be chosen the sensitivity could be slightly higher or lower \\citep[see][]{Schilizzi2007}. The core of the SKA will be densely filled with 50\\% of the collecting area within a 5\\,km diameter area. The remaining elements will be placed at locations extending up to several thousands of kilometres from the core. The field of view (FoV) of the dishes will be around 0.64\\,deg$^2$ at 1.4\\,GHz for a single-pixel receiver and might be extended up to 20\\,deg$^2$ by means of phased array feeds, although for the long baselines only the single-pixel receivers will be available. Having both a dense core and long baselines allows the SKA to both find radio pulsars and to perform accurate astrometric measurements on them. In particular, the SKA will be able to perform parallax measurements on a great number of pulsars, allowing accurate determination of their distances from the Earth. Most distances to pulsars are currently derived from the measured column density of free electrons between the pulsars and Earth, known as the dispersion measure (DM). Using a model for the interstellar medium and its free electron content, the DM can be converted into a distance. However, the Galactic electron distribution is not accurately known, so that the distance derived from DM values is often uncertain, the errors in the distance estimates being as large as 100\\% \\citep[see e.g.][]{Deller09}. The SKA will perform a Galactic census of pulsars \\citep{ckl+04} that will potentially yield a population of 20,000 or more radio pulsars across the Galaxy \\citep{Smits09}. Parallax measurements of many of these pulsars with the SKA will provide direct measurements of their distances, which will then allow accurate studies of the ionised component of the interstellar medium, given the measurement of the DM. In this way, the interstellar electron model can be calibrated and will not only provide reliable estimates of the DM distances of pulsars without a parallax measurement in return, but will simultaneously provide a map of the free electron content in the Milky Way that can be combined with HI and HII measurements to unravel the Galactic structure and the distribution of ionised material. Combined with Faraday rotation measurements of the same pulsars, the Galactic magnetic field can also be studied in much greater detail as today, since the precise distance measurements potentially allow us to pin-point field-reversals occurring with some accuracy. The determination of accurate pulsar distances is also important for those pulsars that are part of binary systems, in particular for those with another compact object, such as the Double pulsar \\citep{Burgay2003,Lyne2004}. Relativistic effects can be used to determine the distance to some of these systems when the validity of general relativity is assumed. In reverse, to perform precision tests of general relativity, kinematic effects have to be removed for which it is often required to know the distance precisely \\citep{lk05,ksm+06,Deller09}. In this paper, we review the possible ways to determine pulsar distances with the SKA. To compare the technical capabilities with the expected pulsar population, we simulate an SKA pulsar survey, following \\citet{Smits09}. Using the results of this simulation we calculate the accuracy with which the SKA can determine the distance to each pulsar by measuring its parallax. We examine two basic methods for performing the astrometry, the first that uses imaging and the long baselines of the SKA, and the second that uses the timing of the pulsar radio signal. Here we distinguish between parallax measurements (applicable to all pulsars) and distance-related effects for binary pulsar parameters. In Sect. \\ref{sec:simulation}, we describe the simulation of the SKA pulsar survey. The two different methods for measuring the parallax of pulsars are explained in ref{sec:parallax}, where we also explain how we estimate the capability of the SKA to measure the parallaxes of pulsars. The results are presented in \\ref{sec:results}. Finally, we discuss our findings in \\ref{sec:discussion}. ", "conclusions": "\\label{sec:discussion} We have investigated two methods to measure the parallax of radio pulsars using the SKA. The imaging parallax method utilises the long baselines of the SKA to keep track of the position of the pulsar on the sky during different orbital positions of the Earth. The timing method uses the arrival times of the radio signals of pulsars to derive a ``classical'' parallax, an ``orbital-annual'' parallax, or a ``kinematic'' parallax. With imaging parallax, it is possible to determine parallaxes with an error of 20\\% or smaller for 9\\,000 pulsars out to a distance of 13\\,kpc. Measuring the parallaxes of all the pulsars in the Galactic plane alone will require 215 full days of observing. Furthermore, to obtain the required astrometric accuracy of 15\\,$\\mu$as requires several calibrators to be within a few arc-minutes of the source. With the timing parallax method, parallaxes can be measured out to 9\\,kpc for about 3\\,600 MSPs, with 20\\% or higher accuracy. This would require a minimum of roughly 40 days of observation time. However, there will have to be follow-up timing in any case for each MSP that will be detected by the SKA to select the most interesting MSPs. Therefore, obtaining parallaxes for many of these MSPs will not require extra observation time. When we consider the increased timing precision with the SKA and the recent realization that the apparent ``timing noise'' in non-millisecond pulsars is non-random and of magnetospheric origin, possibly allowing to correct for it \\citep{lhk+10}, it seems likely that parallax measurements will even be possible for normal, non-recycled pulsars. In this case, the sample of available pulsars could easily exceed the number of 10\\,000 but this is difficult to estimate at this stage. The imaging and timing parallax methods are complementary, with timing parallax providing high quality measurements for a relatively small number of objects but to large distances, whereas imaging parallax can provide good results for a large number of closer objects. Timing parallax is limited to millisecond pulsars with stable timing characteristics, whereas an imaging parallax depends only on the observed strength of the radio emission, not the rotation characteristics of the pulsar. Imaging parallax measurements can indeed help us significantly with precision tests of general relativity to correct for the same kinematic effects that we can otherwise use for precise distance measurements. Moreover, obtaining imaging information immediately after the discovery of a new pulsar will help us to obtain positional information that will help the typical ``solving procedure'' of a pulsar, i.e.~the first determination of a coherent timing solution. In general, solving pulsars takes several weeks to months of intensive observing, but the positional information from imaging eradicates the correlations between position and period derivative and allows us to reduce the cadence of the observations significantly. From the above results, it is clear that the SKA will become a superb astrometry instrument, the data of which feeds directly back into fundamental astrophysical questions. For instance, the distance to about 9\\,000 pulsars can be measured with an error of 20\\% or smaller. The dispersion measure from each pulsar can then be used to accurately calculate the electron density between the Earth and each of these pulsars. This will result in a map of the distribution of the ionised gas of high accuracy out to a distance of 13\\,kpc. By timing the most stable MSPs that the SKA will discover, the ionised gas can be mapped to a distance of 9\\,kpc. By imaging pulsars, we can locate them with a precision that is comparable with that of GAIA, which is able to measure distances out to 10 kpc with a precision of 10\\%. This is also achievable by timing MSPs and may be able to be, for selected sources, improved upon. This is exciting as we will be able to get independent distance measurements from GAIA and the SKA for pulsars that have an optical companion or are located, e.g. in a globular cluster. The simulations presented in this paper are not corrected for the Lutz-Kelker bias \\citep{Lutz93}. This bias is introduced when the parallax measurement does not take into account the larger volume of space that is sampled at smaller parallax values and leads to a systematic overestimate of the parallax, thus to an underestimate of the distance to the star. This bias is strongly dependent on the accuracy of the parallax measurement, becoming smaller as the measurement becomes more accurate. ~\\citet{vlm10} have studied this bias for the specific case of the parallax measurements of neutron stars, incorporating the bias introduced by the intrinsic radio luminosity function and a realistic Galactic population model for neutron stars. They found a significant bias for measurements with a 50\\% error. In the present paper, we consider parallax accuracies of 20\\% and higher. We find that in our simulations the Lutz-Kelker bias is always smaller than one standard deviation and can be neglected for most pulsars, hence does not affect our conclusions. Nevertheless, since the bias in a pulsar parallax measurement depends on the luminosity of the pulsar and the position on the sky, it is advisable to correct for it in every pulsar parallax measurement\\footnote{The Lutz-Kelker bias for pulsars can be determined online: \\url{http://psrpop.phys.wvu.edu/LKbias}}. Overall, our result shows that the SKA requires both a concentration of collecting area in the core, to search for and detect the Galactic pulsar population and then measure accurate timing parallaxes, and a distribution of collecting area on long baselines, to achieve the angular resolution required to perform imaging parallax measurements. Thus, the current design specifications for the SKA are adequate for achieving both of these objectives. \\acknowledgement{The authors would like to thank an anonymous referee for his/her comments. This effort/activity is supported by the European Community Framework Programme 6, Square Kilometre Array Design Studies (SKADS), contract no 011938. We gratefully acknowledge support from ERC Advanced Grant ``LEAP'', Grant Agreement Number 227947 (PI Michael Kramer). The International Centre for Radio Astronomy Research is a Joint Venture between Curtin University and The University of Western Australia, funded by the Western Australian State Government.}" }, "1101/1101.3550_arXiv.txt": { "abstract": "We investigate how a range of physical processes affect the cosmic metal distribution using a suite of cosmological, hydrodynamical simulations. Focusing on redshifts $z=0$ and 2, we study the metallicities and metal mass fractions for stars as well as for the interstellar medium (ISM), and several more diffuse gas phases. We vary the cooling rates, star formation law, structure of the ISM, properties of galactic winds, feedback from AGN, supernova type Ia time delays, reionization, stellar initial mass function, and cosmology. In all models stars and the warm-hot intergalactic medium (WHIM) constitute the dominant repository of metals, while for $z \\ga 2$ the ISM is also important. In models with galactic winds, predictions for the metallicities of the various phases vary at the factor of two level and are broadly consistent with observations. The exception is the cold-warm intergalactic medium (IGM), whose metallicity varies at the order of magnitude level if the prescription for galactic winds is varied, even for a fixed wind energy per unit stellar mass formed, and falls far below the observed values if winds are not included. At the other extreme, the metallicity of the intracluster medium (ICM) is largely insensitive to the presence of galactic winds, indicating that its enrichment is regulated by other processes. The mean metallicities of stars ($\\sim Z_\\odot$), the ICM ($\\sim 10^{-1}\\,Z_\\odot$), and the WHIM ($\\sim 10^{-1}\\,Z_\\odot$) evolve only slowly, while those of the cold halo gas and the IGM increase by more than an order of magnitude from $z=5$ to 0. Higher velocity outflows are more efficient at transporting metals to low densities, but actually predict lower metallicities for the cold-warm IGM since the winds shock-heat the gas to high temperatures, thereby increasing the fraction of the metals residing in, but not the metallicity of, the WHIM. Besides galactic winds driven by feedback from star formation, the metal distribution is most sensitive to the inclusion of metal-line cooling and feedback from AGN. We conclude that observations of the metallicity of the low-density IGM have the potential to constrain the poorly understood feedback processes that are central to current models of the formation and evolution of galaxies. ", "introduction": "The spatial distribution of elements synthesised in stars (henceforth `metals') provides an archaeological record of past star formation activity and of the various energetic phenomena that stirred and mixed these metals. Recent cosmological simulations of galaxy formation follow the different stellar evolutionary channels through which metals are produced, and include some processes that cause metals to escape from their parent galaxy. Here we will investigate to what extent these simulations produce realistic enrichment patterns, which physical processes affect the metal distribution, and how robust the predictions are. Observations can constrain integrated stellar and nebular (ISM) abundances of galaxies (for $z\\la 1$ see e.g.\\ \\citealt{Kobulnicky1999, Kunth2000, Teta04, Dunne2003}, for $z\\approx 1$ see \\citealt{Churchill2007}, for $z\\ga 2$ see \\citealt{Eeta06,Mannucci2009}). Metals in cold gas can be observed in absorption against a background source, see the damped Lyman-alpha observations at $z \\approx 3$ \\citep[e.g.][]{Pettini1994, Prochaska2003} and at lower \\citep[e.g.][]{Meiring2009} and higher redshift \\citep[e.g.][]{Ando2007,Price2007}. If the background source is sufficiently bright, then this technique can also be applied to more diffuse gas \\citep[e.g.][] {SC96, Cowie1998, Ellison2000, Seta00a, Seta03, Schaye2007, Simcoe2004, Scannapieco2006b, Aguirre2008}. The metallicity of the ICM is inferred from X-ray observations \\citep[e.g.][]{Mushotzky1996, dp2006, Sato2007a, Rasmussen2007, Maughan2008, Leccardi2008, Snowden2008}. These observations use different tools to observe a variety of elements in a variety of ionization states, and it is not always obvious how to convert all these measurements to a common `metallicity'. Often this is done assuming that the relative abundances of elements equal those measured in the Sun. Unfortunately even the assumed metallicity of the Sun itself varies between authors. Here we assume $Z_\\odot=0.0127$, and solar abundances from the default settings of {\\sc cloudy} (version 07.02, last described by Ferland et al.\\ 1998). Convolving metallicity with the fraction of baryonic mass in each of the different phases then yields a census of cosmic metals \\citep[e.g.][]{Fuku98}. It must, however, be kept in mind that a large fraction of the metals are potentially not accounted for by the current data. For example, a large portion of $z=0$ baryons are thought to reside in the warm-hot intergalactic medium (WHIM), which has not yet been convincingly detected. If, as is likely, the WHIM is also metal enriched, then it could harbour a significant amount of metals. Similarly, the prevalence of hot ($T\\ga 10^5$~K) metals at $z\\ga 2$ is currently poorly constrained. Gas cooling, star formation, galaxy interactions, ram pressure stripping, stellar and galactic winds all cause gas and hence metals to be cycled between the different phases, which makes it complicated to relate the observed metal distribution to the original source and/or enrichment process. A theoretical calculation of the census of cosmic metals must take into account their production by nucleosynthesis, their initial distribution, and the mixing that occurs in later stages. Numerical simulations that include such \\lq chemo-dynamics\\rq\\ have become increasingly sophisticated since the early work by \\citet{TBH92}, usually concentrating on the evolution of a single galaxy \\citep{SM95, RVN96, Berczik1999,Reta01,K01a,KG03,Kobayashi2004,MU06a, Seta06a, Geta07,Bekki2009a, Rahimi2010, Pasetto2010}, or a cluster of galaxies \\citep[e.g.][]{LPC02, Valdarnini2003,Toeta04, Sommer-larsen2005, Reta06, Tornatore2007b, Fabjan2010} in a cosmological context. These simulations include interpolation tables from stellar evolution calculations for the production of metals, but implement the enrichment processes explicitly, for example by injecting metal-enriched gas near a site of star formation in a galactic wind. The fluxes of metals between phases due to gas cooling and galaxy/gas interactions are computed explicitly by these hydro-codes. Early simulations that were also used to look at the metals outside galaxies include \\cite{Cen1999a}, \\cite{Mosconi2001} and \\cite{Theuns2002}. Subsequent authors implemented more physics while increasing resolution and/or box size in order to minimise numerical effects \\citep[e.g.][]{Seta05, OD06, Cen2006, Brook2007, Kobayashi2007,Oppenheimer2008, Wiersma2009b,Wiersma2010,Shen2009, Tornatore2009}. In this paper we use cosmological hydrodynamical simulations of the formation of galaxies to attempt to answer the question `{\\em Where are the metals}?', by computing the evolution of the fraction of metals in stars and various gas phases. Our suite of numerical simulations (the OWLS suite; \\citealt{Schaye2010}), includes runs that differ in terms of their resolution, input physics, and numerical implementation of physical processes. We use it to investigate what physical processes are most important, how reliable the predictions are, and to what extent these depend on the sometimes poorly understood physics \\cite[see][for a similar invesitgation]{Sommer-Larsen2008}. In \\cite{Wiersma2009b} we introduced the method and described some of the numerical issues involved, but analysed only a single physical model (the OWLS reference model). Here we vary the cooling rates, star formation law, structure of the ISM, properties of galactic winds, feedback from AGN, supernova type Ia time delays, reionization, stellar initial mass function, and cosmology. This paper is organised as follows. Section~\\ref{sec-method} introduces the simulations used. Those familiar with the OWLS project may skip to the results in Section~\\ref{sec-WATM}, which begins with an overview before discussing the physical variations which are the most relevant to the metal distribution. This section closes with a summary of all the simulations in the OWLS suite in Section~\\ref{sec-sum}. In Section~\\ref{sec-concs} we present our conclusions. ", "conclusions": "\\label{sec-concs} We have investigated the distribution of metals predicted by simulations taken from the OverWhelmingly Large Simulations project (\\textsc{OWLS}; \\citealt{Schaye2010}). We considered different recipes and physical parameters in order to determine what the factors are that shape the distribution of metals as a function of the gas density, temperature, and metallicity. This builds on our previous work where we introduced our method and demonstrated the effect of simulation box size and resolution \\citep{Wiersma2009b}. Our results can be divided between statements about the effects of physical processes on the metal mass distribution and statements about the metal mass distribution itself. The former can be summarised as follows: \\begin{itemize} \\item Metal-line cooling increases the fraction of the metals that reside in stars and the ISM, at the expense of the fraction in highly overdense warm-hot gas. Metal-line cooling increases the metallicities of all cold phases. In particular, it increases the metallicities of the diffuse IGM and cold halo gas by about a factor of two. \\item While other mechanisms can enrich the warm-cold IGM (e.g.\\ ram pressure and tidal stripping), without SN feedback, its metallicity is far smaller and is strongly ruled out by observations of QSO absorption lines. \\item Processes other than outflows driven by feedback from star formation and AGN contribute significantly to, or even dominate, the enrichment of the ICM. \\item Even for a constant energy per unit stellar mass formed (i.e.\\ a fixed feedback efficiency), the freedom given by the parameters of sub-grid implementations of galactic winds is very large. At a fixed efficiency, higher wind velocities (and thus correspondingly low mass loading factors) increase the fraction of the metals that reside in low-density ($\\rho < 10^2\\left <\\rho\\right >$) gas. However, since higher wind velocities also shift metals to higher temperatures, the metallicity of the diffuse, warm-cold IGM is in fact higher for lower wind velocities. \\item The metallicity of the warm-cold diffuse IGM is most sensitive to the implementation of galactic winds. \\item By varying the parameters of the wind model with local properties, one can significantly affect the metal distribution in the IGM. For example, a model in which low-mass galaxies drive highly mass-loaded, slow winds while high-mass galaxies drive fast winds with low mass loading factors \\citep[e.g.][]{Oppenheimer2008} can efficiently enrich the diffuse warm-cold IGM while still limiting the build up of metals in the stellar components of more massive galaxies. \\item Feedback from AGN strongly reduces the fractions of metals that reside in the ISM and stars as well as the metallicities of these two components. Fast galactic winds in massive galaxies, e.g.\\ resulting from a top-heavy IMF at high gas pressures, work in the same direction but have a smaller effect than AGN. Efficient feedback in high-mass galaxies only has a relatively small effect on the metallicities of the diffuse gas components. \\item Changes in the star formation law, the structure of the ISM, the cosmology, and the IMF all play a role, but are less important than the inclusion and implementation of metal-line cooling and, most importantly, outflows driven by feedback from star formation and accreting supermassive black holes. We also find a small effect for changes in the reionisation history, but for the case of the IGM this may be due to our limited resolution. \\end{itemize} We summarise the metal mass distribution as follows: \\begin{itemize} \\item Stars and the WHIM are the dominant depositories of metals. In all models they contain together at least 78 \\% of the metals at $z = 0$ and at least 53 \\% at $z = 2$. At high redshift the ISM also contains a large fraction of the metals: at least 11 \\% in all simulations at $z=2$. In all our models, and at both $z=0$ and $z=2$, the remaining components (diffuse cold-warm IGM, cold halo gas, and ICM) together contain less than a quarter of the metals. \\item The mean metallicities of the WHIM and ICM are nearly constant in time. In most models they increase slightly with time to $\\sim 10^{-1}\\,Z_\\odot$ at $z = 0$. In contrast, the metallicities of the cold halo gas and the diffuse IGM, increase by more than an order of magnitude from $z=5$ to 0. \\item All our models predict, both for $z=0$ and 2, mean metallicities of $Z \\sim Z_\\odot$ for stars and ISM, and $10^{-1}\\,Z_\\odot \\la Z < Z_\\odot$ for the ICM and for cold halo gas. Except for the models without SN feedback, all simulations predict, both for $z=0$ and 2, mean metallicities of $Z \\sim 10^{-1}\\,Z_\\odot$ for the WHIM and $10^{-3}\\,Z_\\odot \\la Z \\la 10^{-2}\\,Z_\\odot$ for the diffuse cold-warm IGM. \\end{itemize} The inclusion and implementation of outflows driven by feedback from star formation and AGN is most important for predictions of the distribution of metals. As cosmological simulations will need to continue to use sub-grid implementations for these processes for the foreseeable future, this implies that predicting the distribution of metals will remain difficult for ab initio models. However, as we have shown, provided that some feedback is included, the mean metallicities of most different components can be robustly predicted to order of magnitude. The sensitivity of some of the results to feedback can also be regarded as fortunate. It gives us a useful tool to study the effects of outflows, which are an essential but very poorly understood ingredient of models of galaxy formation and evolution. The metallicity of the diffuse warm-cold IGM, which can be constrained using QSO absorption lines, is particularly sensitive to variations in the models and is therefore a very promising probe of the physics of galactic outflows." }, "1101/1101.3599_arXiv.txt": { "abstract": "We revisit the problem of low-mass pre-main-sequence (PMS) stellar evolution and its observational consequences for where stars fall on the Hertzsprung-Russell diagram (HRD). In contrast to most previous work, our models follow stars as they grow from small masses via accretion, and we perform a systematic study of how the stars' HRD evolution is influenced by their initial radius, by the radiative properties of the accretion flow, and by the accretion history, using both simple idealized accretion histories and histories taken from numerical simulations of star cluster formation. We compare our numerical results to both non-accreting isochrones and to the positions of observed stars in the HRD, with a goal of determining whether both the absolute ages and the age dispersions inferred from non-accreting isochrones are reliable. We show that non-accreting isochrones can sometimes overestimate stellar ages for more massive stars (those with effective temperatures above $\\sim 3500$ K), thereby explaining why non-accreting isochrones often suggest a systematic age difference between more and less massive stars in the same cluster. However, we also find the only way to produce a similar overestimate for the ages of cooler stars is if these stars grow from $\\sim 0.01$ $\\msun$ seed protostars that are an order of magnitude smaller than predicted by current theoretical models, and if the size of the seed protostar correlates systematically with the final stellar mass at the end of accretion. We therefore conclude that, unless both of these conditions are met, inferred ages and age spreads for cool stars are reliable, at least to the extent that the observed bolometric luminosities and temperatures are accurate. Finally, we note that the time-dependence of the mass accretion rate has remarkably little effect on low-mass stars' evolution on the HRD, and that such time-dependence may be neglected for all stars except those with effective temperatures above $\\sim 4000$ K. ", "introduction": "\\label{sec:intro} Pre-main-sequence (PMS) stars in low-mass star forming regions show a sizable luminosity spread when placed on the Hertzsprung-Russell diagram (HRD) \\citep[e.g.,][]{hillenbrand09}. This spread translates into a significant dispersion in inferred stellar ages that records the past star formation activity in each region (e.g., \\citealt{dantona94, baraffe98, siess00, palla00, palla99, hartmann01,hartmann03}). However, the idea that star clusters form over an extended period is subject to extensive debate on both observational and theoretical grounds \\citep[e.g.,][]{elmegreen00a, hartmann01a, tan06a, krumholz07a, evans09a}, and several authors have claimed that the dispersion of stellar luminosities does not reflect a real age spread. Members of young binaries and multiples exhibit a tighter age correlation, supporting the existence of an intrinsic age distribution. However, a luminosity spread persists even among such systems, and some companions display a substantial age mismatch \\citep{prato03, stassun08,krauss09}. Deriving stellar ages is complicated from an observational standpoint. Calculation of stellar bolometric luminosities is beset by uncertainties in extinction, photometric variability, and unresolved multiplicity. Calibration between the stellar spectral type and effective temperature is also not trivial. In some cases, it can be demonstrated that observational uncertainties alone are sufficient to induce an age spread of $> 10$ Myr and mask a coeval stellar population \\citep{slesnick08}. However, \\citet{dario10b, dario10a} carefully model these uncertainties and conclude that these effects alone cannot reproduce the entire spread. Other age indicators such as stellar rotation rate \\citep{jeffries10}, surface gravity \\citep{slesnick08}, and lithium abundances \\citep{sestito08} also support the idea that the inferred age spreads are real, but are each subject to significant challenges. Apart from the observational uncertainties, physical mechanisms may be responsible for a portion of the observed HRD scatter. For the purpose of inferring stellar ages, it is usually assumed that PMS stars first appear along a ``birthline'' in the HRD when mass accretion ceases (e.g., \\citealt{PS90,hartmann97}). However, luminosities of younger embedded stars (Class 0 and I sources) that are presumably still accreting also show a wide spread, and a fraction of them have luminosities much lower than the values expected from the standard birthline (e.g., \\citealt{kenyon90, evans09, enoch09}). A solution for this ``luminosity problem'' is the scenario that mass accretion takes place very time-dependently, repeating burst-like accretion phases and quiescent phases. Recent numerical simulations suggest that such episodic mass accretion is caused by gravitational fragmentation of a circmustellar disk (e.g., \\citealt{vorobyov05, machida11}), though radiative warming from protostars alleviates it \\citep[e.g.,][]{offner09}. Regardless of the ultimate explanation for the luminosities of Class 0 and I sources, the existence of young stars that fall well below the putative birth line is strong evidence that we must extend our PMS evolution models to include the accretion phase. \\citet[][hereafter BCG09]{baraffe09} study protostellar evolution with various episodic mass accretion histories and examine the resultant spread of PMS stars in the HRD. They argue that PMS stars of the same mass and age show some scatter in the HRD owing to variation of the early evolution resulting from complex accretion histories. However, BCG09 simultaneously vary not only their accretion histories, but also their initial stellar models and the radiative properties of the accretion flow. Because they change these parameters in correlated ways and do not perform a systematic survey of parameter space, it is not clear which of these effects drives their results. Nor is it clear whether the results they generate via their parameter choices are consistent with observed HRDs of clusters. Consequently, it is still unclear how much vigorous time-dependent accretion histories influence protostellar evolution. In this paper, we aim to resolve this question by performing a systematic study of how PMS evolutionary tracks change as we alter the accretion history, the initial models, and the thermal efficiencies of mass accretion. We perform a systematic survey of parameter space in order to understand how each of these factors affects protostellar evolution. This enables us to answer the question of whether variation in any of these quantities could produce the appearance of an age spread in a population that is actually coeval. The structure of the paper is as follows. In Section \\ref{sec:method}, we briefly explain our numerical method for modeling protostellar evolution. Section \\ref{sec:results} is the main part of the paper, where the numerical results are presented. First, we investigate how different accretion histories influence protostellar evolution in Section \\ref{ssec:acchist}. We next investigate protostellar evolution with differing initial models in Section \\ref{ssec:ini}, and with differing thermal efficiencies in \\ref{ssec:eff}. In Section \\ref{sec:reliability} we combine all these results to draw general conclusions about the reliability of age and age spread estimates from PMS evolutionary tracks. Section \\ref{sec:sum} contains the summary and discussion. ", "conclusions": "\\label{sec:sum} In this paper, we have examined a variety of low-mass protostellar evolutionary tracks with varying accretion histories, initial models, and thermal efficiencies of mass accretion. We have also compared the resultant spread of PMS stars in the HRD to that observed in nearby low-mass star forming regions \\citep{peterson08, gatti06,gatti08, muzerolle05}. We first calculate protostellar evolution models with varying accretion histories but fixed initial stellar models and boundary conditions (Sec.~\\ref{ssec:acchist}). Our results show that if mass accretion is thermally inefficient, variation in the accretion history hardly influences protostellar evolution. Although isochrones for non-accreting protostars, such as those calculated by \\citet{dantona94, baraffe98, siess00}, do not necessarily provide us with correct stellar ages, models with differing accretion histories nevertheless form a tight sequence in the HRD. Thus variable accretion histories alone cannot explain the observed spread of PMS stars in the HR. Moreover, the errors in absolute ages arise because non-accreting isochrones are not good descriptions of stars growing with thermally inefficient mass accretion. They are not a result of variable accretion histories. However, we note that this does suggest that accreting isochrones may resolve the problem of systematically larger inferred ages for high-mass cluster members \\citep[e.g.,][]{hillenbrand09,covey10}. Second, we examine protostellar evolution with different initial models and thermal efficiencies of mass accretion, using a constant accretion rate $\\dot{M} = 10^{-5}~\\msunyr$ (Sec.~\\ref{ssec:ini} and \\ref{ssec:eff}). We find that the spread of PMS stars in the HRD that results from varying the initial radius (or entropy) or the thermal efficiency is much larger than the spread that arises from different accretion histories. We find that a coeval population of stars with significant star-to-star variation in thermal efficiency or initial radius could conceivably occupy the entire observed luminosity range for protostars with effective temperatures $\\ga 3500$ K. Thus ages and age spreads observed in this temperature range may be unreliable. At lower effective temperature, however, the situation is very different. The only models we found in our parameter space survey that are capable of producing false old ages at low $T_{\\rm eff}$ are those with purely cold accretion starting from very small initial radii. However, we point out that these models require that $\\sim 0.01~M_\\odot$ second protostellar cores have radii $<0.3~R_\\odot$, more than an order of magnitude smaller than any formed in simulations to date (e.g., MI00), and that they can be rendered consistent with observations only if such small initial radii are realized only for stars that end up growing to small final masses. Neither possibility can be definitively ruled out, but neither is supported by any current observations or theory either. If we exclude very small initial radii on these grounds, we find that all the remaining models indicate that ages and age spreads inferred from non-accreting isochrones are reliable for cool stars, at least to the extent that the observationally-determined luminosities and temperatures are reliable \\citep[e.g.,][]{dario10a, dario10b}. In varying the thermal efficiency, we find that models with only a small amount of hot accretion, e.g., while $M_* \\le 0.03~\\msun$, nonetheless show similar evolution to models with entirely hot accretion. Thus, in order to explain the HRD spread at low-masses with cold accretion models, the models must begin with small radii and be thermally inefficient for nearly their entire evolution. Observational constraints on the thermal efficiency of accretion are somewhat tenuous. During the earliest stages during which accretion rates of $10^{-6}-10^{-4}~\\msun$ yr$^{-1}$ occur, protostars are too deeply embedded with too much radiation reprocessing to measure accretion signatures directly. Observations of T Tauri stars, for which accretion rates have declined to $\\lesssim 10^{-7}~ \\msun$ yr$^{-1}$, suggest that for these rates the accretion column covers only $\\sim$1-10\\% of the stellar surface. This supports a cold accretion senario later in the accretion history. However, observations also find that the covering fraction increases with accretion rate \\citep{gullbring00, ardila00}, suggesting that accretion may be more thermally inefficient at early times. Higher mass stars, which experience higher initial accretion rates, seem unlikely to avoid hot accretion, while it is more observationally probable that very low-mass stars experience purely cold accretion. Our conclusions are different from those of BCG09, who stressed the significance of episodic mass accretion for explaining the observed spread of PMS stars in the HRD. However, our numerical results are actually consistent with theirs. First, BCG09 calculated protostellar evolution using simple, non-episodic accretion histories (their Fig.~1) and found that, with thermally efficient accretion ($\\alpha \\geq 0.2$ in their notation), isochrones for non-accreting protostars give the correct ages at $t \\gtrsim 1$~Myr. Disagreement with the isochrones arises only for thermally inefficient accretion flows ($\\alpha = 0$). This is consistent with the results shown in our Figure \\ref{fig:ri_pms}. Next, they calculated the evolution with more vigorous, episodic mass accretion histories (their Fig.~2). However, the spread they obtain is no broader than that shown in their Figure 1, indicating that episodic accretion does not increase the HRD spread beyond what they had already introduced by using varying initial conditions and thermal efficiencies. Indeed, they state that the time-dependence of the accretion rates is not essential for their results, and we confirm this finding.\\footnote{\\citet{baraffe10} argue that episodic mass accretion also significantly influences the lithium depletion of low-mass stars. However, our present work suggests that variation of the initial models will be more significant than variation of the accretion histories for the problem of lithium depletion as well.} Our results suggest that BCG09 were able to obtain small luminosities for stars with $T_{\\rm eff} \\la 3500$ K, and thus claim to reproduce the observed HRD using a coeval population, because they used cold accretion starting from initial conditions with entropies far lower than what current theoretical models predict. They also did not continue runs with these initial conditions up to higher masses and effective temperatures.\\footnote{For example, the only cases BCG09 show where stars with $T_{\\rm eff} < 3500$ K lie near the 10 Myr non-accreting isochrone after 1 Myr are their models A-C. However, they only use the initial conditions and accretion rates for models A-C to produce stars up to $0.2$ $\\msun$. In comparison, all the models that they run to masses above $0.5$ $\\msun$ have much higher starting entropies.} Our results suggest that, had they done so, the resulting stars would have fallen well below the locus of observed stars in the HRD, as do our comparable low initial entropy models. Finally, we stress that we do not reject the episodic mass accretion as a possible solution for the ``luminosity problem'' of young embedded sources \\citep{dunham10, mckee10, offner11}. Episodic accretion may well occur. It is simply not capable of explaining the broad spread of optically-visible PMS stars in the HRD. {" }, "1101/1101.4146_arXiv.txt": { "abstract": "Irradiation with high energy photons (10.2 $-$ 11.8 eV) was applied to small diamondoids isolated in solid rare gas matrices at low temperature. The photoproducts were traced via UV absorption spectroscopy. We found that upon ionization the smallest of these species lose a peripheral H atom to form a stable closed-shell cation. This process is also likely to occur under astrophysical conditions for gas phase diamondoids and it opens the possibility to detect diamond-like molecules using their rotational spectrum since the dehydrogenated cations possess strong permanent dipole moments. The lowest-energy electronic features of these species in the UV were found to be rather broad, shifting to longer wavelengths with increasing molecular size. Calculations using time-dependent density functional theory support our experimental findings and extend the absorption curves further into the vacuum ultraviolet. The complete $\\sigma - \\sigma^*$ spectrum displays surprisingly strong similarities to meteoritic nanodiamonds containing 50 times more C atoms. ", "introduction": "Diamond-like material is expected to be abundant in the interstellar medium \\citep{Henning98}. Small nanodiamonds (2$-$3 nm) were extracted from meteoritic material, and they are the most abundant presolar grains in the primitive meteorites \\citep{lewis87, anders93, jones04}. Recently, diamondoid molecules were the subject of different experimental and theoretical studies revealing their spectroscopic properties \\citep{oomens06, bauschlicher07, lenzke07, pirali07, landt09a, landt09b}. These special hydrocarbons can be considered as faced-fused diamond cages where all carbon atoms are sp$^3$ hybridized with hydrogens saturating the dangling bonds on the surface. The diamond-like structure results in a remarkable rigidity, strength, and thermodynamic stability, especially compared to other hydrocarbons. The smallest of these species, adamantane C$_{10}$H$_{16}$, consists of only one diamond cage, followed by diamantane C$_{14}$H$_{20}$ with two, and triamantane C$_{18}$H$_{24}$ with three faced-fused cages. Diamondoids can be found in some natural gas reservoirs, and especially diamantane is one of the deposits in gas pipelines \\citep{reiser96}. Lately, diamondoids consisting of up to 11 diamond cages were isolated from petroleum by \\citet{dahl03} making these species accessible to laboratory investigations. \\citet{oomens06} measured the infrared spectroscopic properties of powders of higher diamondoids (up to hexamantane) at room temperature. In accordance with calculations applying density functional theory (DFT), the by far strongest features in the IR spectra of neutral diamondoids were found to be the C-H stretching bands between 3.4 and 3.6 $\\mu$m arising from the hydrogen-terminated surfaces. \\citet{pirali07} additionally measured the IR emission spectra of hot (500 K) gas phase adamantane, diamantane, and triamantane in the wavelength region of the C-H stretching modes revealing a small redshift of the bands in the solid-state spectra. Based on the IR spectra obtained by \\citet{oomens06} along with additional DFT calculations, \\citet{pirali07} also made assignments for two IR features of two different classes of astronomical objects. The first one is the unusual IR emission at 3.43 and 3.53 $\\mu$m originating from two objects whose spectra are elsewhere dominated by the well-known infrared emission bands of polycyclic aromatic hydrocarbons (PAHs), Elias 1 and the inner region of the circumstellar disk around HD 97048 \\citep{Habart04}. Although these bands were already assigned to nanodiamonds of at least 50 nm diameter \\citep{guillois99}, it was shown that also tetrahedral diamondoid molecules containing around 130 C atoms (close to the size of the smallest meteoritic nanodiamonds) exhibit the 3.43 and 3.53 $\\mu$m bands with the proper intensity ratio \\citep{pirali07}. The second feature mentioned is the broad (FWHM 0.09 $\\mu$m) absorption band centered at 3.47 $\\mu$m which is observed in the absorption spectra of various dense clouds in lines of sight toward young stellar objects \\citep{allamandola92, allamandola93}. By co-adding all diamondoid solid-state spectra obtained by \\citet{oomens06}, \\citet{pirali07} showed that the different C-H stretching modes merge into a broad structure centered around 3.47 $\\mu$m, and they argued that small diamond-like molecules may therefore be a major contributor to the interstellar absorption band which is only observed in or behind dense molecular clouds, but not in the diffuse interstellar medium. As the intensity of the interstellar band could be correlated with the intensity of the 3.08 $\\mu$m band of water ice, it was speculated that its carriers may be formed on icy grains in the shielded environment of molecular clouds \\citep{brooke96} which is furthermore supported by recent laboratory experiments where nanodiamond crystallites were created upon UV irradiation of interstellar ice analogs \\citep{kouchi05}. Based on the computed intensities of the diamondoid C-H stretching bands, \\citet{bauschlicher07} deduced that only 1$-$3\\% of the cosmic C has to be locked in diamondoids in order to explain the observed intensity of the interstellar band. Furthermore, they calculated ionization potentials (IPs) of neutral diamondoids, as well as IR spectra and electronic transition energies of neutral and cationic diamondoids, concluding that cations may also contribute to the 3.47 $\\mu$m absorption band. However, the bands of the cations are somewhat weaker than those of their neutral counterparts. Furthermore, the cations feature further bands with comparable strengths at longer wavelengths (6$-$18 $\\mu$m), e.g. due to C-H bending vibrations, which could be used in the future to trace ionized diamondoids. Compared to the strong IR emission of PAHs triggered by the absorption of UV-vis photons, the IR emission of neutral diamondoids is rather inefficient because their electronic absorption onset lies far in the UV between 6 and 7 eV \\citep{landt09a, landt09b}. Calculations imply that diamondoid cations would feature absorption bands in the visible and near-IR due to their open-shell structure \\citep{bauschlicher07}. However, these transitions are very weak and, as for the neutrals, efficient IR emission can only be expected in regions of space experiencing high-energy and high-flux UV radiation fields, usually in close proximity of the exciting stars \\citep{bauschlicher07}. Therefore the 3.43 and 3.53 $\\mu$m emission features are so rarely observed, to date only in HD 97048 and Elias 1. The IP of diamondoids \\citep[8$-$9 eV;][]{lenzke07} is only slightly higher than their band gap. \\citet{lenzke07} have shown that the photoion yield of diamondoids reaches its maximum between 10 and 11 eV, almost exactly around the hydrogen Ly$\\alpha$ emission. Hence, ionization of neutral diamondoids in strongly irradiated regions of space (HD 97048 and Elias 1) may be very efficient. Consequently, one should address the following questions: are these cationic species stable and, if so, what are their spectroscopic fingerprints? As already stated, one can expect weak absorption bands in the visible and near-UV wavelength range for the cations due to low-energy transitions to semi-occupied molecular orbitals, which might be detectable by electronic spectroscopy methods. Unlike in PAHs, there is no delocalized electron cloud. Removing one electron due to photoionization weakens the bonds between the atoms, and fragmentation (C$-$H bond breaking) may occur. Indeed, such behavior was observed for the three smallest diamondoids. By comparing the IR absorption bands of positively charged, gas phase adamantane \\citep{polfer04}, diamantane, and triamantane \\citep{pirali10} with theoretical spectra of dehydrogenated diamondoid cations, it was shown that these molecules easily lose a hydrogen atom upon ionization to form stable closed-shell species. The loss preferentially happens on a tertiary carbon (CH group) rather than on a secondary carbon (CH$_2$ group). Since in these experiments, the ionization was performed via charge transfer using cationic agents with high IPs, it is not obvious whether this dehydrogenation will also occur under astrophysical irradiation conditions. In this work, we investigated the electronic transitions of the four smallest diamondoids, namely adamantane C$_{10}$H$_{16}$, diamantane C$_{14}$H$_{20}$, triamantane C$_{18}$H$_{24}$, and tetramantane C$_{22}$H$_{28}$ (three isomers), and their photoproducts. For this purpose, we used matrix isolation spectroscopy (MIS). Cationic species were formed via UV irradiation using a hydrogen-flow discharge lamp to simulate the interstellar UV photon field. Our experimental findings are supported by theoretical calculations applying DFT and time-dependent DFT (TD-DFT). These results will help astronomers to search for the spectroscopic fingerprints of diamond-like molecular species. Furthermore, the UV absorption cross sections of neutral and cationic diamondoids can be used to accurately predict IR emission processes caused by stochastic heating due to the absorption of UV photons in strongly irradiated regions of space. ", "conclusions": "\\label{results} \\subsection{Adamantane} \\label{adam_section} In Figure \\ref{fig2}, the structure and calculated ground-state energy of adamantane, as well as the energies of the adamantane and adamantyl cations are displayed. Adamantane possesses two structurally different H sites and therefore two different adamantyl isomers. The numbers on the adamantane structure shown in Figure \\ref{fig2} indicate the C atoms from which the H atoms are removed to form the corresponding adamantyl structures. According to the calculations, only minor distortions of the carbon framework are expected upon H removal. This also applies to the larger species we have investigated. For the ground states of the adamantyl cations, different spin states are principally possible. However, the triplet state of the 1-adamantyl cation is about 3 eV higher in energy than the singlet state. Because of this rather high energy difference, we expect all other dehydrogenated diamondoid cations to be closed-shell singlet species as well and the calculations were performed accordingly. For completeness, we also calculated the structures and ground-state energies of the neutral 1- and 2-adamantyl radicals. In both cases, an energy of 4.2 eV would be necessary to remove one H atom. However, direct photodissociation is unlikely since the absorption onset is further in the UV. Basically, the photons delivered by the hydrogen lamp provide enough energy to cause photoionization, as well as the removal of one H atom from the ionized molecule. Further dissociation cannot be accomplished with a single photon. Considering the low FUV doses applied during the experiments, comprising typically 15$-$30 minutes of irradiation, further processing of already ionized molecules can be excluded. This was ensured by increasing the irradiation time up to 2 hr without noticeable change of the absorption bands. \\begin{figure}[t]\\begin{center} \\epsscale{1.1} \\plotone{f2.eps} \\caption{Zero-point-corrected ground-state energies of adamantane and its derivatives calculated at the B3LYP/6$-$311$++$G(2d,p) level of theory. The formation routes of the adamantyl cations in the experiment are indicated by solid arrows.} \\label{fig2} \\end{center}\\end{figure} The calculated spectra of neutral and cationic adamantane and adamantyl, using the B3LYP/6$-$311$++$G(2d,p) level of theory, are displayed in the upper two panels of Figure \\ref{fig3}. The theoretical spectra were obtained by convolving the shown stick spectra, representing the oscillator strength of each transition at its corresponding transition wavelength, with Lorentzian functions. While the area of each Lorentzian is proportional to the calculated oscillator strength, the width was chosen to be constant (3000 cm$^{-1}$). The chosen bandwidth is rather arbitrary, but the so-computed spectrum indicates in which energy range strong electronic transitions of the real molecule can be expected. The measured spectrum of FUV-processed adamantane in solid Ne (6.8 K) can be found in the bottom panel of Figure \\ref{fig3}. The ratio of Ne to adamantane in terms of absolute numbers of atoms or molecules was determined as described in Section \\ref{misfuv}. The densities of solid Ne and adamantane were taken to be 45 atoms nm$^{-3}$ \\citep{timms96} and 1.2 g cm$^{-3}$ \\citep{yashonath86}, respectively. Using these values, the isolation ratio (Ne to adamantane) was estimated to be better than 190 $n_{\\text{Ada}}$ $n^{-1}_{\\text{Ne}}$. The factor $n_{\\text{Ada}}$ $n^{-1}_{\\text{Ne}}$ is the ratio of the refractive indices of the pure solid materials in the visible and should be between 1 and 2. Nevertheless, we also performed measurements at lower isolation ratios and did not notice appreciable spectral differences. Since it is unknown what fraction of the neutral precursor molecules is actually transformed into cations\\footnote{Usually, the conversion rate is lower than 10\\% under these experimental conditions.} we cannot provide experimental values for the absorption cross section. \\begin{figure}[t]\\begin{center} \\epsscale{1.15} \\plotone{f3.eps} \\caption{Calculated (B3LYP/6$-$311$++$G(2d,p)) spectra of neutral and cationic adamantane (top panel) and singly dehydrogenated cationic adamantyl (middle panel). The measured spectrum of FUV-irradiated adamantane isolated in solid Ne (6.8 K) is displayed in the bottom panel. Note that the calculated spectra are just representations of the calculated excited states (stick spectra in top and middle panels). They have been computed by convolution using Lorentzians with a width of 3000 cm$^{-1}$ to indicate the range where strong electronic transitions can be expected.} \\label{fig3} \\end{center}\\end{figure} As indicated in Figure \\ref{fig3}, an additional baseline correction was applied to remove the strong scattering background in the UV. Because of its wavelength dependence ($\\sim \\lambda^{-4}$), we attribute this background mainly to an increased Rayleigh scattering of the charged species compared to their neutral counterparts. The derived absorption spectrum consists of four broad features above 200 nm. The broad band between 280 and 350 nm with maximum around 308 nm is not an artifact of the measurement. It has been confirmed by repeatedly performed experiments. The same applies to the two somewhat narrower features at 252 and 261 nm. The strongest band in the accessible wavelength region extends from 200 to 240 nm and peaks at 223.5 nm. Considering the theoretical results, we exclude the presence of open-shell adamantane cations in the photo-processed matrix due to the absence of bands in the visible. Instead, the observed spectrum points toward the creation of the more stable, closed-shell, singly de-hydrogenated cation. Therefore, we assign the strongest band at 223.5 nm to the S$_0 \\rightarrow$ S$_2$ transition of the 1-adamantyl cation (point group C$_{3\\text{v}}$) which is the isomer with lower ground-state energy. The calculated oscillator strength of this band is $f = 0.091$. Its position is 0.3 eV away from the measured value in Ne matrix which is a reasonable error for the applied theoretical model. The other weaker transitions may be related to the S$_0 \\rightarrow$ S$_1$ transition of the 1-adamantyl cation ($f = 0.0028$) and, more likely, to the first four excited states of the 2-adamantyl cation (point group C$_{\\text{s}}$) with calculated oscillator strengths below $f = 0.033$. A more reliable assignment on the basis of computed spectra would be possible via IR spectroscopic investigations, because IR-active vibrations can be more easily and accurately calculated using quantum chemical models. However, this is beyond the scope of the present investigation. Nevertheless, taking into account the weaker band strengths of the 2-adamantyl cation and comparing the calculated with the measured spectra, it seems that the second isomer is almost as abundantly created as the 1-adamantyl cation. The bottom panel of Figure \\ref{fig3} contains a sum spectrum of both isomers (each contributing 50 \\%). Besides the 0.3 eV redshift of the 1-adamantyl band, it closely resembles the measured spectrum. To some extent, this would be in contradiction to the results obtained by \\citet{polfer04} who used an indirect charge transfer method to create the ions and subsequently observed only the isomer with the lowest ground-state energy. A possible explanation may be found in the different experimental techniques (charge transfer versus photoionization) and conditions (gas phase versus matrix at low temperature) that were applied. Considering the previous results, the adamantane molecules were subjected to dissociative photoionization upon irradiation with FUV photons of energy 10.2 $-$ 11.8 eV according to \\begin{equation} \\text{C}_{10}\\text{H}_{16} + h \\nu \\rightarrow \\text{C}_{10}\\text{H}_{15}^{+} + \\text{e}^{-} + \\text{H}. \\end{equation} This process has been described for a few smaller molecules, e.g., H$_2$O \\citep{cairns71} or CH$_4$ \\citep{samson89}. After ionization, some excess energy is stored in the vibrational degrees of freedom of the cationic molecule which subsequently leads to the destruction of one terminal C$-$H bond. During this process, electrons and neutral H atoms are released which are usually trapped on defects or impurities in the matrix. Because of recombination reactions between positively charged molecules and released electrons (and H atoms), the ion yield saturates when a certain irradiation dose is reached. The formation of negatively charged adamantane or neutral adamantyl molecules due to electron attachment can be ruled out for the following reasons. First, these species possess an open-shell electronic structure and, like the adamantane cation, would feature absorption bands in the visible part of the spectrum. And second, as far as neutral adamantane is concerned, the negative electron affinity \\citep{drummond07} hampers further electron attachment. Notably, we observed neither sharp absorption bands nor a clear vibrational pattern. We want to remark that the measured bands are much broader than what is expected for typical matrix-induced broadening (at least for 7 K neon matrices) which is mainly due to site effects.\\footnote{Molecules in different sites of the matrix exhibit different redshifts of their absorption bands, effectively resulting in a broadening.} We tentatively attribute this to an intrinsic property of the molecule, i.e., a very short lifetime of the excited state which is not entirely caused by the interaction with the rare gas atoms. This would lead to the conclusion that the main difference to astrophysically more relevant spectra of cold gas phase adamantyl cations is a small matrix-induced redshift, but not a broadening of the absorption bands. By coincidence, the spectral shape is in surprisingly good agreement with the computed spectra of the purely electronic (vertical) transitions. Therefore, the absolute values of the absorption cross sections, as they appear in the calculated spectra, may be regarded as representative for real gas phase molecules.\\footnote{Otherwise, only the integrated cross section or the oscillator strength could be taken.} Finally, we want to highlight another important property of the adamantyl cations. Unlike their neutral precursor adamantane, these species possess rather strong permanent dipole moments pointing from the molecular center toward the C atom from which the H atom has been removed. The calculated dipole moments of the 1- and 2-adamantyl cations amount to 0.96 and 2.57 Debye, even stronger than the dipole moment of the open-shell cation (0.61 Debye). This opens the possibility to detect and identify molecular diamond-like species in space using their rotational spectra. \\subsection{Diamantane} The calculated ground-state energies of diamantane and its singly dehydrogenated cationic derivatives are displayed in Figure \\ref{fig4}. Again, the numbers on the diamantane structure shown indicate the positions from where the hydrogen atoms are removed to form the corresponding closed-shell cations. There are three possible isomers for the diamantyl cation, possessing quite strong permanent dipole moments of 1.76 (D1), 3.09 (D4), and 3.71 Debye (D3) due to the localized charge at the edge of the molecule. \\begin{figure}[t]\\begin{center} \\epsscale{1.0} \\plotone{f4.eps} \\caption{Zero-point-corrected ground-state energies of diamantane and its derivatives calculated at the B3LYP/6$-$311$+$G(d) level of theory. The diamantyl structures are labeled according to the IUPAC numbering system for diamondoids.} \\label{fig4} \\end{center}\\end{figure} The electronic spectra of diamantane and its related species are presented in Figure \\ref{fig5}. An additional background correction has been applied on the red side of the measured spectrum to remove non-reproducible bumps and fringes due to baseline variations. Using 1.2 g cm$^{-3}$ \\citep{karle65} as mass density of the solid diamantane deposit, the isolation ratio (Ne to diamantane) varied in different experiments between 450 $n_{\\text{Dia}}$ $n^{-1}_{\\text{Ne}}$ and 750 $n_{\\text{Dia}}$ $n^{-1}_{\\text{Ne}}$, where $n_{\\text{Dia}}$ is the refractive index of the diamantane film. Compared to the photoprocessed adamantane, the spectrum of irradiated diamantane reveals a slightly broader peak (7800 cm$^{-1}$ versus 5900 cm$^{-1}$), positioned further to the red at 255 nm. As is evident upon inspection of the calculated and measured spectra, this feature cannot be explained by the presence of open-shell cation radicals. Regarding the formation of negatively charged diamantane or neutral diamantyl radicals, the same reasoning applies as in Section \\ref{adam_section}. Principally, the photons from the hydrogen lamp carry enough energy to induce dissociative photoionization and create all three diamantyl isomers. However, the conclusion may be drawn that the main photoproduct in the matrix experiment is the 4-diamantyl cation (point group C$_{\\text{3v}}$). Its first strong transitions S$_0 \\rightarrow$ S$_{2,3}$ are predicted at 284 nm ($f=0.084$) and 269 nm ($f=0.033$), 0.5 and 0.26 eV away from the peak maximum of the measured band at 255 nm. Alternatively, the 255 nm band may partly or completely originate from transitions caused by the 1-diamantyl cation (C$_\\text{s}$) for which several close-lying absorptions at 289 nm (S$_0 \\rightarrow$ S$_3$, $f=0.022$), 271 nm(S$_0 \\rightarrow$ S$_4$, $f=0.027$), and 266 nm (S$_0 \\rightarrow$ S$_5$, $f=0.003$) are predicted. Due to the lack of certain absorption features in the measured spectrum, the presence of the 3-diamantyl isomer (C$_1$) in the matrix, i.e. the H removal from a CH$_2$ group, can rather be excluded. Because of the apparent absence of fine structure, it seems difficult to draw further conclusions. \\begin{figure}[t]\\begin{center} \\epsscale{1.15} \\plotone{f5.eps} \\caption{Calculated (B3LYP/6$-$311$+$G(d)) spectra of neutral and cationic diamantane (top panel) and singly dehydrogenated cationic diamantyl (middle panel). The measured spectrum of FUV-irradiated diamantane isolated in solid Ne (6.8 K) is displayed in the bottom panel.} \\label{fig5} \\end{center}\\end{figure} \\subsection{Triamantane} Figure \\ref{fig6} displays the calculated ground-state energies of triamantane and its seven isomers of singly dehydrogenated cations. The necessary energies to remove one electron and one H atom from the parent molecule are comparable to the previously discussed diamondoids. Basically, the FUV lamp delivers photons with energies high enough to create all seven isomers. Their dipole moments, again quite strong, range between 1.13 and 5.72 Debye. The dipole moment of the open-shell cation amounts to 0.51 Debye. \\begin{figure}[t]\\begin{center} \\epsscale{1.0} \\plotone{f6.eps} \\caption{Zero-point-corrected ground-state energies of triamantane and its derivatives calculated at the B3LYP/6$-$311$+$G(d) level of theory. The triamantyl structures are labeled according to the IUPAC numbering system for diamondoids.} \\label{fig6} \\end{center}\\end{figure} The corresponding calculated electronic spectra, as well as the measured spectrum of FUV-irradiated triamantane in Ne are displayed in Figure \\ref{fig7}. The isolation ratio in the matrix experiment was on the order of 500 $-$ 1000 $n_{\\text{Tria}}$ $n^{-1}_{\\text{Ne}}$ with unknown refractive index of the pure triamantane ($n_{\\text{Tria}}$) film. The photodissociation of trace amounts of water in the matrix is responsible for the narrow bands from the OH radical at 308 and 283 nm \\citep{tinti68}. Compared to the previous measurements of irradiated adamantane and diamantane, the lowest-energy spectral feature shifts further to the red, extending roughly from 300 to 500 nm. It peaks at 368 nm and has a shoulder around 450 nm. A strong FUV rise also slides into the accessible wavelength region. An assignment to a certain isomer of the triamantyl cation is rather difficult. Oddly, the best match seems to be possible with the calculated spectrum of the 5-triamantyl cation which is 10.56 eV higher in energy than the triamantane neutral. (In the previously discussed measurements, the strongest bands seemed to be caused by species which were 10.35 eV (adamantane) and 10.27 eV (diamantane) away from the parent molecule.) Besides dehydrogenated triamantyl cations, an alternative explanation for the measured broad band would be the formation of the open-shell triamantane cation as can be seen from the comparison with the corresponding calculated spectrum. Possibly, triamantane is already big enough, and the energy stored in the molecule upon absorption of an FUV photon is distributed over sufficient vibrational modes to avoid H abstraction. Nevertheless, a clear identification of the created species on the basis of electronic absorption spectroscopy is not possible and one should take into account the possibility that the applied quantum chemical model deviates more strongly from reality than expected. \\begin{figure}[t]\\begin{center} \\epsscale{1.15} \\plotone{f7.eps} \\caption{Calculated (B3LYP/6$-$311$+$G(d)) spectra of neutral and cationic triamantane (top panel) and singly dehydrogenated cationic triamantyl (two middle panels). The measured spectrum of FUV-irradiated triamantane isolated in solid Ne (6.8 K) is displayed in the bottom panel.} \\label{fig7} \\end{center}\\end{figure} \\subsection{Tetramantane} In contrast to the smaller diamondoids, there are already three different isomers of neutral tetramantane C$_{22}$H$_{28}$, while one of them has actually two enantiomers (P and M [123]-tetramantane). Their structures, the calculated spectra of their cations, as well as the measured spectra of their photoproducts, are displayed in Figure \\ref{fig8}. We did not calculate the structures and spectra of the tetramantyl cations because of the increasing number of isomers and therefore escalating computational effort. Regarding the necessary energies for H abstraction, we do not expect large deviations from the smaller diamondoids. With almost no difference among the three species, the measured spectra are very similar to what has been measured for triamantane. Besides [121]-tetramantane, the broad feature extending roughly from 300 to 500 nm could be assigned to the open-shell cations as is obvious upon comparison with the calculations. The [121]-tetramantane cation, however, should have stronger bands at longer wavelengths suggesting that the measured spectrum is actually caused by the corresponding singly dehydrogenated cation.\\footnote{On the other hand, it should be taken into account that the rise in the baseline beyond 510 nm could indicate a very broad and, compared to the calculated spectrum, rather weak band around 570 nm.} Whether the other two tetramantanes ([123] and [1(2)3]) lost a peripheral H atom upon ionization or not, cannot completely be clarified, as there are too many tetramantyl isomers and, as is the case with triamantane, a comparison with TD-DFT theory would not provide unambiguous insights. \\begin{figure}[t]\\begin{center} \\epsscale{1.15} \\plotone{f8.eps} \\caption{Isomers of tetramantane, calculated spectra of their cations, and measured spectra of their photoproducts isolated in solid Ne (6.8 K).} \\label{fig8} \\end{center}\\end{figure} \\subsection{Complete $\\sigma - \\sigma^*$ Absorption Spectra} \\begin{figure*}[t]\\begin{center} \\epsscale{1.15} \\plotone{f9.eps} \\caption{Complete calculated electronic $\\sigma - \\sigma^*$ absorption spectra of neutral and ionized small diamondoids. For comparison, the bottom right panel contains the IR to VUV spectrum of meteoritic nanodiamonds from the Allende meteorite ($\\sim$ 2 nm) as derived from combined absorption and EELS (electron energy loss spectroscopy) measurements \\citep{mutschke04}. The IR bands of the nanodiamonds are actually too weak compared to the electronic absorption to be seen in this scale.} \\label{fig9} \\end{center}\\end{figure*} The electronic spectra of the neutral and cationic diamondoids resulting from (all possible) bound-bound ($\\sigma - \\sigma^*$) transitions, as calculated with the Octopus code, are displayed in Figure \\ref{fig9}. These spectra may be used to model photophysical interactions of diamond-like molecules in the interstellar medium. For a brief discussion about their physical relevance refer to Section \\ref{theory}. Regarding the energy range below 8.5 eV, more detailed, measured gas phase absorption spectra of neutral diamondoids can be found in the publication of \\citet{landt09b}. The vibrational structure that can be seen in these spectra can hardly be predicted with current theoretical methods. The spectra presented here feature broader bands that are purely artificial. An additional energy-dependent broadening by convolving the spectra with Lorentzians of increasing bandwidths has been applied to account for an increased lifetime broadening which is expected at higher energies. Due to the high density of states above $\\sim$ 10 eV, changing the bandwidths does not substantially alter the absolute cross section values. While the positions and shapes of resonances appearing in the spectra may be affected by uncertainties of the computational method the general trend of the absorption curves and the cross section values may be regarded as real (at least within the limitations discussed before). Comparing the spectra of the neutral and ionized molecules with each other, it is obvious that there are not many differences, especially for the transitions at higher energies, as the electronic structure of the C skeleton is equivalent. For all species, the high-energy absorption is dominated by a broad hump with a maximum between 15 and 20 eV. Also other features on the red and blue tails of this $\\sigma - \\sigma^*$ hump, like peaks at 11, 14, and 28.5 eV for the cations or 9, 12.5, and 27.5 eV for the neutrals are very much comparable. As already discussed in the previous sections, the absorption onset of the cations appears further to the red compared to their neutral precursors which is also evident from Figure \\ref{fig9}. By increasing the molecular size, two effects are obvious: rising values for the absolute absorption cross section and a trend toward a less-structured absorption curve due to an increased density of states. We want to point out an interesting aspect of these results. Even though laboratory experiments in the discussed energy range for these molecular species are lacking, experimental data on nanodiamonds, extracted and isolated from the Allende meteorite \\citep{mutschke04}, display surprising resemblances (see Figure \\ref{fig9}). These nanodiamonds possess an average size of less than 2 nm, corresponding to $\\approx$ 500 C atoms, which is much bigger than the molecular diamonds presented here, the largest of which contains 22 C atoms. Their electronic absorption spectra solely consist of the broad $\\sigma - \\sigma^*$ band with maximum at 17.1 eV. Furthermore, a shoulder can be seen around 30 eV which may have its equivalent in a band at 28.5 eV in the diamondoid spectra. The peak mass absorption coefficient $\\kappa$ for the meteoritic nanodiamonds was found to be $1.1 \\times 10^6 \\text{ cm}^2\\text{ g}^{-1}$, very close to the calculated values of $\\kappa = 1.2 - 1.3 \\times 10^6 \\text{ cm}^2 \\text{ g}^{-1}$ for the molecular diamond. The main differences of the molecular compared to the nanoscopic material are the redshifted absorption onset and a more structured absorption curve, especially on the red wing of the collective $\\sigma - \\sigma^*$ hump." }, "1101/1101.4649.txt": { "abstract": "We present the first characterization of the excess continuum emission of accreting T Tauri stars between optical and near-infrared wavelengths. With nearly simultaneous spectra from 0.48 to 2.4~\\micron\\ acquired with HIRES and NIRSPEC on Keck and SpeX on the IRTF, we find significant excess continuum emission throughout this region, including the $I$, $Y$, and $J$ bands, which are usually thought to diagnose primarily photospheric emission. The $IYJ$ excess correlates with the excess in the $V$ band, attributed to accretion shocks in the photosphere, and the excess in the $K$ band, attributed to dust in the inner disk near the dust sublimation radius, but it is too large to be an extension of the excess from these sources. The spectrum of the excess emission is broad and featureless, suggestive of blackbody radiation with a temperature between 2200 and 5000 K. The luminosity of the $IYJ$ excess is comparable to the accretion luminosity inferred from modeling the blue and ultraviolet excess emission and may require reassessment of disk accretion rates. The source of the $IYJ$ excess is unclear. In stars of low accretion rate, the size of the emitting region is consistent with cooler material surrounding small hot accretion spots in the photosphere. However, for stars with high accretion rates, the projected area is comparable to or exceeds that of the stellar surface. We suggest that at least some of the $IYJ$ excess emission arises in the dust-free gas inside the dust sublimation radius in the disk. ", "introduction": "Classical T Tauri stars (CTTS) are low-mass stars in the final stages of disk accretion. Accretion from the disk to the star is thought to be governed by the stellar magnetosphere, where a sufficiently strong magnetic field truncates the disk at several stellar radii and guides infalling material along field lines to the stellar surface at high latitudes, terminating in accretion shocks. Historically, a major factor contributing to the understanding that T Tauri stars are accreting matter from their circumstellar disks was the spectrum of their excess continuum emission from ultraviolet to radio wavelengths \\citep{ber89}. While the disks themselves are the primary source of continuum radiation at wavelengths longward of 2~\\micron, optical and ultraviolet continuum emission in excess of the photosphere is attributed to accretion shocks on the stellar surface. Spectral lines also played a key role in developing the current magnetospheric accretion paradigm, providing the kinematic evidence that the stellar magnetosphere channels accreting material from the disk to the star and that accretion-powered winds arise from near-stellar regions \\citep{naj00}. In this paper we focus on the excess continuum emission in a region that has not received close scrutiny to date -- the spectral region between the optical and the near infrared, where the photospheric emission from these late-type stars peaks. The well studied optical and ultraviolet excesses are used to derive a fundamental parameter of T Tauri stars -- the rate of accretion from the disk to the star. The inferred accretion rates span several orders of magnitude, from extremes of $10^{-10}$ to $10^{-7}$ \\msunyr, with a median of $\\sim 10^{-8}$ \\msunyr\\ at 1 Myr that declines with increasing age \\citep{har98}. The most reliable estimates for disk accretion rates come from comparisons of spectrophotometric observations over a broad wavelength range to shock models. This was done successfully by Calvet and Gullbring (1998, hereafter \\citealt{cal98}), who accounted for the excess continuum between 0.32 and 0.52~\\micron\\ with shock models that attribute the Paschen continuuum to optically thick post-shock gas in the heated photosphere and the Balmer continuum and Balmer jump to optically thin gas in the pre-shock and attenuated post-shock regions. This combination of spectrophotometry with shock models gives a fairly uniform result for the spectrum of the excess in the Paschen continuum, with temperatures $\\sim6000-8000$ K and a shape that is essentially blackbody. A more commonly used approach to determine accretion rates is to evaluate the excess emission over a limited range of wavelength in the Paschen continuum by comparing the depth of photospheric absorption lines to those of a template matched in temperature, gravity, and projected rotational velocity. In the presence of a continuum excess, photospheric features will be weakened, and a quantity known as {\\em veiling}, defined as the flux ratio $r=F_{\\rm excess}/F_*$, can be derived \\citep{bas90,har91}. Accretion rates are then determined from a bolometric correction based on an isothermal slab model at an assumed temperature, density, and optical depth (\\citealt{val93,har95}; Gullbring et al.\\ 1998a, hereafter \\citealt{gul98a,har03,her08}). At near-infrared wavelengths, CTTS show excess emission from 2 to 5 \\micron\\ that is well described by $\\sim1400$~K blackbody radiation and is attributed to a raised rim of dust at the dust sublimation radius in the inner disk (Muzerolle et al.\\ 2003, hereafter \\citealt{muz03,fol01,joh01}). The magnitude of the excess is proportional to the accretion luminosity, requiring that the dust be heated by radiation from both the photosphere and the accretion shock. $K$-band interferometry, which locates dust in a ring at a few tenths of an AU from the star \\citep{eis05, mil07}, provides further evidence that 2 to 5 \\micron\\ emission in CTTS comes from sublimating dust in the inner disk. Evidence has been accumulating that there is an unexplained source of excess emission in CTTS between the optical and near infrared that cannot be attributed to shocks at the base of magnetic funnel flows or dust sublimating in the disk. As early as 1990, \\citeauthor{bas90} found that the veiling longward of 0.5 \\micron\\ does not steadily decline with wavelength as expected from hot accretion shock models but instead is relatively constant between 0.5 and 0.8 \\micron, and high veiling at 0.80--0.85 \\micron\\ has more recently been reported by \\citet{har03} and \\citet{whi04}. Furthermore, \\citet{edw06} found high veiling at 1~\\micron, which is unlikely to come from 1400 K dust in the inner disk since this emission will fall rapidly to short wavelengths from its peak around 3~\\micron. Also, $K$-band interferometric studies of T Tauri stars find the angular size of the near-infrared emission in modeling of multiple-baseline observations \\citep{ake05} and the size-wavelength behavior in spectrally dispersed observations \\citep{eis09} to suggest that gaseous material inside the dust sublimation radius is present at a temperature higher than that of the dust. However, the lack of a systematic study of the wavelength dependence of excess emission between 0.5 and 2 \\micron\\ has inhibited the understanding of its role in accreting systems. The presence of unexplained continuum emission in CTTS between 0.5 and 2 \\micron\\ not only offers the opportunity to improve our understanding of the structure of accretion disk systems; it also poses a practical problem. If we do not understand the shape of the veiling spectrum at wavelengths longer than 0.5 \\micron, then we cannot reliably convert veiling measurements at these wavelengths via a simple bolometric correction into accurate accretion luminosities and disk accretion rates. There may also be ramifications for extinction determinations of CTTS if they are based on far-red colors assumed to be predominantly photospheric. In this paper we determine the excess spectrum over a broad wavelength region between 0.48 and 2.4 \\micron\\ for a sample of 16 classical T Tauri stars in Taurus. The presentation includes \\S\\ 2 describing the sample and data reduction, \\S\\ 3 describing the derivation of the veiling and the excess emission spectra, \\S\\ 4 characterizing the behavior of the excess emission spectra, and \\S\\ 5 describing simple models for the excess, followed by a discussion in \\S\\ 6 and conclusions in \\S\\ 7. We demonstrate that the two traditional sources of optical and near-infrared excess emission, a hot component arising from an accretion shock-heated photosphere with small filling factor and a cool component arising from the dust sublimation radius of the accretion disk, are not sufficient to describe the observed excess. A third component of intermediate temperature seems to be required to fit the excess emission in this region, with a luminosity of the same order of magnitude as derived from the hot shock-heated gas. ", "conclusions": "We have characterized the excess continuum emission in accreting T Tauri stars between 0.48 and 2.4 \\micron. At 1 \\micron\\ the magnitude of the excess ranges from 0.1 to 3 times the photospheric flux, and it scales with the excess at shorter and longer wavelengths. We interpret this excess as arising from a source of continuum emission in addition to small hot ($T\\sim8000$ K) accretion spots on the stellar surface at the base of accretion columns and large rings of warm ($T\\sim1400$ K) dust at the sublimation radius in the disk. The broad, smooth distribution of the $IYJ$ excess suggests it comes from a third component with a temperature between 2200 K and 5000~K and roughly blackbody emission characteristics. We cannot pin down a unique combination of temperature and filling factor to account for the emission, but among stars with high $IYJ$ excess, the filling factor is comparable to the surface area of the star if the temperature is 5000~K or more than an order of magnitude larger if the temperature is 2500 K. Possible sources of the emission include warm annuli surrounding hot accretion spots in the shock-heated photosphere, accreting gas in funnel flows, or disk gas inside the dust sublimation radius. The luminosity from this region is comparable to the accretion luminosity found from ultraviolet and blue excesses, suggesting that accretion luminosities have been underestimated by a factor of about two. The implication for disk accretion rates is more complicated, but a comprehensive campaign to obtain simultaneous spectrophotometry from the ultraviolet to the near infrared for a large sample of CTTS would provide further insight. Whatever its source, this excess emission contains a large fraction of the accretion energy released in the vicinity of the star. \\vspace{-0.1in}" }, "1101/1101.0832_arXiv.txt": { "abstract": "We present analysis of high-resolution spectra of a sample of stars in the globular cluster M5 (NGC 5904). The sample includes stars from the red giant branch (seven stars), the red horizontal branch (two stars), and the asymptotic giant branch (eight stars), with effective temperatures ranging from 4000 K to 6100 K. Spectra were obtained with the HIRES spectrometer on the Keck I telescope, with a wavelength coverage from 3700 \\AA{} to 7950 \\AA{} for the HB and AGB sample, and 5300 \\AA{} to 7600 \\AA{} for the majority of the RGB sample. We find offsets of some abundance ratios between the AGB and the RGB branches. However, these discrepancies appear to be due to analysis effects, and indicate that caution must be exerted when directly comparing abundance ratios between different evolutionary branches. We find the expected signatures of pollution from material enriched in the products of the hot hydrogen burning cycles such as the CNO, Ne$-$Na, and Mg$-$Al cycles, but no significant differences within these signatures among the three stellar evolutionary branches especially when considering the analysis offsets. We are also able to measure an assortment of neutron-capture element abundances, from Sr to Th, in the cluster. We find that the neutron-capture signature for all stars is the same, and shows a predominately r-process origin. However, we also see evidence of a small but consistent extra $s$-process signature that is not tied to the light-element variations, pointing to a pre-enrichment of this material in the protocluster gas. ", "introduction": "Among globular clusters of the northern sky, M5 is one of the nearest, and the element abundance patterns among its member stars have received considerable attention. On the basis of its observed proper motion, M5 actually appears to be an outer halo globular cluster on an eccentric orbit with a large apogalactic distance of $\\sim 60$ kpc \\citep{scholz96}. It is one of the most metal-rich globular clusters of the outer Galactic halo, with ${\\rm [Fe/H]} = -1.34 \\pm 0.09$ \\citep{carretta09c}. M5 was one of the first globular clusters in which a sub-population of red giant branch (RGB) stars whose spectra exhibit enhanced $\\lambda$4215 CN bands were discovered via DDO photometry \\citep{osborn71, hesser77, pike78}. The CN anomalies in M5 have been traced lower down the giant branch \\citep{briley92} and to the base of the RGB \\citep{cbs02}. Abundance variations of O, Na, and Al also exist among the RGB stars (e.g. \\citealt{norris83,ivans01, yong08, yong08b, carretta09a, carretta09b}), and inhomogeneities in Na abundance have been traced from the tip of the RGB to the main sequence turnoff by \\citet{rc03}. In all these respects, the abundance inhomogeneities found among RGB stars in M5 appear to be typical of the broad patterns found in other globular clusters of the Milky Way (e.g, \\citealt{carretta04, gratton01, bragaglia10}) and also the Local Group \\citep{mucciarelli09}. However, abundance anomalies among the asymptotic giant branch (AGB) stars of globular clusters, including M5, have not been as well studied as those on the RGB. In color-magnitude diagrams (CMDs) of M5, the loci of the RGB and AGB are relatively clearly separated (see, for example, \\citealt{simoda70, buonanno81, sandquist96, sandquist04}), making it a particularly useful cluster for studying AGB stars. \\citet{zinn77} classified a number of AGB stars as having very weak $G$-bands, and subsequently \\citet{smith93} found that a substantial fraction of AGB stars in M5 have enhanced CN band strengths. The presence of CN-strong stars on the AGB of various globular clusters has been reviewed by \\citet{snedenAGB} and \\citet{campbell96}, based on the relatively sparse literature available. Table 1 in \\citet{campbell96} suggests that in clusters more metal-poor than M5 the AGB stars tend to have weak CN bands. \\citet{campbell10} confirmed a relatively large population of both CN-weak and CN-strong stars on the AGB of M5. This cluster therefore offers an opportunity for a more extensive study of N-Na-Mg-Al element enhancements on the AGB. Two general mechanisms have been proposed to explain these abundance patterns (e.g., \\citealt{kraft94}). The first is that the surfaces of these stars are polluted during the RGB phase by the interior products of proton-capture reactions which have been consequently mixed to the surface. The second is that the stars with high N, Na, and Al and low O and C are part of a second generation of stars, formed out of gas ejected by polluters in which hot H burning took place. The nature of the polluters is still debated, including intermediate-mass AGB stars from $\\sim$4-8 $M_{\\odot}$ \\citep{ventura09}, fast-rotating massive stars with 20-60 $M_{\\odot}$ \\citep{decressin07}, and massive binaries with $\\sim$ 20 $M_{\\odot}$ \\citep{demink09}. The discovery that these abundance patterns continue to at least the main-sequence turnoff indicates that these must be second-generation stars with the abundance anomalies throughout the entire star. However, the surface abundances of some elements and isotopes are affected as stars go through the later phases of evolution (e.g., \\citealt{gratton00,sm03}). In addition to the first dredge-up on the low RGB, ``deep mixing'' or ``extra mixing'' on the upper reaches of the RGB causes C and Li abundance drops, N increases, and $^{12}$C/$^{13}$C ratio decreases which were not predicted in original stellar models. This requires an additional physical effect that has not been conclusively identified. Possibilities include magnetic buoyancy (e.g. \\citealt{busso07, denissenkov09}) and mean molecular weight gradients which lead to ``thermohaline'' mixing \\citep{eggleton06, eggleton08, charbonnel07}. Part of the uncertainty lies in the efficiency of the mixing by either mechanism. For example, \\citet{charbonnel07} showed that thermohaline mixing could account for the abundance patterns on the RGB if the efficiency for mixing was high, but two-dimensional simulations by \\citet{denissenkov10} of thermohaline mixing found that the actual efficiency was much smaller, closer to the magnitude in \\citet{kippenhahn80}. There is also an ongoing debate about whether extra mixing happens on the AGB. Models without AGB extra mixing may have difficulty explaining observations of C/N and $^{12}$C/$^{13}$C ratios in AGB stars (e.g., \\citealt{lambert86, lebzelter08, milam09}) and O isotope ratios in pre-solar grains (e.g, \\citealt{hoppe97}). \\citet{karakas10} argued that if extra mixing on the RGB was included in the models, then no extra mixing on the AGB was needed to explain the C/N and C and O isotope ratios dredged up to the surface and observed in stars, at least at solar metallicities. \\citet{busso10} found, however, that extra mixing in AGB stars was necessary to match isotope ratios in pre-solar grains, along with the C isotope ratios in C(N) stars, even if extra-mixing was included on the first ascent RGB. However, the mechanism for this extra mixing, like its counterpart on the RGB, is not yet known. If deep mixing can also occur in AGB stars, then there is the possibility of CNO abundance differences being produced between the RGB and AGB. Recently, stellar models have been evolved from the main-sequence to the thermally pulsing AGB that includes mechanisms for mixing and extra mixing to trace the evolution of surface abundances in low-mass stars. \\citet{stancliffe10} focused on low-metallicity stars. Thermohaline efficiency was adopted from \\citet{charbonnel07} and is therefore very high. They found that $^3$He is not all depleted, so thermohaline mixing could persist on the AGB. In addition, on the early AGB, the deepening of the convective envelope changes slightly the surface Li and $^3$He abundances and the C isotope ratios. In summary, theoretical work shows there are potentially interesting changes in the light elements resulting from mixing and extra mixing throughout the RGB and AGB. Observational evidence of the existence and size of these effects will constrain the mechanism of mixing and its efficiency. The AGB stars of M5 are cleanly separated from RGB stars, and the stars at the tip of the RGB provide a good reference for light-element abundances that may change on the AGB. In this paper we present measurements of light-element abundances for stars on both the RGB and AGB of M5 in an effort to determine whether there is any variation with stellar evolution. The star-to-star inhomogeneities in the light elements of M5, however, do not appear to extend to the heavy elements formed by neutron-capture processes. This suggests that their production is divorced from the nucleosynthesis in the self-polluting cluster stars that made the second generation stars. The first in-depth study of Ba, La, and Eu in M5 from \\citet{ivans01} found small internal scatter and good overall agreement with halo field subdwarfs with similar [Fe/H]. Interestingly, \\citet{ivans01} noted that the abundances in M4 \\citep{ivans99} showed enhanced $s$-process contributions from AGB stars that enriched the natal gas of all stars in the cluster. \\citet{yong08b} measured 27 elements heavier than Fe in two RGB stars in M5 and 12 RGB stars in M4. In addition to confirming the differences found by \\citet{ivans01}, they found that the abundance ratios could be explained by some $s$-process in M5 as well, though at a much smaller fraction than M4. No freshly $s$-processed material is expected to appear on the surface of the present-day M5 AGB stars, because third dredge-up does not occur for stars with $M< 1.5M_{\\odot}$. However, as is clear from the discussion of extra mixing, we do not fully understand the possible mixing events that can occur outside of the long-established dredge-up events, and the $s$-process elements for stars with a range of evolutionary states in M5 provide an opportunity to test models. For example, \\citet{masseron06} suggest that the $s$-process enhancements seen in the extremely metal-poor AGB star CS 30322-023 are the result of an unknown mixing process that has brought this just-produced material to the surface. In this study, we also explore the origin of the neutron-capture elements in M5, and if there are any signatures of an $s$-process contribution. ", "conclusions": "We have performed detailed abundance analyses for a sample of evolved stars in the globular cluster M5, covering the RGB, RHB, and AGB branches of the CMD. Our conclusions can be broken into three parts. The first is a cautionary note. It appears that there can be systematic abundance offsets induced when analyzing stars on the different evolutionary branches using standard onedimensional LTE abundance analysis and atmospheres. This can manifest itself not only in absolute abundances, but more worryingly in abundance ratios such as [Ca/Fe]. The largest offset found was in [V/Fe], with a 0.15 dex difference between the AGB versus RGB stars. This puts a limit on how well abundances among the different evolutionary branches of M5 can be compared. The second conclusion is that our sample clearly shows the signatures of star-to-star abundance differences related to the CNO, Ne$-$Na, and Mg$-$Al nuclear reaction cycles, and that there are no discernible differences in these element patterns with stellar evolution phase. This agrees with theoretical predictions that self-pollution and mixing within present-day globular cluster stars will not begin until the thermal pulsation phase of the AGB. Taking this into context with the M5 turn-off stars analyzed in \\citet{rc03}, the (N-,Na-,Al)-rich cluster stars are present throughout all post-main-sequence phases of evolution. This is consistent with such stars having acquired their surface element enhancements from external sources, and not from the outwards transport of material processed through H-burning reactions within their interiors. The third conclusion is that we find the neutron-capture abundances of M5 to be $r$-process dominated, but with what we interpret as a small uniform addition of $s$-process material. This neutron-capture signature is constant through all stars in our sample, depending on neither evolution nor the light-element variations. This suggests that low-mass AGB stars contributed heavy elements to the primordial cluster environment. However, the lack of correlation with the light-element variations also seems to preclude low-mass AGB stars from having much of a contribution once star formation had begun." }, "1101/1101.1143_arXiv.txt": { "abstract": "The principal Hugoniot for liquid hydrogen was obtained up to 55 GPa under laser-driven shock loading. Pressure and density of compressed hydrogen were determined by impedance-matching to a quartz standard. The shock temperature was independently measured from the brightness of the shock front. Hugoniot data of hydrogen provide a good benchmark to modern theories of condensed matter. The initial number density of liquid hydrogen is lower than that for liquid deuterium, and this results in shock compressed hydrogen having a higher compression and higher temperature than deuterium at the same shock pressure. ", "introduction": "The properties of hydrogen at high pressure and high density are of great scientific interest. The equation of state (EOS) of hydrogen at these conditions is essential for modeling of the interior structure of gas giant planets \\cite{saumon04,nettelmann08,militzer08}. The large diversity in the estimation of Jupiter's core mass is resulted from the uncertainty in the EOS data especially in the region around the insulator-to-metal transition. The EOS of hydrogen-isotopes has important practical applications for inertial confinement fusion \\cite{lindl95}, and metallic hydrogen is suggested as a prospective candidate of high-temperature superconductor \\cite{ashcroft68}. Chemical free-energy models \\cite{saumon95,ross98,kerley03} and $ab$ $initio$ simulations \\cite{collins01prb,holst08} have been used to predict the properties of warm dense hydrogen, but the results vary widely and have not converged yet. Therefore accurate experimental data for the hydrogen EOS are required for evaluation of the theoretical models and for further understanding of the fundamental nature of hydrogen. It is more difficult to generate high pressures in hydrogen than in deuterium because of its lower shock impedance. For this reason, most of the recent experimental measurements by shock compression have focused on the heavier isotope \\cite{collins98,mostovych00,boehly04,knudson04,boriskov05,hicks09}. However, it should be noted that the Hugoniots for the two isotopes do not scale in density. Owing to the difference in zero-point energy, the mole volume of liquid hydrogen is larger than deuterium \\cite{wiberg55,kerley03}. As a result, hydrogen is expected to have higher compression and higher Hugoniot temperature than deuterium at the same pressure. There is a large gap in the experimental achievement of shock compression between liquid hydrogen and deuterium. The principal Hugoniot for liquid deuterium was measured up to 220 GPa using laser-driven shock waves \\cite{hicks09}. For the case of liquid hydrogen, the Hugoniot was studied experimentally only to 10 GPa by a gas gun and explosive method more than two decades ago \\cite{dick80,nellis83}. The metallization of hydrogen on the Hugoniot is expected to occur at much higher pressure. In this work, we carried out laser-shock experiments of liquid hydrogen to pressures exceeding 10 GPa in order to make a quantitative comparison of the hydrogen Hugoniot around the metal transition with the deuterium data. ", "conclusions": "Drude-type models are often applied to parameterize the optical properties of liquid metal \\cite{hodgson72,celliers00}. Within the Drude description, the complex index of refraction is given by $\\hat{n}_s^2 = 1 - (\\omega_p^2/\\omega^2)(1 + i / \\omega\\tau_e)^{-1}$ where $\\omega_p = (4 \\pi n_e e^2 / m_e)^{1/2}$ is the plasma frequency, $n_e$ is the carrier density, $e$ is the electron charge, $m_e$ is the electron mass, and $\\omega = 2 \\pi c / \\lambda$ is the optical frequency. Here the electron relaxation time is assumed to be $\\tau_e = R_0 / v_F$ where $R_0$ is the interparticle spacing and $v_F$ is the electron Fermi velocity. We adopt $R_0=0.126$ nm corresponding to a density $0.4$ g/cm$^3$. Then the Drude reflectivity of liquid metallic hydrogen can be derived as a function of $n_e$, which is depicted in Fig. 7. High reflectivity is produced when the carrier density exceeds the critical density $n_c$ defined by $\\omega^2_p (n_c) = \\omega^2$. The critical density is $n_c = 3.9 \\times 10^{21}$ cm$^{-3}$ at $\\lambda = 532$ nm. Taking $R = 23$\\%, for example, the carrier density is given by $3.7 \\times 10^{22}$ cm$^{-3}$, which is about 31\\% of total hydrogen number density $n_{\\max}=1.2 \\times 10^{23}$ cm$^{-3}$ and equivalent with $\\omega \\tau_e = 0.37$. The Fermi energy estimated from this carrier density is 4.0 eV, which is higher than the shock temperature. This indicates that the observed highly reflective state would be characteristic of a degenerate liquid metal. Based on the Drude model, the reflectivity increases with the carrier density and saturates at $R = 36$\\%. This implies that the maximum reflectivity of hydrogen may be slightly lower than that of deuterium \\cite{celliers00}, although it depends on the assumed compression ratio. It will be interesting to confirm the saturated value of the reflectivity by future experiments. The hydrogen Hugoniot obtained in this work is strongly dependent on the quartz EOS. Recently the quartz Hugoniot was reexamined by using magnetically driven flyer impact on the Z machine and a new fitting function was derived \\cite{knudson09}. For comparison, the hydrogen Hugoniots calculated by using the Z-fit and Kerley release ($U'_{p{\\rm H}}$, $P'_{\\rm H}$, and $\\rho'_{\\rm H}$) are listed in Table I. The stiffer Hugoniot of quartz reduces the initial pressure of the release isentrope (see Fig. 2). At the initial shock state of quartz, the pressure and density inferred from the Z-fit Hugoniot, $P_1 = 716$ GPa and $\\rho_1 = 6.61$ g/cm$^3$, are lower than those of the OMEGA-fit case, $P_1 = 746$ GPa and $\\rho_1 = 7.05$ g/cm$^3$). In our analysis, the quartz release cannot be approximated by a reflection of the Hugoniot in the $P$-$U_p$ plane, and thus the off-Hugoniot EOS of Kerley's model determines the shape of the curve. The particle velocity along the release isentrope can be calculated by using the Riemann integral: \\begin{equation} U_p = U_{p1} - \\int_{\\rho_1}^{\\rho} c_s \\rho^{-1} d{\\rho} \\;, \\label{eqn:riemann} \\end{equation} where $U_{p1}$ is the particle velocity at the initial shock state, and the sound speed is obtained from the pressure derivative of the density at constant entropy $S$, $c_s^2 = (\\partial P / \\partial \\rho)_S$. The second term of Eq. (\\ref{eqn:riemann}) is larger for the Z-fit case due to the lower density, and then the particle velocity increases faster with the decrease of density. Therefore, two release curves shown in Fig. 2 are gradually approaching at the lower pressure. The resulting differences in the Hugoniot pressure and density for hydrogen are $\\sim$ 1\\% and $\\sim$ 4\\% compared to those derived by the OMEGA-fit. Apparently, improvement of the quartz EOS is essential for the further development of Hugoniot measurements using quartz standards \\cite{boehly07}. In summary, we have obtained the principal Hugoniot data $P$-$\\rho$-$T$ for liquid hydrogen, not deuterium, in an unexplored range of pressure up to 55 GPa. The results demonstrate that the hydrogen Hugoniot cannot be scaled by density from the deuterium data in consequence of the initial density effects \\cite{hicks09,militzer06,eggert08}. As for the study of planetary interiors, the hydrogen EOS data at much higher pressure are required since the transition to metallic hydrogen is anticipated to be at $P \\sim$ 200-400 GPa in Jupiter \\cite{saumon04}. However, the hydrogen temperature must be kept lower because the Hugoniot temperature at this pressure range is too high to reproduce Jupiter's conditions. Therefore off-Hugoniot measurements of hydrogen by means of reflection shocks \\cite{boehly04} and/or precompressed samples \\cite{eggert08,loubeyre04,kimura10} will be a quite important next step." }, "1101/1101.4766_arXiv.txt": { "abstract": "\\os\\ absorption is observed in a wide range of astrophysical environments, including the Local ISM, the disk and halo of the Milky Way, high-velocity clouds, the Magellanic clouds, starburst galaxies, the intergalactic medium, damped \\lya\\ systems, and gamma-ray-burst host galaxies. Here a new compilation of \\ntot\\ \\os\\ absorbers drawn from the literature is presented, all observed at high resolution (instrumental FWHM$\\le$20\\kms) and covering the redshift range $z$=0--3. In galactic environments [log\\,$N$(\\hi)$\\ga$20], the mean \\os\\ column density is shown to be insensitive to metallicity, taking a value log\\,$N$(\\os)$\\approx$14.5 for galaxies covering the range --1.6$\\la$[O/H]$\\la$0. In intergalactic environments [log\\,$N$(\\hi)$<$17], the mean \\os\\ component column density measured in datasets of similar sensitivity shows only weak evolution between $z$=0.2 and $z$=2.3, but IGM \\os\\ components are on average twice as broad at $z$=0.2 than at $z$=2.3. The implications of these results on the origin of \\os\\ are discussed. The existence of a characteristic value of log\\,$N$(\\os) for galactic \\os\\ absorbers, and the lack of evolution in log\\,$N$(\\os) for intergalactic absorbers, lend support to the ``cooling-flow'' model of \\citet{He02}, in which all \\os\\ absorbers are created in regions of initially-hot shock-heated plasma that are radiatively cooling through coronal temperatures. These regions could take several forms, including conductive, turbulent, or shocked boundary layers between warm ($\\sim$10$^4$\\,K) clouds and hot ($\\sim$10$^6$\\,K) plasma, although many such layers would have to be intersected by a typical galaxy-halo sightline to build up the characteristic galactic $N$(\\os). The alternative, widely-used model of single-phase photoionization for intergalactic \\os\\ is ruled out by kinematic evidence in the majority of IGM \\os\\ components at low and high redshift. ", "introduction": "In a ground-breaking paper discussing the theoretical basis for the hot Galactic corona, \\citet{Sp56} correctly predicted that the lithium-like O$^{+5}$ ion ``might be sufficiently abundant to produce measurable absorption'' in the UV spectra of background sources, and could therefore be used to detect and analyze such a corona. \\os\\ absorption in its far-ultraviolet resonance doublet at 1031.926 and 1037.617\\,\\AA\\ is now routinely detected in the interstellar medium (ISM) of the Milky Way, the extended halos of other galaxies, and the intergalactic medium (IGM) from $z$=0 to $\\approx$3. The high cosmic abundance of oxygen, the intrinsic strength of the \\os\\ doublet, and the ability to probe warm/hot gas all combine to make \\os\\ a powerful and well-studied tracer of the diffuse Universe. It is common practice to study the detailed properties of \\os\\ absorption in a specific interstellar or intergalactic location. An alternative, global approach is to compare \\os\\ measurements from many different locations, and search for correlations with other observable parameters (such as redshift and metallicity). Along these lines, \\citet[][hereafter H02]{He02} synthesized \\os\\ observations across several low-redshift environments, and proposed a unified model in which all \\os\\ absorbers (both galactic and intergalactic) trace radiatively-cooling regions of initially-hot gas. In the last eight years, many new high-resolution (instrumental FWHM$\\le$20\\kms) \\os\\ datasets have become available. In order to draw them together, and to gain insight into the nature of \\os\\ absorbers, a new, heterogeneous compilation of \\ntot\\ \\os\\ absorbers is presented in this paper. The compilation includes galactic \\os\\ covering galaxy metallicities in the range $\\approx$0.01--1 times solar, and intergalactic \\os\\ covering redshifts from $z$=0 to $z\\!\\approx\\!2.5$. In \\S2 the relevant \\os\\ ionization physics is briefly reviewed. In \\S3, a survey of all published interstellar and intergalactic \\os\\ absorbers is presented. The new galactic and intergalactic compilations are presented in \\S4 and \\S5, respectively, together with a discussion of their properties. The key results are summarized in \\S6. ", "conclusions": "An extensive heterogeneous compilation of \\ntot\\ \\os\\ absorbers observed at high-resolution (instrumental FWHM$\\le$20\\kms) has been presented, covering the Local ISM, the disk and halo of the Milky Way, HVCs, the SMC, the LMC, starburst galaxies, the intergalactic medium from $z$=0 to $z\\!\\approx\\!2.5$, DLAs at $z$=2--3, and GRB host galaxies at $z$=2--4, using results drawn from the literature. This compilation is divided into a galactic sample [defined by log\\,$N$(\\hi)$>$17, and usually log\\,$N$(\\hi)$\\ga$20] consisting of \\ngal\\ \\os\\ measurements in galaxies with metallicities from $\\approx$0.01 solar to solar, and an intergalactic sample [log\\,$N$(\\hi)$<$17] consisting of \\nigm\\ \\os\\ components covering redshifts from $z$=0 to $z\\!\\approx\\!2.5$. The key observational findings are: \\begin{enumerate} \\item For \\os\\ measured in extended sight lines through {\\bf galactic} halos, the mean value of log\\,$N$(\\os) is surprisingly insensitive to metallicity (and mass), with the Milky Way halo, the LMC, the SMC, starburst galaxies, and DLAs at $z$=2--3 all showing a mean log\\,$N$(\\os) within 0.4\\,dex of 14.5, even though they span $\\approx$2\\,dex in [O/H]. While this characteristic value has been noticed before in the local Universe, the new result is that it applies down to the DLA regime at $-$1.6$\\la$[O/H]$\\la$--0.6, though there is a suggestion from the lowest-metallicity DLAs that $\\langle N$(\\os)$\\rangle$ falls off at [O/H]$<$--1.6. \\item For {\\bf intergalactic} \\os, there is surprisingly little evolution in the mean \\os\\ component column density over cosmic time. Over the 8.2\\,Gyr interval between $z$=2.28 and $z$=0.21, the mean log\\,$N$ % of intervening \\os\\ components with log\\,$N$(\\os)$>$13.2 (corresponding to rest equivalent widths $>$20\\,m\\AA) increases by only 0.17\\,dex, a factor of 1.5. % The distributions of $N$(\\os) at log\\,$N\\!>\\!13.5$, where the samples are complete, are shown to be insensitive to redshift. Furthermore, at both low and high redshifts, there is no difference in the mean log\\,$N$(\\os) between the intervening and proximate samples (so long as truly intrinsic absorbers are excluded). The insensitivity of log\\,$N$(\\os) to redshift and quasar proximity indicates an insensitivity to the strength of the ionizing radiation field, which (using the extragalactic background) is a factor of $\\approx$20 higher at $z$=2.28 than at $z$=0.21. \\item Intergalactic \\os\\ components are, on average, almost twice as broad at low-$z$ than at high-$z$, with the mean $b$(\\os) rising from 14\\kms\\ at $\\langle z\\rangle$=2.28 to 26~\\kms\\ at $\\langle z\\rangle$=0.21. % \\end{enumerate} The observation of a ``characteristic'' \\os\\ column density in many diverse galactic halos covering a range of mass and metallicity, is suggestive of a common origin or regulation mechanism. One such potential origin is the cooling-flow model of H02. In this model both galactic and intergalactic \\os\\ absorbers trace regions of initially-hot shock-heated plasma that are now radiatively cooling through coronal temperatures. The key advantage of this model is that it naturally explains the insensitivity of $N$(\\os) to metallicity and redshift, though it is unclear how it could explain the evolution in IGM $b$-values (result 3). The general framework of the cooling-flow theory allows the regions to take several forms, including conductive, turbulent, or shocked interfaces between warm ($\\sim$10$^4$\\,K) clouds and hot ($\\sim$10$^6$\\,K) plasma. However, many such regions would have to be intersected by a typical galaxy-halo sightline to build up the characteristic galactic $N$(\\os), which is $\\approx$30 times larger the column predicted in a single conductive interface. The idea that much of the galactic \\os\\ arises in coronal-temperature boundary layers is well-known and well-supported (see references in \\S3.1). The idea that \\emph{intergalactic} \\os\\ absorbers are also produced in such boundary layers, instead of by photoionization, is more controversial. The newly-demonstrated insensitivity of the mean intergalactic $N$(\\os) to $z$ can only be explained by photoionization if the gas density tracks the ionizing photon density, which would require the high-$z$ \\os\\ absorbers to have smaller sizes, lower metallicities, and/or lower ionization fractions at high-$z$ than the low-$z$ absorbers. While this is plausible (even expected), the kinematics of most intergalactic \\os\\ absorbers observed at low and high redshift (specifically the $b$-value differences and velocity-centroid offsets observed between \\os, \\cf, and \\hi) rule out single-phase photoionization, and require multi-phase models such as the cooling-flow scenario. Intergalactic \\os\\ absorbers are often discussed in the context of the elusive warm-hot intergalactic medium (WHIM), predicted by cosmological simulations to contain a substantial fraction of the present-day baryons \\citep{CO99, CO06, Ce01, Da01, FB01, Ka05, CF06}. The results presented here support the view that although \\os\\ absorbers do not directly trace the (hotter) bulk of the WHIM \\citep{OD09, TG10, Sm10}, they do trace the boundary layers where the WHIM interfaces with cooler, metal-enriched regions.\\\\ {\\it Acknowledgments.} The author acknowledges valuable conversations with Blair Savage, Benjamin Oppenheimer, Linda Smith, and Bart Wakker, and is grateful to the anonymous referee for an insightful report. He thanks Jacqueline Bergeron for providing her \\os\\ sample in electronic form. He appreciates the support of an ESO Fellowship, and a Marie Curie Intra-European Fellowship (contract MEIF-CT-2005-023720) awarded to investigate \\os. % This paper has made considerable use of the NASA Astrophysics Data System Abstract Service." }, "1101/1101.1233_arXiv.txt": { "abstract": "Recently, surface magnetic field maps had been acquired for a small sample of active M dwarfs, showing that fully convective stars (spectral types $\\sim$ M4 and later) host intense ($\\sim$ kG), mainly axi-symmetrical poloidal fields. In particular, the rapidly rotating M dwarf V374~Peg (M4), believed to lie near the theoretical full convection threshold, presents a stable magnetic topology on a time-scale of $\\sim 1$~yr. The rapid rotation of V374~Peg ($P=0.44$~days) along with its intense magnetic field point toward a magneto-centrifugally acceleration of a coronal wind. In this work, we aim at investigating the structure of the coronal magnetic field in the M dwarf V374~Peg by means of three-dimensional magnetohydrodynamical (MHD) numerical simulations of the coronal wind. For the first time, an observationally derived surface magnetic field map is implemented in MHD models of stellar winds for a low-mass star. We self-consistently take into consideration the interaction of the outflowing wind with the magnetic field and vice versa. Hence, from the interplay between magnetic forces and wind forces, we are able to determine the configuration of the magnetic field and the structure of the coronal winds. Our results enable us to evaluate the angular momentum loss of the rapidly rotating M dwarf V374~Peg. ", "introduction": "The rotational evolution of M dwarf (dM) stars can be inferred from observations of open clusters at different ages \\citep{2006MNRAS.370..954I, 2007MNRAS.381.1638S, 2009ApJ...691..342H, 2009MNRAS.400..451, 2009ApJ...695..679M}. In young ($\\lesssim 700~$Myr) open clusters, dM stars still present high rotation rates, which suggests that angular momentum losses at the early main-sequence phase are negligible for them \\citep{2009IAUS..258..363I}. However, as the cluster ages ($\\gtrsim 700~$Myr), the number of rapidly rotating dM stars decreases, implying that there should exist a mechanism of angular momentum removal that acts on time-scales of a few hundred million years \\citep{2007MNRAS.381.1638S}. For solar-like main sequence stars, the magnetised stellar wind is believed to spin down the star by carrying away stellar angular momentum. It has been observationally established that the angular velocity rate $\\Omega_0$ for solar-like stars varies as a function of age $t$ as $\\Omega_0 \\propto t^{-1/2}$ \\citep{1972ApJ...171..565S}. However, it seems that the empirical \\citeauthor{1972ApJ...171..565S}'s law is not valid for low-mass stars, suggesting that a solar-type wind (i.e., with low velocities and mass-loss rates) cannot reproduce the rotational evolution of fully-convective stars. The existence of hot coronae, rapid rotation, and high levels of magnetic activity in dM stars suggests the presence of winds with an enhanced mass loss as compared to the solar wind. However, the low-density, optically thin winds of these stars prevents the observation of traditional mass-loss signatures, such as P~Cygni profiles. The still unobserved high mass-loss rates from dM stars could be able to disperse debris discs, explaining why discs around dM stars older than $\\gtrsim 10$~Myr are scarcely found \\citep{2005ApJ...631.1161P}. Estimates of mass-loss rates from dM stars vary considerably. It has been suggested that the coronal winds of dM stars, despite of being very tenuous, possess mass-loss rates ($\\dot{M}$) that can considerably exceed the solar value ($\\dot{M}_\\odot \\simeq 2 \\times 10^{-14}~{\\rm M}_\\odot ~{\\rm yr}^{-1}$) by factors of $10$ to $10^4$ \\citep{1992ApJ...397..225M, 1992SvA....36...70B, 1996ApJ...462L..91L, 1997A&A...319..578V, 2001ApJ...546L..57W}, although \\citet{2001ApJ...547L..49W} claim an upper limit of $\\dot{M} \\lesssim 4 \\times 10^{-15}~{\\rm M}_\\odot ~{\\rm yr}^{-1}$ for Proxima Centauri (dM5.5e), $5$ times below the value of the solar wind mass-loss rate. In this work, we investigate the coronal wind of a specific fully-convective dM star, V374~Peg, for which observed surface magnetic maps have been acquired \\citep{2006Sci...311..633D, 2008MNRAS.384...77M}. For this, we use three-dimensional (3D) magnetohydrodynamics (MHD) simulations based on our previous models developed for solar-like stars \\citep{2009ApJ...699..441V} and weak-lined T Tauri stars \\citep{2009ApJ...703.1734V,vidotto10}. For the first time, an observationally derived surface magnetic field map is implemented in MHD models of stellar winds for a low-mass star. V374~Peg is a suitable case for modelling as a first step, because its surface magnetic distribution is close to potential, which implies that the adopted boundary conditions match the observed map closely. We self-consistently take into consideration the interaction of the outflowing wind with the magnetic field and vice-versa. Hence, from the interplay between magnetic forces and wind forces, we are able to determine the configuration of the magnetic field and the structure of its coronal wind. More details of this work can be found in \\citet{vidottomn}. ", "conclusions": "Observations of the rotation evolution of dM stars in open clusters at different ages provide a way to constrain the time-scale $\\tau$ for the angular-momentum loss. It has been suggested that $\\tau \\sim 200$~Myr or, mostly likely $400$ -- $800$~Myr, \\citep{2007MNRAS.381.1638S} for dM stars. Angular momentum of the star is carried away by the stellar wind. Because in our simulations there is no axi-symmetry, the torque $\\dot{\\bf J}$ on the star has $x$, $y$ and $z$ components. Here, we are interested only on the $z$-component, as it is the only one responsible for the rotational braking (because the angular velocity of the star points in the $z$-direction). The $z$-component of the angular momentum carried by the wind is \\citep{1970MNRAS.149..197M} \\begin{eqnarray}\\label{eq.jdot} \\dot{J} &=& \\left[\\alpha{\\bf \\hat{z}} \\times \\int_{V_A} {\\bf r} \\times \\rho ({\\bf V }+ \\alpha{\\bf \\hat{z}} \\times {\\bf r}){\\rm d}V_A \\right]_z + \\int_{S_A} \\left( p + \\frac{B^2}{8\\pi} \\right) ({\\bf r} \\times {\\bf \\hat{n}})_z {\\rm d}S_A \\nonumber \\\\ &+& \\int_{S_A} \\left[ {\\bf r} \\times (\\alpha{\\bf \\hat{z}} \\times {\\bf r})\\right]_z \\rho {\\bf V} \\cdot {\\bf \\hat{n}} {\\rm d} { S_A} , \\end{eqnarray} where ${\\bf V} = {\\bf u} - \\alpha {\\bf \\hat{z}}\\times {\\bf r}$ is the velocity vector in the frame rotating with angular velocity $\\alpha {\\bf \\hat{z}}$, ${\\bf \\hat{z}}$ is the unit vector that points in the $z$-direction, $S_A$ is the Alfv\\'en surface that delimits the volume $V_A$, and ${\\bf \\hat{n}}$ is the normal unit vector to the Alfv\\'en surface. The first term on the right of Eq.~(\\ref{eq.jdot}) does not contribute to the $z$-component torque and is therefore null. The second term disappears in the case of a spherical Alfv\\'en surface, but it is non-null in the cases where a surface magnetic map is considered and it becomes relatively more important for the cases with larger adopted $\\beta_0$. The third term is the dominant term in our simulations. We can estimate the time-scale for rotational braking as $\\tau = {J}/{\\dot{J}}$, where $J$ is the angular momentum of the star. If we assume a spherical star with a uniform density, then $J = 2/5 M_* R_*^2 \\Omega_0$ and the time-scale is \\begin{equation} \\tau \\simeq \\frac{9 \\times 10^{36}}{\\dot{J}} \\left( \\frac{M_*}{M_\\odot}\\right) \\left( \\frac{1~{\\rm d}}{P_0} \\right) \\left( \\frac{R_*}{R_\\odot}\\right)^2 ~{\\rm Myr}, \\end{equation} where $P_0=2\\pi/\\Omega_0$ is the rotational period of the star. For V374~Peg, this results in \\begin{equation}\\label{eq.tau} \\tau \\simeq \\frac{6.45 \\times 10^{35}}{\\dot{J}} ~{\\rm Myr}. \\end{equation} Because $\\dot{J}$ depends on the mass flux crossing a given surface, i.e., on the mass-loss rate of the wind $\\dot{M}$, from Eq.~(\\ref{eq.windmassloss}), we have a rough scaling relation between $\\dot{J}$ and $\\dot{M}$ for cases 1Map, 2Map, and 3Map \\begin{equation}\\label{windangloss} \\dot{J} \\propto \\dot{M} \\propto n_0^{1/2} , \\end{equation} which implies in a time-scale [Eq.~(\\ref{eq.tau})] for rotational braking that scales as \\begin{equation}\\label{eq.windtimescale} \\tau \\propto n_0^{-1/2} . \\end{equation} For cases 1Map, 2Map, and 3Map, $\\tau \\simeq 18 {n_{12}^{-1/2}}$~Myr, well below the estimated solar spin-down time $\\tau_\\odot \\simeq 7$~Gyr \\citep{1967ApJ...148..217W}. Table~\\ref{table} presents the mass and angular momentum loss rates, and the time-scale for rotational braking calculated for all simulations, where we verify the approximate scaling given by Eqs.~(\\ref{eq.windmassloss}), (\\ref{windangloss}), and (\\ref{eq.windtimescale}). Comparing to the observationally derived rotational braking time-scales of a couple of hundreds of Myr for dM stars is open clusters \\citep{2007MNRAS.381.1638S}, we tend to rule out cases with larger coronal base densities (i.e., $n_0 \\gtrsim 10^{11}~{\\rm cm}^{-3}$). According to this comparison, the most plausible wind density is the one used for models 1Map. Such a density is also able to reproduce typical emission measures of dM stars (${\\rm EM} \\approx 10^{51}~{\\rm cm}^{-3}$) and comparatively (with the remaining cases) smaller mass-loss rates and higher wind velocities. Ultimately, when the star ages, the stellar rotation brakes, reducing the stellar surface magnetic field intensity, and therefore the wind velocity. With the inclusion of an observed distribution of surface magnetic field, we head towards a more realistic modelling of magnetised coronal winds. Never the less, our model presents limitations, such as the neglect of a detailed energy balance. Instead, we consider a polytropic relation between pressure and density parametrised through $\\gamma$ in the derivation of the energy equation of the wind. Once the magnetic field distribution is set, the thermal pressure adjusts itself in order to provide a distribution of heating/cooling that is able to support the MHD solution obtained \\citep{1986ApJ...302..163L}. If the wind of V374~Peg is able to cool down, e.g., by radiative cooling not considered in our models, the terminal velocities of the wind could be considerably smaller. Depending on where in the wind energy deposition (or removal) occurs, the wind velocity may change, without affecting the mass-loss rates. For instance, if a substantial cooling occurs above the Alfv\\'en surface, the velocity profile of the wind from that point outwards will be affected. As the information of what is happening above the Alfv\\'en point cannot be transmitted to the sub-Alfv\\'enic region, the wind density and velocity profiles in the proximity of the star will not be changed, and consequently neither the stellar mass-loss/angular momentum-loss rates." }, "1101/1101.1254.txt": { "abstract": "% In this paper we consider how non-Gaussianity of the primordial density perturbation and the amplitude of gravitational waves from inflation can be used to determine parameters of the curvaton scenario for the origin of structure. We show that in the simplest quadratic model, where the curvaton evolves as a free scalar field, measurement of the bispectrum relative to the power spectrum, $\\fNL$, and the tensor-to-scalar ratio can determine both the expectation value of the curvaton field during inflation and its dimensionless decay rate relative to the curvaton mass. % If we further assume the initial value of the curvaton field is given by a stochastic distribution then the curvaton mass and decay rate may be % determined separately. % We show how these predictions are altered by the introduction of self-interactions, in models where higher-order corrections are determined by a characteristic mass scale and discuss how additional information about primordial non-Gaussianity and scale dependence may constrain curvaton interactions. % ", "introduction": "%[THE INTRODUCTION SHOULD GO THIS WAY: INFLATION AS A SOLUTION FOR DIFFERENT PROBLEMS; THEN DIFFERENT SCALAR FIELDS ON THE EARLY UNIVERSE CAN BE THE ORIGIN OF STRUCTURE, I.E., INFLATION VS OTHER FIELDS; CURVATON PROPOSAL; GW-INFLATION, FNL-CURVATON; GW+FNL GIVE PARAMETER CONSTRAINTS; MINIMAL MODEL = QUADRATIC POTENTIAL; SELF INTERACTIONS (AXION + COSH + POLYNOMIAL)] Inflation solves the horizon problem, the flatness problem and the monopole problem. Furthermore, it gives a simple way to source primordial perturbations from quantum vacuum fluctuations. Any light scalar field during a period of inflation with an almost constant Hubble expansion acquires an almost scale-invariant power spectrum of fluctuations that could be the origin of primordial density perturbations \\cite{Bassett:2005xm,Lyth:2009zz}. The curvaton is one such field which is only weakly coupled and hence decays on a time-scale much longer than the duration of inflation \\cite{Linde:1996gt,Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct,Wands:2005aa}. Its lightness enables the field to acquire super-Hubble perturbations from vacuum fluctuations during inflation. When it decays into radiation some time after inflation has ended, its decay can source the perturbations in the radiation density of the universe, and all other species in thermal equilibrium \\cite{Lyth:2002my,Weinberg:2003sw}. One of the distinctive predictions of the curvaton scenario for the origin of structure is the possibility of non-Gaussianity in the distribution of the primordial density perturbations \\cite{Bartolo:2004ty,Boubekeur:2005fj,Lyth:2005fi}. Treating the curvaton as a pressureless fluid one can estimate the resulting non-Gaussianity either analytically by treating the decay of the curvaton as instantaneous \\cite{Lyth:2002my,Bartolo:2004ty,Lyth:2005fi}, or numerically \\cite{Malik:2006pm,Sasaki:2006kq}, showing that the non-Gaussianity parameter $\\fNL$ becomes large when the curvaton density at the decay time becomes small. The non-linear evolution of the field before it decays can also contribute to the non-Gaussianity of the final density perturbation. % The authors of \\cite{Enqvist:2005pg,Enqvist:2008gk,Huang:2008zj,Enqvist:2009zf,Enqvist:2009eq,Enqvist:2009ww} look at the effect of polynomial corrections to the quadratic curvaton potential. In some cases the curvaton density can be significantly subdominant at decay and still yield small $\\fNL$ \\cite{Enqvist:2009zf}. For small values of $\\fNL$, the non-Gaussianity can instead be probed by the trispectrum parameter, $\\gNL$. % [WHAT ABOUT OTHER INVESTIGATIONS OF NON-GAUSSIANITY IN, FOR EXAMPLE, COSINE POTENTIALS? SHOULD REFER TO ALL RELEVANT WORK.] Primordial gravitational waves on super-Hubble scales are also present since they are an inevitable byproduct at some level of an inflationary expansion. Non-Gaussianity alone could distinguish between the curvaton scenario and the conventional inflaton scenarios for the origin of structure since a single inflaton field is not capable of sourcing significant non-Gaussianity \\cite{Maldacena:2002vr}. But non-Gaussianity and gravitational waves together can give tight constraints on curvaton model parameters. Nakayama {\\em et al} \\cite{Nakayama:2009ce} recently studied the effects of the entropy released by the decay of a curvaton field with a quadratic potential on the spectrum of gravitational waves that are already sub-horizon scale at the decay and consider the possibilities of future direct detection experiments, such as BBO or DECIGO, to constrain the parameter space. We will restrict our attention to gravitational waves on super-Hubble scales when the curvaton decays which are not affected by the decay, and consider self-interactions of the curvaton field in addition to the quadratic potential \\cite{Enqvist:2010dt,Huang:2008zj}. This includes scales which contribute to the observed CMB anisotropies, where the power in gravitational waves is typically given by the tensor-to-scalar ratio for the primordial metric perturbations, $\\rT$. In this paper we will investigate how non-Gaussianity and gravitational waves provide constraints on curvaton model parameters. For any value of the curvaton model parameters we can obtain the observed amplitude of primordial density perturbations on large scales by adjusting the Hubble scale of inflation, which we assume to be an independent parameter in the curvaton model. However observational constraints on the tensor-to-scalar ratio places an upper bound on the inflationary Hubble scale, while non-Gaussianity constrains the remaining model parameters. % [NOT SURE THAT THE FOLLOWING PARAGRAPH CONTAINS ANYTHING ESSENTIAL AT THIS STAGE.] % Previously there have been studies of the curvaton decays into radiation using a fluid description of the system \\cite{Gupta:2003jc}. We are only %interested in the evolution of the curvaton prior to decay to determine the non gaussian parameter $f_{nl}$ in term of curvaton model parameter. To do %so we will take a field description of the system. Similarly we need to relate the tensor-to-scalar ratio with the same parameters. We do so by saying %that the power spectrum of gravitational waves only depends on the Hubble parameter during inflation, $H_*$. This can be determine from the power %spectrum predicted for the curvaton. Doing so we are able to relate the set of model parameters $(\\chi_*,m/\\Gamma)$ with the set of observables %$(f_{nl},r_{gw})$. Therefore, non-gaussianity and gravitational waves can be used in a complementary way to find which parts of parameter space are %still allowed. We numerically solve the evolution of the curvaton field in a homogeneous radiation-dominated era after inflation allowing for non-linear evolution of the curvaton field due to both explicit self-interaction terms in the potential and the self-gravity of the curvaton. In particular we consider quadratic and non-quadratic potentials which reduce to a quadratic potential about the minimum with self-interaction terms governed by a characteristic mass scale, corresponding to cosine or hyperbolic-cosine potentials. Cosine potentials arise for PNGB axion fields and are often considered as candidate curvaton fields \\cite{Chun:2004gx,Dimopoulos:2005bx,Kawasaki:2008mc,Chingangbam:2009xi,Huang:2010cy}. The hyperbolic cosine is representative of a potential where self-interaction terms become large beyond a characteristic scale. In each case we show how the non-linearity parameter $\\fNL$ and tensor-to-scalar ratio, $\\rT$, can be used to determine model parameters. In Section~II we review the perturbations generated during inflation and how these are transfered to the primordial density perturbation in the curvaton scenario. In Section~III we present the numerical results of our study for three different curvaton potentials. We conclude in Section~IV. ", "conclusions": "%[THE CONCLUSION MORE THAN STATING RESULTS SHOULD FOCUS ON THE FREE PARAMETERS OF THE MODEL AND THE OBSERVABLES. INFLATION HAS 3 FREE PARAMETERS: $H_*,\\epsilon_*,\\eta_*$. THE CURVATON HAS 6 FREE PARAMETERS:$H_*,\\chi_*,m,\\Gamma,\\eta_*,\\eta_{\\phi,*}$] In this work we have investigated the numerical evolution of a curvaton field from its overdamped regime after inflation until it decays into radiation. We have shown how measurement of both the non-linearity parameter, $\\fNL$, and the tensor-to-scalar ratio, $\\rT$, provide complementary constraints on the model parameters. We did this for three different curvaton potentials: the quadratic potential, axion-like cosine potentials and hyperbolic potentials. As expected both the cosine and the hyperbolic potentials recover the quadratic regime when $\\chi_*\\ll f$. For the simplest quadratic potential for the curvaton, bounds on the tensor-to-scalar ratio place an upper bound on the dimensionless decay rate, ruling out large regions of parameter space that would yield a large primordial non-Gaussianity in the distribution of scalar perturbations. % Simultaneous measurement of both the non-linearity parameter, $\\fNL$, and the tensor-to-scalar ratio, $\\rT$, can determine both the expectation value of the field during inflation, $\\chi_*$, and the dimensionless decay rate, $\\Gamma/m$. In the conventional inflaton scenario for the origin of structure we have three free parameters: the inflation scale $H_*$ and two slow-roll parameters, $\\epsilon$ and $\\eta_{\\phi}$. These can be determined by power of the primordial scalar perturbations, $\\P_{\\zeta}$, the tensor perturbations, $\\P_{T}$, and the spectral index of the scalar spectrum, $n_{\\zeta}$. The spectral index of the tensor spectrum, if measurable, would give a valuable consistency check \\cite{Baumann:2008aq}. Another important consistency condition for canonical, slow-roll inflation is that the primordial density perturbations should be Gaussian and the non-linearity parameter, $\\fNL$, should be much less that unity \\cite{Maldacena:2002vr}. In the curvaton scenario with a simple quadratic potential we have 5 free parameters: the inflation scale $H_*$, the expectation value of the curvaton during inflation $\\chi_*$, the decay rate of the curvaton relative to its mass, $\\Gamma/m$, and the slow roll parameters $\\epsilon$ and $\\eta_{\\chi}=m_*^2/3H_*^2$. % For a curvaton, we find that $H_*$, $\\chi_*$ and $\\Gamma/m$ are determined by the primordial scalar perturbations, $\\P_{\\zeta}$, the tensor perturbations, $\\P_{T}$, and the non-linearity parameter, $\\fNL$, but the mass and decay rate of the curvaton are not separately determined. The two slow-roll parameters $\\epsilon$ and $\\eta_{\\chi}$ are then determined by the two spectral indices $n_{\\zeta}$, and $n_{T}$. Another natural observable in the curvaton model is the scale dependence of the non-linearity parameter, defined as \\cite{Byrnes:2009pe} % \\be n_{\\fNL} \\equiv \\frac{d \\ln|\\fNL|}{d\\ln k} \\ee % In the curvaton scenario this is given by a simple expression \\cite{Byrnes:2010xd,Huang:2010cy} \\be n_{\\fNL} = \\eta_3 \\frac{g}{\\mPl g'} \\frac{5}{4\\R\\fNL} \\,. \\ee where we define $\\eta_3 \\equiv \\mPl^3 V'''/V$. This can be rewritten in terms of observable quantities and $\\eta_3$ \\be n_{\\fNL} = \\eta_3 \\frac{5}{12\\sqrt{2}} \\frac{\\sqrt{r_T}}{\\fNL} \\,. \\ee Thus it offers the possibility of testing the curvaton self interactions. Future observations may be able to detect $|\\fNL n_{\\fNL}|> 5$ \\cite{Sefusatti:2009xu}, corresponding to $|\\eta_3|\\sqrt{\\rT}>17$. For the quadratic potential we have the consistency condition $n_{\\fNL}=0$. Deviations from a quadratic potential introduce at least one further model parameter, $f$, corresponding to the mass scale associated with the non-linear corrections. This leads to a degeneracy in model parameters consistent with the five observables $\\P_{\\zeta}$, $\\P_{T}$, $\\fNL$, $n_{\\zeta}$ and $n_{T}$, but this can be broken by a measurement of $n_{\\fNL}$. In the case of a cosine-type curvaton potential the self interaction corrections became important near the top of the potential, i.e., when $\\chi_* \\sim \\pi f$ \\cite{Huang:2010cy} and the tensor-to-scalar ratio no longer places an upper bound on $\\Gamma/m$. As for a quadratic curvaton, we still find $\\fNL>-5/4$ and hence any large non-Gaussianity, $|\\fNL|\\gg1$, has positive $\\fNL$. But for $\\chi_*\\sim f$ we have $\\eta_3 \\sim -(\\mPl/f^3)<0$, and if $f$ is well below the Planck scale there could be strong scale dependence. %This can be avoided at a low inflation scale. But, as seen above, this requires an high $m/\\Gamma$ ratio to agree with the bounds on $\\fNL$. On the %other hand we can cross the upper bounds on $m/\\Gamma$ as stated before. In the case of a hyperbolic-type potential $\\fNL$ can become large and negative, for $\\chi_*\\sim f$. However the tensor-to-scalar ratio again plays an important role, in this case placing a lower bound on $\\fNL$, e.g., $\\fNL>-100$ for $\\rT<0.1$ when $f=10^{16}$~GeV. In this regime we find $\\eta_3 \\sim (\\mPl/f^3)>0$, which can be large, leading to strong scale dependence for $f\\ll \\mPl$, with $n_{\\fNL}<0$ for $\\fNL<0$. Running of either the scalar tilt, $\\alpha_\\zeta$, or the non-linearity, $\\alpha_{\\fNL}$ \\cite{Huang:2010cy}, yields additional information about the higher derivatives of the potential, and in particular curvaton-inflaton interactions which we have assumed are negligible in our analysis. Significant non-Gaussianity in the primordial perturbations opens up the possibility to extract information from the higher-order correlations in the scalar spectrum, such as the trispectrum \\cite{Byrnes:2006vq,Sasaki:2006kq,Enqvist:2008gk,Huang:2008bg,Enqvist:2009ww} % \\bea T_{\\zeta}(k_1,k_2,k_3,k_4) = \\frac{54}{25} \\gNL \\left[ P_{\\zeta}(k_2)P_{\\zeta}(k_3)P_{\\zeta}(k_4) +3 ~ {\\rm perms} \\right] + \\frac{36}{25} \\fNL^2 \\left[P_{\\zeta}(k_{13})P_{\\zeta}(k_3) P_{\\zeta}(k_4) + 11~{\\rm perms} \\right] \\,. \\eea % which are sensitive to higher-order derivatives of the expansion history with respect to the curvaton field value during inflation through $\\gNL= (25/54)N'''/N'^3$. % Differentiating Eq.~(\\ref{fNLRused}) we obtain % \\be \\label{gNLRused} \\gNL= \\frac {25}{24} \\left[ \\frac{\\R''}{\\R^3}\\frac{g^2}{g'^2}+2\\frac{\\R'}{\\R^3}\\left(\\frac{g^2g''}{g'^3}-\\frac{g}{g'}\\right)+\\frac 1 {\\R^2} \\left(\\frac{g^2g'''}{g'^3}-3\\frac{gg''}{g'^2}+2\\right)\\right] \\,. \\ee which using the sudden-decay approximation can be written as \\cite{Sasaki:2006kq,Byrnes:2006vq} % \\be \\gNL = \\frac{25}{54} \\left[ \\frac{9}{4\\R^2} \\left( \\frac{g^2g'''}{g^{\\prime3}} + 3\\frac{gg''}{g^{\\prime2}} \\right) - \\frac9\\R \\left( 1 + \\frac{gg''}{g^{\\prime2}} \\right) + \\frac12 \\left( 1 -9 \\frac{gg''}{g^{\\prime2}} \\right) +10\\R +3\\R^2 \\right] \\,. \\ee % $\\gNL$ and its scale dependence $n_{\\gNL}$ \\cite{Byrnes:2010ft,Byrnes:2010xd}, thus provide additional observable parameters which then offer consistency conditions for generalised curvaton models such as the cosine or hyperbolic potentials. % In practice we require more accurate numerical simulations than those used in this work to reliably determine the required higher-derivatives with respect to the initial field value across the range of model parameters used in this paper and we leave this for future work." }, "1101/1101.5161_arXiv.txt": { "abstract": "{ % Some disturbed galaxy clusters host diffuse elongated radio sources, also called radio relics. It is proposed that these relics trace shock waves in the intracluster medium (ICM). Within the shock waves, generated by cluster merger events, particles are accelerated to relativistic energies, and in the presence of a magnetic field synchrotron radiation will be emitted. CIZA~J2242.8+5301 is a disturbed galaxy cluster hosting complex diffuse radio emission, including a so-called double radio relic. Here we present new Giant Metrewave Radio Telescope (GMRT) radio observations of CIZA~J2242.8+5301 at 325 and 150~MHz. We detect the double radio relic at 150 and 325 MHz. The very deep 150~MHz image reveals the presence of large-scale diffuse emission between the two radio relics. ", "introduction": "Galaxy clusters grow by mergers with other clusters or galaxy groups, as well as through the accretion of gas from the intergalactic medium (IGM). Both these two processes shock the ICM. It has been proposed that within these shocks particles can be accelerated to highly relativistic energies by the diffusive shock acceleration (DSA) mechanism \\citep{1977DoSSR.234R1306K, 1977ICRC...11..132A, 1978MNRAS.182..147B, 1978MNRAS.182..443B, 1978ApJ...221L..29B, 1983RPPh...46..973D, 1987PhR...154....1B, 1991SSRv...58..259J, 2001RPPh...64..429M}. It is thought that radio relics, elongated steep-spectrum radio sources \\citep[e.g.,][]{1991A&A...252..528G, 2006Sci...314..791B,2008A&A...486..347G,2011ApJ...727L..25B}, trace these shock waves generated by cluster merger events \\citep{1998A&A...332..395E, 2001ApJ...562..233M}. \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics[angle =90, trim =0cm 0cm 0cm 0cm,width=1.0\\textwidth, clip=true]{ciza22_325}} \\caption{\\footnotesize GMRT 325~MHz image made with robust weighting set to $-1.0$ \\citep{briggs_phd}. The image has a rms noise of 59~$\\mu$Jy~beam$^{-1}$. Contour levels are drawn at $\\sqrt{[1, 2, 4, 8, \\ldots]} \\times 4\\sigma_{\\mathrm{rms}}$. The beam size is $8.7\\arcsec \\times 7.4\\arcsec$ and shown in the bottom left corner of the image. A thinner contour from an image with robust weighting set to $0.5$, with a resolution of $11.5\\arcsec \\times 9.8\\arcsec$, is drawn at a level of $0.2$~mJy~beam$^{-1}$. } \\label{fig:325} \\end{figure*} CIZA~J2242.8+5301 is a disturbed galaxy cluster located at $z=0.1921$ \\citep{2007ApJ...662..224K}. The cluster hosts a large double radio relic system, as well as additional large-scale diffuse radio emission \\citep{2010Sci...330..347V}. The spectral index ($\\alpha$) across the bright northern relic steepens systematically in the direction of the cluster center, across the full length of the narrow relic. This is expected for outwards moving shock waves, with synchrotron and inverse Compton losses behind the shock front. Here we present new low-frequency GMRT radio images at 150 and 325~MHz, which complement our previous higher frequency observations at 610 and 1400~MHz \\citep{2010Sci...330..347V}. ", "conclusions": "We detect both the northern and southern radio relics in the galaxy cluster \\object{CIZA~J2242.8+5301} with the GMRT at 150 and 325~MHz . The 150~MHz image reveals the presence of additional diffuse emission throughout the cluster with a total extent of about 3~Mpc, confirming the results from our a previous WSRT 1.4~GHz observation. In a future paper we will present a more detailed spectral analysis of the radio emission in the cluster." }, "1101/1101.0744_arXiv.txt": { "abstract": "% It is known that scalar-tensor theory of gravity admits regular crossing of the phantom divide line $w_{\\DE}=-1$ for dark energy, and existing viable models of present dark energy for its particular case -- $f(R)$ gravity -- possess one such crossing in the recent past, after the end of the matter dominated stage. It was recently noted that during the future evolution of these models the dark energy equation of state $w_{\\DE}$ may oscillate with an arbitrary number of phantom divide crossings. In this paper we prove that the number of crossings can be infinite, present an analytical condition for the existence of this effect and investigate it numerically. With the increase of the present mass of the scalaron (a scalar particle appearing in $f(R)$ gravity) beyond the boundary of the appearance of such oscillations, their amplitude is shown to decrease very fast. As a result, the effect quickly becomes small and its beginning is shifted to the remote future. ", "introduction": "% The accelerating expansion of the present Universe is confirmed by current precise observational data such as type Ia supernovae~\\cite{Perlmutter:1998np,Riess:1998cb}, anisotropy of cosmic microwave background~\\cite{Komatsu:2010fb}, large scale structure~\\cite{Tegmark:2003ud} and baryon acoustic oscillations~ \\cite{Eisenstein:2005su,Percival:2007yw}. The standard cosmological constant ($\\rm \\Lambda$)-Cold-Dark-Matter (${\\rm CDM}$) model is indeed able to explain these observational results within observational errors. In this model a cosmological constant is regarded as a new fundamental physical constant. However, the required value of the cosmological constant is very tiny compared with any known physical scales. Thus, its relation to the standard quantum theory of known particles and fields is not understood today although some nonperturbative effects may naturally generate such a small quantity, see e.g. Refs.~\\cite{Yokoyama:2001ez,Kiefer:2010pb}. More generally, a source of the current cosmic acceleration is called dark energy (DE). Further we shall use the more detailed term \"present DE\" for it to distinguish it from primordial DE which was responsible for another accelerated expansion regime, dubbed inflation, which occurred in the very early Universe~ \\cite{Starobinsky:1980te,Sato:1980yn,Guth:1980zm}. The relation of primordial DE to the known elementary particles has not been established, too. Both primordial and present DE can have either a physical origin (some new physical fields of matter) or a geometrical one. In the latter case, the Einstein gravity becomes modified. One of the simplest and self-consistent generalizations of the Einstein gravity is $f(R)$ gravity which incorporates a new phenomenological function $f(R)$ of the Ricci scalar $R$ (with $d^2f/dR^2$ not identically zero) into the action, see Eq.~\\eqref{action} below. For a long time this theory of gravity was known to contain viable inflationary models, among them the simplest one introduced in Ref.~\\cite{Starobinsky:1980te} that remains in agreement with the most recent observational data. Thus, $f(R)$ gravity can successively describe primordial DE. Rather recently, after many unsuccessful trials, viable models of present dark energy were found \\cite{Hu:2007nk,Appleby:2007vb,Starobinsky:2007hu} which provide non-trivial alternatives to the standard $\\lcdm$ model.\\footnote{In order not to destroy the standard early Universe cosmology, including the recombination, the correct Big Bang nucleosynthesis and inflation of any kind, these models of present DE possessing a non-trivial form of $f(R)$ in the low-$R$, $R>0$ region have to be further generalized by changing the behaviour of $f(R)$ at large $R$ and by extending it to the region of negative $R$, see Ref.~\\cite{Appleby:2009uf}. However, this generalization is not important for our study.} This theory is a special class of the scalar-tensor theory of gravity with the vanishing Brans-Dicke parameter $\\omega_{BD}$. It contains a new scalar degree of freedom dubbed \"scalaron\" in Ref.~\\cite{Starobinsky:1980te}, thus, it is a {\\em nonperturbative} generalization of the Einstein gravity. From the quantum point of view, scalaron is a massive spin-$0$ particle which mass depends on $R$. We consider $f(R)$ gravity as a phenomenological macroscopic theory of gravity, alternative to the Einstein one, without discussing its microscopic origin.\\footnote{Note the simplest possible mechanism that has attracted a new interest recently: a scalar field $\\phi$ with some potential and the non-minimal coupling $-\\xi R\\phi^2/2$ to the Einstein gravity in the limit of a very large negative $\\xi$ (i.e. the sign of coupling is opposite to that of the conformally coupled case). However, this mechanism leads to $df/dR > 1$. So, while sufficient to produce the functional form of $f(R)$ needed for successful inflationary models, it is not useful for construction of viable models of present DE.} The existence of this additional degree of freedom imposes a number of constraints on the functional form of $f(R)$ in viable cosmological models. In particular, in order to have the correct Newtonian limit, as well as the standard matter-dominated stage with the scale factor behaviour $a(t)\\propto t^{2/3}$ driven by cold dark matter and baryons, the following conditions should be fulfilled for $R\\gg R_0$ where $R_0\\equiv R(t_0)\\sim H_0^2$, $t_0$ is the present moment and $H_0$ is the Hubble constant, and up to curvatures in the centre of neutron stars: \\be |f(R)-R|\\ll R,~~|f'(R)-1|\\ll 1,~~Rf''(R)\\ll 1. \\ee Here the prime denotes the derivative with respect to the argument $R$. Furthermore, $f(R)$ should satisfy the following conditions to guarantee both that Newtonian gravity solutions are stable and that the standard matter-dominated Friedmann-Robertson-Walker (FRW) stage remains an attractor with respect to an open set of neighbouring generic cosmological solutions in $f(R)$ gravity: \\be f'(R)>0,~ f''(R)>0. \\ee In quantum language, the first condition means that gravity is attractive and graviton is not a ghost, while the second one -- that scalaron is not a tachyon. Specific functional forms of $f(R)$ satisfying these conditions, as well as laboratory and Solar system tests of gravity, and possessing a future stable (or at least metastable) de Sitter stage that is required for correct description of observable properties of present DE, have been proposed in Refs.~\\cite{Hu:2007nk,Appleby:2007vb,Starobinsky:2007hu}, and much work has been carried out on their cosmological consequences. In order to describe the difference between FRW background solutions of $f(R)$ gravity and the $\\lcdm$ model, it is useful to introduce the effective equation-of-state~(EoS) parameter for DE $w_{\\DE}\\equiv P_{\\DE}/\\rho_{\\DE}$ where the effective pressure $P_{\\DE}$ and the effective energy density $\\rho_{\\DE}$ of DE are determined using the Einsteinian representation of gravitational field equations, see Eqs.~\\eqref{E1},~\\eqref{E2} below. Another independent parameter which describes scalar (density) perturbations on a FRW background is the gravitational growth index $\\gamma$ defined as $d\\ln \\d/d\\ln a\\equiv \\Omega_m(z)^{\\gamma(z)}$ where $\\d\\equiv \\d\\rho_m/\\rho_m$ and $\\Omega_m\\equiv 8\\pi G\\rho_m/3H^2$ are a matter density fluctuation and the density parameter for matter, respectively. In $f(R)$ gravity, $w_{\\DE}$ is time dependent and $\\gamma$ is time and scale dependent whilst they keep the constant value $w_{\\DE}=-1$ and $\\gamma\\approx 6/11$ in the $\\lcdm$ model. Time and scale dependency of $\\gamma$ generate an additional transfer function for matter density fluctuation that constrains the model parameter region~\\cite{Motohashi:2010tb,Motohashi:2009qn,Motohashi:2010sj}. One of the most interesting features of geometrical DE distinguishing it from physical DE based on non-ghost physical fields minimally coupled to gravity, like quintessence, is the possibility of phantom behaviour, $w_{\\DE}<-1$, of DE. Moreover, this behaviour may well be temporary with DE becoming normal, $w_{\\DE}> -1$, after smooth crossing of the phantom boundary $w_{\\DE}=-1$. In particular, models of geometrical DE based on scalar-tensor gravity were long known to admit this property~\\cite{BEPS00}. $f(R)$ gravity is a particular case of scalar-tensor gravity, so it permits phantom behaviour of DE and smooth crossing of the phantom boundary, too. Existing observational data do not exclude the possibility of phantom behaviour of DE (though they do not specifically favour it, too) for the following simple reason: as was noted above, DE in the particular form of an exact cosmological constant ${\\rm \\Lambda}$ is in a good agreement with all data. But since $w_{\\rm \\Lambda}\\equiv -1$, it lies exactly at the phantom boundary. Thus, any small deviation of DE from ${\\rm \\Lambda}$ to the direction of decreasing $w_{\\DE}$ results in its phantom behaviour. So, theory has to be prepared for this possibility that explains large interest in DE models admitting it. However, it is clear already that this \"phantomness\" should be small. In particular, if it is assumed for simplicity that $w_{\\DE}={\\rm const}$, when $|w_{\\DE}+1|<0.1$ at the approximately $2\\sigma$ confidence level~\\cite{Komatsu:2010fb}. Moreover, viable $f(R)$ models of present DE~\\cite{Hu:2007nk,Appleby:2007vb,Starobinsky:2007hu} generically exhibit phantom behaviour during the matter-dominated stage and one recent crossing of the phantom divide $w_{\\DE}=-1$ even in the case of the smoothest behaviour of a FRW scale factor $a(t)$, when there were no superimposed small oscillations of $a(t)$ in the past (in quantum language, no condensate of primordial scalarons with the zero momentum) ~\\cite{Hu:2007nk,MMA09,Appleby:2009uf,Motohashi:2010tb}. From the physical point of view, the absence of primordial scalarons in the viable $f(R)$ models of present DE is needed in order not to destroy the standard early Universe cosmology and it can be achieved by primordial inflation of any kind, see Ref.~\\cite{Appleby:2009uf} for a detailed discussion. Using numerical calculations, it has been recently shown that even in this smoothest case the EoS parameter $w_{\\DE}$ can oscillate around the future de Sitter solution in these DE models~\\cite{Bamba:2010iy}, see also Ref.~\\cite{LKM10}. To investigate the phenomenon of multiple crossing of the phantom divide in more detail and analytically, in the present paper we prove that this crossing can indeed occur infinitely many times during the future evolution of viable $f(R)$ models of present DE if the scalaron mass at a future stable de Sitter stage in these models is sufficiently large. Though this phenomenon is not directly observable since it refers to remote future, it is interesting and important from the theoretical point of view. Also, it is possible to check from observational data at the present moment if the derived analytical criterion for the existence of the infinite number of crossings is satisfied or not. Thus, the present paper focuses on the oscillatory behaviour of $w_{\\DE}$ around the phantom divide $w_{\\DE}=-1$ in the future. In Sec.~II, we review the stability conditions and the condition of the existence of a stable future de Sitter stage in $f(R)$ gravity, and derive the condition for the existence of the infinite number of these oscillations analytically using the perturbation theory around the de Sitter solution. In Sec.~III, we focus on the specific viable model of present DE in $f(R)$ gravity and present results from numerical calculations relating this condition to observable properties of the Universe at the present time. Sec.~IV is devoted to conclusions and discussion. ", "conclusions": "% We have investigated conditions under which the effective EoS parameter $w_\\DE$ of present DE in $f(R)$ gravity can oscillate an infinite number of times around the phantom boundary $w_{\\DE}=-1$ during the future evolution of the Universe. The analytical condition of the existence of this phenomenon, Eq.~\\eqref{onc}, is derived that depends on the properties of $f(R)$ near a future stable de Sitter stage only. The physical sense of this condition is that the rest mass of the scalaron (a massive scalar particle which arises in $f(R)$ gravity in addition to massless spin-2 graviton) should be sufficiently large at the future de Sitter stage. Thus, this phenomenon is generic. However, the amplitude of these oscillations has been shown to decrease fast with the increase of the scalaron mass beyond the boundary of the appearance of such oscillations. As a result, the effect quickly becomes small and its beginning is shifted to the remote future. For real scalaron masses lying below this boundary, the future stable de Sitter stage is reached without the phantom boundary crossing. Analytic solutions for the behaviour of $w_{\\DE}$ near the phantom boundary have been obtained in the first order of the small quantity $|w_{\\DE}+1|$. Generically they have a monotonically decaying part $\\d w_{\\rm dec}$ and a damped harmonic oscillatory part $\\d w_{\\rm osc}$. For a specific viable $f(R)$ model of present DE energy, numerical integration of FRW background evolution has been performed which future behaviour is in a good agreement with the analytic formulas. All calculations have been done for the smoothest initial conditions in the past corresponding to the absence of a primordial homogeneous oscillating scalaron component. So, even in this case, an oscillating scalaron component (the condensate of scalarons with the zero momentum) arises around the present moment when the scalaron mass is comparable to the Hubble constant $H_0$ (in the units where $\\hbar = c=1$), and it quickly becomes dominant over the non-relativistic matter component (cold dark matter and baryons) in the future. But its effective energy-momentum tensor in turn soon becomes negligible compared to an effective cosmological constant producing the future stable de Sitter stage. For less smooth initial conditions, more phantom boundary crossings may occur in the past. But these initial conditions are hardly compatible with the standard cosmology of the early Universe confirmed by numerous observational data. We hope to return to this question elsewhere. Finally, though the very phenomenon of multiple (and even infinite) number of phantom boundary crossings in the future is not directly observable, it is very interesting and important from the theoretical point of view. Also, as follows from our numerical calculations of the full evolution from the remote past to the remote future, the scalaron mass at the future de Sitter stage is close to its present value. Therefore, in principle it is possible to check from observational data describing the present and the past of our Universe if the derived analytical criterion for the existence of an infinite number of oscillations in $w_{\\DE}$ is satisfied or not." }, "1101/1101.5357_arXiv.txt": { "abstract": "We study the low-frequency timing properties and the spectral state evolution of the transient neutron star low-mass X-ray binary EXO 1745--248 using the entire {\\it Rossi X-ray Timing Explorer} Proportional Counter Array data. We tentatively conclude that EXO 1745--248 is an atoll source, and report the discovery of a $\\approx 0.45$ Hz low-frequency quasi-periodic oscillation and $\\sim 10$ Hz peaked noises. If it is an atoll, this source is unusual because (1) instead of a `C'-like curve, it traced a clear overall clockwise hysteresis curve in each of the colour-colour diagram and the hardness-intensity diagram; and (2) the source took at least 2.5 months to trace the softer banana state, as opposed to a few hours to a day, which is typical for an atoll source. The shape of the hysteresis track was intermediate between the characteristic `q'-like curves of several black hole systems and `C'-like curves of atolls, implying that EXO 1745--248 is an important source for the unification of the black hole and neutron star accretion processes. ", "introduction": "The spectral states and the correlated timing properties of neutron star and black hole low-mass X-ray binaries (LMXBs) can be very useful to understand the extreme environments of these sources \\citep{vanderKlis2006}. An excellent way to study these properties is to track these sources in the colour-colour diagram (CD; hard colour (HC) vs. soft colour (SC)) and in the hardness-intensity diagram (HID; hard colour vs. intensity; see \\S~\\ref{DataAnalysisandResults}). From the beginning of an outburst, the intensity of a transient black hole source increases, typically keeping the HC at a near-constant value. Near the highest intensity the HC value quickly decreases, followed by an intensity decrease at a lower HC value, and a soft-to-hard transition at a lower intensity value. Thus a black hole LMXB typically traces a `q'-like hysteresis curve in the HID \\citep{vanderKlis2006, Belloni2009}. It is usually believed that neutron star LMXBs do not trace hysteresis curves in CD/HIDs \\citep{vanderKlis2006}. For example, the near-Eddington Z sources trace out roughly `Z' shaped tracks on time scales of hours to a day, while the less luminous atoll sources have `C' shaped tracks \\citep{vanStraatenetal2003, vanderKlis2006}. The lower HC banana-like portion (BS) of the `C' track can be divided into upper banana (UB), lower banana (LB) and lower left banana (LLB) based on spectral and timing properties. The BS is traced out on time scales of hours to a day without any hysteresis \\citep{vanderKlis2006}. On the other hand, the higher HC extreme island state (EIS) is traced out in days to weeks, and secular motions in the form of parallel tracks are seen in EIS. An atoll source moves from EIS to BS via an island state (IS). Probably the only transient atoll source showing a `q'-like hysteresis HID curve is Aql X-1 (\\citet{MaitraBailyn2004, Reigetal2004}; see also \\citet{Bellonietal2007} for 4U 1636--53 tracks). Such neutron star LMXBs, and more importantly sources showing intermediate tracks between `q' and `C', can be very useful (1) to unify the black hole and neutron star accretion processes, and (2) to sort out the mismatch between the standard EIS-IS-BS framework and the general hysteresis phenomena. In this Letter, we show that the bursting neutron star LMXB EXO 1745--248 \\citep{MarkwardtSwank2000, Wijnandsetal2002, Heinkeetal2003} is such an intermediate source with unique properties. ", "conclusions": "\\label{Discussion} In this Letter, we have studied the evolution of spectral states of the neutron star LMXB EXO 1745--248. We tentatively conclude that it is an atoll source because of the following reasons. (1) From the spectral fitting, we find that the observed $2-30$ keV unabsorbed source flux varied in the range $(0.05-2.12)\\times10^{-8}$ ergs cm$^{-2}$ s$^{-1}$. Such a large intensity variation does not happen in a source, which shows an exclusive `Z' behaviour \\citep{vanderKlis2006}. (2) The hard colours (Fig.~\\ref{HID}) of EXO 1745--248 are consistent with those of atoll sources, but different from Z sources \\citep{Munoetal2002}. (3) The source shows parallel tracks for the higher hard colour values in HID (Fig.~\\ref{HID}), which are typical of atoll sources \\citep{vanderKlis2006}. (4) Shape of the CD track for lower hard colour values looks like a banana (Fig.~\\ref{CCD}). (5) Hard state to soft state transition of the source was plausibly quick \\citep{vanderKlis2006}. (6) PRE bursts were found in the softer state (Fig.~\\ref{CCD-HID}), as usually observed for fast spinning neutron star LMXBs \\citep{Munoetal2004}. (7) The kHz QPO was observed in the transitional state (plausibly LB/LLB), which is usual for atolls \\citep{MaitraBailyn2004, vanderKlis2006}. (8) VLFN at $< 1$ HZ was observed in BS, and $\\sim 10$ Hz peaked noise was detected in the transitional state (plausibly LB/LLB), which are usual for atolls \\citep{MaitraBailyn2004, vanderKlis2006}. (9) Red noise RMS is higher in the hard state \\citep{vanderKlis2006}. However, although we cannot confirm, there is some chance that at the most intense state, the source transformed into a Z source (e.g., \\citet{Homanetal2010}). This is because, the estimated source luminosity in this state was $\\sim 0.5$ times the Eddington luminosity \\citep{vanderKlis2006}, for a $2-30$ keV flux of $2.12\\times10^{-8}$ ergs cm$^{-2}$ s$^{-1}$, and assuming a 5.5 kpc source distance, 1.4 M$_\\odot$ neutron star mass, 6.0 stellar radius-to-mass ratio and ionized hydrogenic accreted matter. EXO 1745--248 is very interesting, unusual and important for the following reasons. (1) The source exhibited a clear overall clockwise hysteresis in HID and CD (Fig.~\\ref{HID} and \\ref{CCD}). A local anti-clockwise hysteresis is also observed in the high intensity state (plausibly UB). (2) In the hard state (plausibly EIS), unlike a typical atoll, the hard colour changed largely, and no horizontal track is present at the highest hard colour in HID \\citep{vanderKlis2006}. Moreover, the hard-to-soft transition involved a large change in intensity, unlike several black hole sources. These caused an HID-track-shape intermediate between atoll `C' tracks and black hole `q' tracks. (3) In CD/HID, the source moved from EIS to UB, while usually an atoll moves to LB/LLB from EIS \\citep{vanStraatenetal2003, vanderKlis2006}. (4) The CD/HID tracks of EXO 1745--248 could be segmented in time (Fig.~\\ref{HID} and \\ref{CCD}). Several segments can be distinguished by timing properties (Fig.~\\ref{LF-Powspec}), which shows that these segments are actually in different states, i.e., not in the same state shifted by secular motions. The source typically dwells in a segment for a few days to about a month (Table 1). (5) EXO 1745--248 took at least 2.5 months to trace the BS, as opposed to a few hours to a day, which is typical for an atoll source \\citep{vanderKlis2006}. The CD/HID hysteresis tracks of EXO 1745--248 could be very useful to relate the accretion processes in neutron star systems and black hole systems (\\S~\\ref{Introduction}). Finally, the HID hysteresis track of EXO 1745--248, which is intermediate between the `q'-like hysteresis track of Aql X-1 and `C'-like non-hysteresis tracks of most atoll sources, suggests that the popular EIS-IS-BS framework of `C'-like tracks might be a special case of a more general hysteresis behaviour. However, observations of more such intermediate sources are required to verify this." }, "1101/1101.1097_arXiv.txt": { "abstract": "We present new radial velocities from Keck Observatory and both Newtonian and Keplerian solutions for the triple-planet system orbiting HD 37124. The orbital solution for the system has improved dramatically since the third planet was first reported in \\citet{Vogt05} with an ambiguous orbital period. We have resolved this ambiguity, and show that the outer two planets have an apparent period commensurability of 2:1. A dynamical analysis finds both resonant and non-resonant configurations consistent with the radial velocity data, and constrains the mutual inclinations of the planets to be $< \\sim 30^\\circ$. We discuss HD 37124 in the context of the other 19 exoplanetary systems with apparent period commenserabilities, which we summarize in a table. We show that roughly one in three well-characterized multiplanet systems has a apparent low-order period commensuribility, which is more than would na\\\"ively be expected if the periods of exoplanets in known multiplanet systems were drawn randomly from the observed distribution of planetary orbital periods. ", "introduction": "\\label{Intro} To date, over 50 exoplanetary systems with more than one planet have been discovered, including: the extraordinary detections of the first exoplanets orbiting the pulsar PSR B1257+12 \\citep{Wolszczan92,Wolszczan94}; the imaged system orbiting HR 8799; those discovered during the microlensing event OGLE-2006-BLG-109L \\citep{Gaudi08}; several systems discovered by transit, including four or five\\footnote{KOI 877 may be a blend of two, separately transiting systems.} {\\it multiply} transiting systems from the Kepler mission \\citep{Steffen10}; and 43 systems discovered by radial velocity (RV) searches \\citep{Wright09d}. The RV systems include the four-planet systems $\\mu$ Ara \\citep{Santos04b,Pepe07}, GJ 581 \\citep{Mayor09} and GJ 876 \\citep{Rivera05,Rivera10} and the five-planet system orbiting 55 Cancri \\citep{Fischer08}. Of all these multplanet systems, only four are known to host three or more giant\\footnote{$\\msini > 0.2 \\Mjupmath$} planets with well-determined orbital paramaters: $\\upsilon$ And \\citep{Butler_upsand}, HIP 14810 \\citep{Wright09c}, $\\mu$ Ara \\citep{Pepe07} and HD 37124 \\citep{Vogt05}. HD 37124 (HIP 26381) is a 0.85 \\Msol\\ metal-poor \\citep[\\feh=-0.44,][]{SPOCS} G4 dwarf (V=7.7). \\cite{Vogt00} announced the a Jovian, $P \\sim 150$ d planet orbiting HD 37124 from HIRES data taken at Keck Observatory as part of the California and Carnegie Planet Search. Further monitoring of the star revealed substantial long-term residuals. \\cite{Butler03} fit these residuals with an eccentric, 1940 d planet, but noted that the solution was not unique (and \\cite{Gozdziewski03} showed that this fit was, in fact, unstable.) After collecting two more years of data, \\citet{Vogt05} was able to report the detection of a third planet in the system, though with an ambiguity: while the $b$ and $c$ components had clearly defined periods, the $d$ component could be fit nearly equally well with periods of either 2300 d or 29.32 d, the latter likely being an alias due to the lunar cycle.\\footnote{Time on the Keck telescopes dedicated to observing bright, planet search targets with HIRES is usually assigned during bright or gray time; the resulting scarcity of data points during new moon can interact with planetary signals to create spurious, aliased solutions.} \\citet{Wright09d} reported that recent Keck velocities had resolved the ambiguity qualitatively in favor of the longer orbital period. \\citet{Gozdziewski06b} explored the many possible dynamical configurations consistent with the \\citet{Vogt05} velocities, including many resonant solutions. \\citet{Gozdziewski08} used the system to demonstrate a fast MENGO algorithm, but they did not explore the 2:1 resonance, as the data did not seem to favor it at the time. We present new Keck observations, and these data provide a unique orbital solution for the outer planet. The outer planet period we find is more consistent with the original period reported by \\citet{Butler03} than the refined orbit of \\citet{Vogt05}\\footnote{\\citet{Vogt05} opted to refer to the new, 840 d signal as the $c$ component, despite the prior 1940 d fit of \\citet{Butler03}, because that prior fit was so speculative, and because their new fit put the very existence of a 1940 d periodicity in some doubt.} (though we find a much lower eccentricity). Herein, we present the entire history of Keck velocities obtained for this star, and present self-consistent orbital solutions showing that the outer two planets are in or very near a 2:1 mean-motion resonance (MMR). This is the 20th exoplanetary system to be found near an MMR, and only the tenth system with an apparent 2:1 commensurability. Period commensurabilities (PCs) represent important dynamical indicators in the Solar System and have been linked with observables and formation mechanisms \\citep{Goldreich65}. The near-5:2 PC of Jupiter and Saturn, also known as ``The Great Inequality\", might be the remnant of a divergent resonant crossing that produced the current architecture of the outer Solar System, the Late Heavy Bombardment, and the Trojan Asteroids \\citep{Gomes05,Morbidelli05,Tsiganis05,Tsiganis05b}. The populations of the asteroid belt and the Kuiper Belt, exemplified by the PC and near-PC-populated Kirkwood Gaps \\citep[e.g.][]{Tsiganis02} the Plutinos (3:2 PCs with Pluto and Neptune) and the twotinos \\citep[2:1 PCs with Pluto and Neptune; e.g. ][]{Murray-Clay05,Chiang02b}, have implications for the migratory history of Jupiter and Neptune and the prospect of, e.g. secular resonant sweeping \\citep[e.g., ][]{Nagasawa08}. Near-PCs found in satellite and ring systems have had direct observational consequences; the Saturnian satellite Pandora was $\\sim 19^{\\circ}$ behind its predicted orbital longitude in a 1995 ring plane crossing \\citep{French03} due to its 121:118 PC with neighboring satellite Prometheus. By extension, we may anticipate similar importance in the growing number of exoplanetary systems exhibiting PCs. In extrasolar systems, Mean Motion Resonances (MMRs) have been interpreted as the indication of convergent migration in multi-planet systems, \\citep[e.g.][]{Thommes03,KleyResonance,Papaloizou05}. Several subsequent studies \\citep{Beauge06, Terquem07, Pierens08,Podlewska08,Podlewska09,Libert09,Rein09, Papaloizou10,Rein10,Zhang10b,Zhang10c} exploring convergent migration for a variety of masses, separations and disk properties have found many regions of mass and orbital element phase space in which planets are easily captured through this mechanism. ", "conclusions": "We have resolved the period ambiguity of HD 37124 $d$ from \\citet{Vogt05} and find that HD 37124 $c$ and $d$ are in an apparent 2:1 period commensurability. Our numerical integrations show that both resonant and non-resonant configurations are consistent with the radial velocity data, and that stability requires a nearly circular orbit ($e < 0.3$) for the $d$ component. Our stability analysis shows that the system must be nearly coplanar, and that the three planets have identical minimum masses within the errors (of 3--10\\%). We show that the roughly one in three well-characterized multiplanet systems shows an apparent period commensurability, which is more than a na\\\"ive estimate based on randomly drawing periods from the known exoplanet population would suggest. This offers evidence for some particular proposed scattering and migration mechanisms, and we suggest that the statistics of multiplanet systems may now be sufficiently robust to provide a test and comparison of models of exoplanetary dynamical evolution." }, "1101/1101.1574_arXiv.txt": { "abstract": "{Stars, and more particularly massive stars, have a drastic impact on galaxy evolution. Yet the conditions in which they form and collapse are still not fully understood.} {In particular, the influence of the magnetic field on the collapse of massive clumps is relatively unexplored, it is therefore of great relevance in the context of the formation of massive stars to investigate its impact. } {We perform high resolution, MHD simulations of the collapse of one hundred solar masses, turbulent and magnetized clouds, with the adaptive mesh refinement code RAMSES. We compute various quantities such as mass distribution, magnetic field, and angular momentum within the collapsing core and study the episodic outflows and the fragmentation that occurs during the collapse.} {The magnetic field has a drastic impact on the cloud evolution. We find that magnetic braking is able to substantially reduce the angular momentum in the inner part of the collapsing cloud. Fast and episodic outflows are being launched with typical velocities of the order of 1-3 km s$^{-1}$, although the highest velocities can be as high as 20-40 km s$^{-1}$. The fragmentation in several objects is reduced in substantially magnetized clouds with respect to hydrodynamical ones by a factor of the order of 1.5-2.} {We conclude that magnetic fields have a significant impact on the evolution of massive clumps. In combination with radiation, magnetic fields largely determine the outcome of massive core collapse. We stress that numerical convergence of MHD collapse is a challenging issue. In particular, numerical diffusion appears to be important at high density and therefore could possibly lead to an overestimation of the number of fragments.} ", "introduction": "It is believed that stars form during the collapse of prestellar cores inside molecular clouds. Understanding this process is of great relevance as it determines the initial conditions of the protostars as well as the properties of accretion disks which form in their vicinity. It is also during this process that the fragmentation, that is the formation of binaries and clusters rather than a single object, may occur. During the last decades, many studies have been investigating the fragmentation of dense cores using either smooth particles hydrodynamic (SPH) based codes or grid codes (see e.g. Matsumoto \\& Hanawa, 2003, Commer{\\c c}on et al. 2008, or Goodwin et al. 2007 for a review). Until recently, most works have been neglecting the magnetic field and assume an isothermal equation of state until the gas becomes optically thick. Under these conditions, various studies infer that the massive cores fragment into several objects. Simulations like the ones performed by Bonnell et al. (2004), Klessen \\& Burkert (2000, 2001) and Dobbs et al. (2005) generally find that the number of fragments is comparable with or even exceeds the number of initial Jeans masses within the clouds, which implies that a massive core may result in a cluster that contains tens objects or even more. Observationally the question as to whether massive cores are fragmented is difficult to investigate because of the large distances at which these objects are located. Preliminary investigations do not appear to show such high levels of fragmentation. For example, Bontemps et al. (2010) report 1700 AU-resolution observations using PdBI of IR-quiet massive cores in Cygnus X, and find that although one of them does break up to some degree when observed at high resolution, most of them do not have most of their collapsed mass in low mass objects. Some of them do not break up at all, and remain single compact objects even at 1700 AU resolution. Recent submillimeter array observations by Longmore et al. (2010) reach similar", "conclusions": "We performed high resolution numerical simulations of collapsing magnetized and turbulent hundred solar masses cores assuming a barotropic equation of state using the RAMSES code. Three different magnetic intensities corresponding to mass-to-flux ratio, $\\mu$ equal to 120, 5, and 2 were explored. The simulations were repeated with two different resolutions to investigate the impact of the numerical method and the issue of numerical convergence. These simulations confirm the strong impact that the magnetic field has, in particular regarding the byproduct of the collapse. The main effects of the magnetic field are i) to significantly reduce the angular momentum in the inner part of the cloud, ii) to launch episodic and relatively fast outflows, even when the value of the magnetic intensity is initially weak, iii) to reduce the fragmentation of the cloud in several objects (by about a factor 2 when $\\mu$, the mass-to-flux ratio is equal to 2). While the collapse is relatively well organized in the outer part of the cloud and exhibites a classical $r^{-2}$ density profile, the inner part is very turbulent and the infall is dominated by high velocity fluctuations. In this region, the density profile is steeper and typically goes as $r^{-\\simeq 2.5}$. The magnetic field is amplified by gravitational contraction that leads to roughly $B \\propto \\sqrt{\\rho}$, which in turn implies that the Alfv\\'en velocity is nearly constant on average although it significantly fluctuates at all scales. When the magnetic field is very weak ($\\mu=120$), the amplification is stronger which makes in the cloud inner part the Alfv\\'en speed of the order of the sound speed. The outflows appear to be episodic and are usually non-bipolar. Not only their velocities evolve with time, but there are events of intense ejections followed by periods without significant outflow motions. The typical velocity of these flows is of the order of 1-3 km s$^{-1}$ but much higher velocities (5 to 10 times higher) can be reached for a small fraction of the mass, in particular when the field is weak. The total mass they carry is of the order of a tenth to a few solar masses, depending on the time and the resolution. There is a clear influence of the numerical resolution, implying that these numbers are {\\it probably} underestimated. The strongly magnetized clouds tend to fragment less (factor 1.5-2) than the weakly magnetized ones, implying that the mass is more concentrated in the more massive stars. The region in which fragmentation occurs is also more compact when the magnetic intensity is stronger. We stress however that numerical diffusion is clearly reducing the magnetic flux in the dense part of the clouds, which makes it possible that the fragmentation is indeed overestimated for $\\mu=5$ and 2." }, "1101/1101.3571_arXiv.txt": { "abstract": "We analyze the intrinsic polarization of two classical Be stars in the process of losing their circumstellar disks via a Be to normal B star transition originally reported by Wisniewski et al. During each of five polarimetric outbursts which interupt these disk-loss events, we find that the ratio of the polarization across the Balmer jump (BJ+/BJ-) versus the V-band polarization traces a distinct loop structure as a function of time. Since the polarization change across the Balmer jump is a tracer of the innermost disk density whereas the V-band polarization is a tracer of the total scattering mass of the disk, we suggest such correlated loop structures in Balmer jump-V band polarization diagrams (BJV diagragms) provide a unique diagnostic of the radial distribution of mass within Be disks. We use the 3-D Monte Carlo radiation transfer code HDUST to reproduce the observed clockwise loops simply by turning ``on/off'' the mass decretion from the disk. We speculate that counter-clockwise loop structures we observe in BJV diagrams might be caused by the mass decretion rate changing between subsequent ``on/off'' sequences. Applying this new diagnostic to a larger sample of Be disk systems will provide insight into the time-dependent nature of each system's stellar decretion rate. ", "introduction": "\\label{intro} Classical Be stars are well known to be characterized by having gaseous circumstellar decretion disks which are fed from mass-loss from their rapidly rotating central stars \\citep{por03}. As detailed in the review of \\citet{car10}, the viscous decretion disk model (VDDM; \\citealt{lee91}) is adept at explaining many of the observational properties of Be star disks, including the observed Keplerian rotation of the disk, and is generally considered the most promising model to explain the Be phenomenon, although alternate models also have been proposed \\citep{bjo93, cas02}. The mechanism(s) responsible for ejecting material from the stellar photosphere into a disk is also unkown, although both observational \\citep{riv98,nei02} and theoretical \\citep{and86,cra09} studies suggest that non-radial pulsations might act to feed at least some of these decretion disks. While some classical Be stars exhibit observational evidence of large-scale asymmetries \\citep{oka97, vak98, ste09} interpreted as arising from one-armed density waves \\citep{oka91}, the general gas disk density has typically been modeled by a very simple axisymmetric power law (see e.g. \\citealt{bjo97,por99}). Recent contemporaneous optical-infrared (IR) spectroscopic \\citep{wi07a}, polarimetric \\citep{car07}, and optical-IR interferometric studies \\citep{tyc06,gie07, tyc08, pot10}, which are each most sensitive to different physical regions of disk, demonstrate how the radial-dependence of the gas density in these disks can be observationally constrained. Such works offer promise for testing the appropriateness of the single power-law adopted by most models. Some classical Be stars can exhibit stable decretion disks for decades (e.g $\\zeta$ Tau; \\citealt{ste09}); however, they are also known to experience aperiodic ``Be to normal B to Be'' transitions whereby they lose (and subsequently regenerate) all observational signatures of having a disk \\citep{und82,cla03,vin06}. The frequency of these transitions is not well constrained by observations, although the discovery of 12 new transient Be stars in multi-epoch study of eight open clusters suggests that these major events are not rare \\citep{mcs09}. These systems represent ideal testbeds to diagnose the fundamental mechanism which drives disk formation in Be stars precisely because they are known to be actively losing (or gaining) a disk. Studying such disks with techniques capable of diagnosing the radial-dependence of the gas density, especially in the inner-most regions of the disk, could enable an enhanced understanding of how these disks form. In paper I of this series \\citep{wis10}, we analyzed $\\sim$15 years of spectropolarimetric observations of the classical Be stars 60 Cygni and $\\pi$ Aquarii which covered one disk-loss episode in each star, and discussed the time-scale and overall evolution of these events. In this paper, we present a first look at the behavior of the intrinsic polarization during these events and detail a unique new diagnostic which effectively traces the gas density in the inner-most region of the disk. We investigate these ``polarization-loop'' signatures with a well vetted 3-D Monte Carlo NLTE code and present representative model runs which reproduce the observations and support our interpretation of this phenomenon. Finally, we outline our plans to perform more detailed modeling of these signatures. ", "conclusions": "\\subsection{Modeling the Observed Phenomenon} \\label{model} As mentioned in the introduction, the viscous decretion disk model (VDDM) is currently the most promising candidate to explain the structure, formation and evolution of Be disks. While most theoretical approaches so far considered the case of a constant mass decretion rate in the quasi-steady state limit \\citep[e.g.,][]{bjo97,bjo05}, more recent studies investigated the temporal evolution of the disk surface density fed by an arbitrary mass decretion history \\citep{oka07,jon08}. Here we qualitatively investigate whether the VDDM can account for the trends we see in BJV diagrams. In order to theoretically reproduce the observed trends shown in Figure 1 with the VDDM model we used a 1-D hydrodynamical code \\citep{oka07} to compute the time-dependent surface density of the disk. This code solves the viscous diffusion problem given a prescription for the stellar mass decretion rate, and a value for the disk kinematic viscosity (the $\\alpha$ parameter of \\citealt{sha73}). The surface density for chosen epochs of the disk evolution is then fed to our radiative transfer code {\\sc hdust} \\citep{car06} that is capable of turning the structural information thus provided into astrophysical observables, such as emergent spectrum or intensity maps on the sky. The correlation loops observed in the BJV diagram can be qualitatively reproduced assuming a prescription for the mass decretion rate that involves alternating cycles of activity (mass decretion on) and quiescence (mass decretion off). One such an example is shown Fig.~\\ref{modelfig}. Starting from no preexisting disk, this sample model assumes a 3-year long period of activity followed by a 3-year long quiescence. This 6-year cycle were repeated many times. In Fig.~\\ref{modelfig} we show results covering the period between 32 and 39 years after the beginning of disk formation. The polarization forms a clockwise loop in the BJV diagram that can be described as follows. At the end of the active phase, the star had built a large and dense circumstellar disk (phase 1). When the mass decretion stops, the inner disk quickly reaccretes back onto the star; this causes a fast drop of BJ size and the curve follows a track towards the bottom-left of the BJV diagram (phase 2). What follows is a slow secular dissipation of the entire disk along which the BJ size changes little (the inner disk having already reached very low densities) but the V-band polarization diminishes as the disk mass decreases (phase 2 to 3). When the next cycle of activity begins (phase 3) the inner disk quickly fills up again and the curve eventually reaches back the top of the BJV diagram. The detailed shape of the loop depends on several factors: the viewing angle (Fig.~\\ref{modelfig}), the value of $\\alpha$, and the mass decretion history assumed. In the simple model shown here the loops nearly close, because the mass decretion rate assumed for each cycle is the same. The loop would not close if the mass decretion rate of subsequent cycles were to be different and/or if the length of the active/quiescent phases were irregular. We plan to systematically explore a large range of mass decretion rate scenarios in a future publication (Haubois et al 2011, in prep). \\subsection{Comparison with CMD loops} \\citet{dew06} studied the photometric variability of several hundred Be stars in the Magellanic Clouds and found $\\sim$100 stars whose photometric variations traced loop-like patterns in optical color magnitude diagrams (CMDs). Most ($\\sim$90\\%) objects traced clockwise loops in these CMDs while $\\sim$10\\% traced counter-clockwise loops. \\citet{dew06} suggested that clockwise loops were indicative of systems actively decreting material from their stellar surfaces whereas counter-clockwise loops were indicative of systems in which material was being re-accreted onto the central star. We note that our models of the loop-like structures we observe in BJV diagrams include the effects of both decretion of material from the central star and the re-accretion of material onto the central star. Hence the large-scale morphological differences we observe, i.e. clockwise versus counter-clockwise loops in BJV diagrams, can not be simply attributed to differences in the radial direction of motion of material in the disk as invoked by \\citet{dew06}. Rather, we speculate that the counter-clockwise loops we sometimes observe in BJV diagrams might be caused by a non-constant mass decretion rate or two nearby cycles which have significantly different mass decretion rates. Our future systematic exploration of the parameter space of our models (Haubois et al 2011, in prep) will enable us to test this speculative hypothesis, as well as other mass ejection scenarios, to explain these counter-clockwise loops. \\subsection{Implications and Future Applications of the Technique} The steady-state \\citep{bjo05} and time dependent \\citep{oka07} surface density of Be disks is driven in large part by the ratio of the stellar mass-loss rate and the $\\alpha$ parameter (both quantities setting the disk decretion rate), although recent modeling work has demonstrated that the disk temperature can also influence the surface density, especially in the inner disk regions, when the stellar mass-loss rate is large \\citep{car08}. As \\citet{car09} noted, observationally constraining the stellar mass-loss rate is very challenging for systems whose disks are truncated by binary companions; moreover, constraining the detailed time dependent decretion rate of non-steady state disk systems (e.g. \\citealt{riv98}) is also challenging. The BJV (and broad-band BUV) diagnostic we have introduced in this paper offers one clear way to better diagnose the detailed time-dependent decretion rate of Be disk systems. Application of this technique, when the requisite low-resolution blue optical spectropolarimetry (BJV diagram) or U- and B-band filter polarimetry is available, to a larger sample of Be systems actively gaining/losing their disks would help elucidate the mechanism(s) responsible for triggering disk formation in Be stars. We encourage the community to obtain this type of well time sampled polarimetry for the mid-2011 periastron passage of $\\delta$ Sco, as it would be provide a powerful diagnostic of the type of decretion outbursts which have been reported in previous periastron passages \\citep{mir01,mir03}. Moreover, inclusion of simple V-band photometry in such analysis could help observationally determine the $\\alpha$ parameter \\citep{car10}." }, "1101/1101.3092_arXiv.txt": { "abstract": "We present the results of the VLT/VIMOS integral-field spectroscopic observations of the inner $28''\\times28''$ ($3.1\\;{\\rm kpc} \\times 3.1\\;{\\rm kpc}$) of the interacting spiral NGC~5719, which is known to host two co-spatial counter-rotating stellar discs. At each position in the field of view, the observed galaxy spectrum is decomposed into the contributions of the spectra of two stellar and one ionised-gas components. We measure the kinematics and the line strengths of the Lick indices of the two stellar counter-rotating components. We model the data of each stellar component with single stellar population models that account for the $\\alpha$/Fe overabundance. We also derive the distribution and kinematics of the ionised-gas disc, that is associated with the younger, less rich in metals, more $\\alpha$-enhanced, and less luminous stellar component. They are both counter-rotating with respect the main stellar body of the galaxy. These findings prove the scenario where gas was accreted first by NGC~5719 onto a retrograde orbit from the large reservoir available in its neighbourhoods as the result of the interaction with its companion NGC~5713, and subsequently fuelled the {\\em in situ\\/} formation of the counter-rotating stellar disc. ", "introduction": "The presence of stars counter-rotating with respect to other stars and/or gas has been detected in several disc galaxies and is commonly interpreted as the end result of a retrograde acquisition of external gas and subsequent star formation (see \\citealt{Bertola+99b} for a review). Nevertheless, some special cases of counter-rotating stellar discs could have an internal origin induced by the presence of a bar (e.g., \\citealt{Evans+94}). The demography of gaseous and stellar counter-rotating components in S0's and spirals is a key to understand their assembly process. The fraction of lenticular galaxies with a counter-rotating gaseous disc is consistent with the 50\\% that we expect if all the gas in S0's is of external origin \\citep{Bertola+92}. In contrast, less than 10\\% of them host a detectable fraction of counter-rotating stars \\citep{Kuijken+96}. Large-scale counter-rotation is a rare phenomenon in spirals. In fact, less than 10\\% of the studied spiral galaxies host a counter-rotating gaseous and/or stellar disc \\citep{Kannappan+01, Pizzella+04}. The particular case of large-scale stellar counter-rotation in disc galaxies has been observed only in NGC 4550 \\citep{Rubin+92, Rix+92, sauron3}, NGC~7217 \\citep{Merrifield+94}, NGC~3593 \\citep{Bertola+96, Corsini+98b, Garcia+00}, NGC~4138 \\citep{Jore+96}, and NGC~5719 \\citep{Vergani+07}. To interpret the observed frequencies of counter-rotations, \\citet{Pizzella+04} argue that the retrograde acquisition of small amounts of external gas can give rise to counter-rotating gaseous discs in gas-poor S0's only, while in gas-rich spirals the newly acquired gas is swept away by the pre-existing gas. Counter-rotating gaseous discs in spirals are formed only from the retrograde acquisition of amounts of gas larger than the pre-existing gas content. This is the case of the purely gaseous counter-rotating components detected in NGC~3626 \\citep{Ciri+95, Garcia+98} and NGC~4826 \\citep{Braun+94, Garcia+03}. Counter-rotating stellar discs are produced by subsequent star formation. Therefore, in this scenario counter-rotating stellar discs are expected to be made by younger stars than those of their host galaxy. This picture can be directly tested in the NGC~5719/13 galaxy pair, which has been recently studied by \\citet{Vergani+07}. NGC~5719 is an almost edge-on Sab galaxy with a prominent skewed dust lane at a distance of 23.2 Mpc. \\citet{Vergani+07} report a spectacular on-going interaction with its face-on Sbc companion NGC~5713. Two \\hi\\ tidal bridges loop around NGC~5719 and connect it to NGC~5713 at a projected distance of 77 kpc. The neutral and ionised hydrogen in the disc of NGC~5719 are counter-rotating with respect to the main stellar disc. The kinematics of the ionized-gas disc and of both the counter-rotating stellar discs are measured out to about $40''$ (4.3 kpc) from the galaxy centre. In conclusion, \\citet{Vergani+07} propose a scenario where \\hi\\ from the large reservoir available in the galactic surroundings was accreted by NGC~5719 onto a retrograde orbit and subsequently fuelled the {\\em in situ\\/} formation of the counter-rotating stellar disc. In this work, we will address the crucial piece of information which is still missing, i.e. proving that the stellar population of the counter-rotating disc of NGC~5719 is younger with respect to that of the stars in the galaxy main disc. \\begin{figure} \\hspace{-0.5cm} \\psfig{file=coccato_fig1.ps,width=8.7cm,clip=} \\caption{Fit of the galaxy spectrum ({\\it black}) in the spatial bin at $10''$ East from the galaxy center. The best fit model ({\\it green}) is the sum of the spectra of the ionised-gas component ({\\it magenta}) and the two stellar components ({\\it blue} and {\\it red}). The latter are obtained convolving the synthetic templates with the best fitting Gaussian LOSVDs and multiplying them by the best fitting Legendre polynomial. The differences in the position of absorption line features and in the \\hb\\ equivalent widths between the two stellar components (indicating different kinematics and stellar population content) are clearly evident.} \\label{fig:kinem_fit} \\end{figure} ", "conclusions": "\\label{sec:3} We present the results of the VLT/VIMOS integral-field spectroscopic observations of the inner $28''\\times28''$ ($3.1 {\\rm kpc} \\times 3.1 {\\rm kpc}$) of the spiral galaxy NGC~5719. The kinematics of the stars and ionised gas are very complex. The observed galaxy spectra are decomposed into the contributions of three distinct kinematic components characterised by a regular disc-like rotation: one main and one secondary stellar component and a ionised-gas component. The spectral decomposition is done using an implementation of the pPXF routine, which allows to measure {\\it simultaneously} the kinematics {\\em and\\/} stellar population properties of the two stellar components. The rotation of the main stellar component is receding towards East, like that of the stellar body of the galaxy as observed by \\citet{Vergani+07} out to $\\sim40''$. We measure a maximum rotation velocity of $\\sim150$ \\kms\\ at about $10''$ (1.1 kpc) from the centre. The secondary stellar and ionised-gas components are counter-rotating with respect to the main stellar component, with a maximum rotation of $\\sim200$ \\kms\\ at about $10''$ from the centre. We are able to resolve the two stellar components down to the innermost $\\sim2''$ (0.2 kpc). At smaller radii, they have a too small velocity separation to be resolved. The median values of light fraction contributed by the two components in the spatial bins where they are resolved are $F_{\\rm main} = 56\\%$ and $F_{\\rm secondary} = 44\\%$ (with $20\\%$ standard deviation), respectively. The ionised gas is detected all over the observed field of view. It is characterised by a strong \\hb\\ emission, which is concentrated in a twin-peaked morphology indicating an edge-on ring with a semi-major axis of $\\sim7''$ (0.8 kpc), and an outer asymmetric $m=1$ and fragmented spiral arc extending southwards with a semi-major axis of $\\sim13''$ (1.5 kpc; see contours in Figs. \\ref{fig:2dkin} and \\ref{fig:2dssp}). The two-stream fluid instability present in systems like NGC~3593 \\citep{Garcia+00} can favor this morphology. An hint of the inner ring is visible also in the \\hi\\ position-velocity diagram measured along the major axis of NGC~5719 \\citep{Vergani+07}. The secondary stellar component is associated to these features of the gas distribution, i.e. the counter-rotating stars are detected in a region enclosed by the ionised-gas structures. A similar phenomenon is observed in NGC~3593, where a concentrated gaseous ring is associated to the counter-rotating stellar disc \\citep{Corsini+98b, Garcia+00}. The two counter-rotating stellar components are characterised by different chemical properties as it results from the measured line strength of the Lick indices (Fig. \\ref{fig:indices}). On average, the main stellar component has lower \\hb\\ and higher [MgFe]$'$, \\mgb, and \\fe\\ with respect to the secondary component. This translates immediately into different properties of their stellar populations (Fig. \\ref{fig:2dssp}). The main stellar component has ages ranging from 2 to 13.5 Gyr (median $\\rm age = 4$ Gyr with 4 Gyr standard deviation). It has nearly-solar metallicity (median $\\rm [Z/H] = 0.08$ dex with 0.20 dex standard deviation). It displays super-solar enhancement (median $\\rm [\\alpha/Fe] = 0.10$ dex with 0.10 dex standard deviation). The counter-rotating stellar population component is younger, with ages ranging from 0.7 Gyr to 2.0 Gyr (median age = 1.3 Gyr with 0.6 Gyr standard deviation). Its metallicity shows a radial gradient: it changes from sub-solar ($\\rm [Z/H] \\sim -1.0$ dex) in the outskirts, to solar ($\\rm [Z/H] = 0.0$ dex) and super-solar ($\\rm [Z/H] \\sim 0.3$ dex) in the centre. The youngest ages and highest metallicities are found in correspondence of the regions where the \\hb\\ emission is more intense. The $\\alpha$-enhancement is super-solar (median $\\rm [\\alpha/Fe] = 0.14$ dex with 0.05 dex standard deviation). These findings extend the results by \\citet{Neff+05} based on GALEX observations. They reported the presence of a young stellar component in the disc of NGC~5719 by analysing UV and optical images. Unfortunately, the poor GALEX angular resolution prevented them to associate it to the \\hb\\ ring as we successfully do (Fig. \\ref{fig:2dssp}). Our estimate is in agreement with their lower limit to the actual age of such a young population ($0.3\\pm0.1$ Gyr). With our new observations and spectral decomposition technique, we prove that the mean age of the counter-rotating disc, which is associated to the neutral and ionised gas disc, is indeed younger than the main stellar disc. This result shows that counter-rotating disc has been recently assembled. The median overabundance of the counter-rotating component (0.14 dex) indicates a star formation history with a time-scale of 2 Gyr \\citep{Thomas+05}. More details about the assembly process could be derived by a further analysis of the east-west asymmetries measured in the maps of the stellar-population properties. The scenario proposed by \\citet{Vergani+07} that NGC~5719 hosts a counter-rotating stellar disc originated from the gas accreted during the ongoing merging with its companion NGC~5713, is finally confirmed. \\label{lastpage}" }, "1101/1101.1604_arXiv.txt": { "abstract": "We present preliminary results of an optical-UV survey of the North Celestial Cap (NCCS) based on $\\sim$5\\% areal coverage. The NCCS will provide good photometric and astrometric data for the North Celestial Cap region ($80^{\\circ}\\leq\\delta\\leq90^{\\circ}$). This region, at galactic latitudes from $17^{\\circ}\\lesssim b\\lesssim37^{\\circ}$, is poorly covered by modern CCD-based surveys. The expected number of detected objects in NCCS is $\\sim$1,500,000. We discuss issues of galactic structure, extinction, and the galaxy clustering in the colour-colour diagrams. ", "introduction": "\\vspace{-8pt} From the dawn of humankind people were interested in studying the surrounding world. In ancient times astronomy, and sky mapping in particular, had not only a world-description function, but had also practical purposes. For example, sky maps were used for navigation and orientation as well as for predicting celestial phenomena. Two inventions contributed significantly to sky mapping: the telescope and the permanent photography. Photographic catalogs, produced from the 1890s to the 2000s, play an important role even today. The first high-quality digital all-sky survey, the \\textbf{Digitized Sky Survey} (DSS) and its extension DSS-II, were produced by scanning photographic survey plates (POSS-I, POSS-II, ESO/SERC) with specific photometric calibrations. The POSS-I, POSS-II and ESO/SERC photographic surveys provided high-precision astrometry used in catalogs such as \\textbf{USNO-A1.0}, \\textbf{USNO-A2.0}, \\textbf{USNO-B1.0}, etc. USNO-A and USNO-B are all-sky high-precision astrometric catalogs including also photometric data from the DSS and DSS-II. In many cases, the USNO DSS-based catalogs are the only optical high-precision data available for a particular sky area. However, photographic catalogs suffer from significant photometric and astrometric systematic and statistical errors due to shortcomings of the photographic emulsion, such as low sensitivity, limited dynamic range and non-linearity (\\citealt{MON03}). These were solved with the introduction of the charge-coupled device (CCD) to astronomy. CCDs eliminated not only the photographic emulsion shortcomings, but solved also other problems, such as providing an effective raw data storage, yielding a short delay between the raw data collection and the final data extraction, and being a reusable photosensitive element. All modern optical sky surveys are CCD-based. The modern optical sky survey in general use is the \\textbf{Sloan Digital Sky Survey} (SDSS), which covers about 10,000 deg$^2$ of the sky and provides high-precision photometric and astrometric data in five Sloan bands (\\citealt{ABA09}). An important extension to the optical catalogs are UV observations. One of the prominent available UV instruments is the \\textbf{Galaxy Evolution Explorer} (GALEX), performing both imaging and low-resolution spectroscopy in two bands: near-UV (NUV). GALEX conducts several pioneering UV sky surveys aimed primarily to understand galaxy evolution, and its publicly-available dataset covers by now about 3/4 of sky (\\citealt{MOR07}). \\vspace{-8pt} ", "conclusions": "\\vspace{-8pt} We have shown that there is valuable new science that can be derived from the NCCS data set, apart from improving the global knowledge about the sky at high declinations. It is possible, therefore, to consider possible extensions of this rather limited project, as defined above. We can relatively easily extend the surveyed region for five more degrees of declination, increasing the surveyed region to more than 700 deg$^2$. This will increase the overlap with SDSS-surveyed regions, allowing an improved cross-calibration between the two surveys. Alternatively, we could repeat the observations of the already surveyed region in one of the R and I bands to check the variability of the detected sources. We could merge the NCCS data with 2MASS extending the wavelength base to IR. Though 2MASS is relatively shallow, it could provide additional information about NCCS objects brighter than $\\sim$16$^m$--17$^m$. Spectroscopic follow-up of the promising NCCS sources could be performed at the Wise Observatory for objects brighter than $\\sim$16$^m$ or at a larger telescope for the faint ones. \\vspace{-8pt}" }, "1101/1101.3437_arXiv.txt": { "abstract": "We present ULTRACAM photometry of ES Cet, an ultracompact binary with a $620$s orbital period. The mass transfer in systems such as this one is thought to be driven by gravitational radiation, which causes the binary to evolve to longer periods since the semi-degenerate donor star expands in size as it loses mass. We supplement these ULTRACAM+WHT data with observations made with smaller telescopes around the world over a nine year baseline. All of the observations show variation on the orbital period, and by timing this variation we track the period evolution of this system. We do not detect any significant departure from a linear ephemeris, implying a donor star that is of small mass and close to a fully degenerate state. This finding favours the double white dwarf formation channel for this AM CVn star. An alternative explanation is that the system is in the relatively short-lived phase in which the mass transfer rate climbs towards its long-term value. ", "introduction": "Ultracompact binaries with periods of the order of tens of minutes or less have attracted much attention in recent years. These systems consist of a primary white dwarf with a companion star that is also at least partially degenerate. Close double-degenerate binaries are one of the proposed progenitor populations of Type Ia supernovae \\citep{Kotak08,Distefano10,Gilfanov10}, as well as the recently proposed sub-luminous `.Ia' supernovae \\citep{Bildsten07,Kasliwal10}. They are of interest from a binary formation and evolution point of view, with the short periods implying at least one common envelope phase in the history of the binary, and the chemical composition suggesting helium white dwarfs, helium stars or cataclysmic variables (CVs) with evolved secondaries as possible progenitors \\citep{Nelemans01,Nelemans10}. The mass transfer in these systems is thought to be driven by angular momentum loss as a result of gravitational radiation. These sources are predicted to be among the strongest gravitational wave sources in the sky \\citep{Nelemans04}, and are the only class of binary with examples already known which will be detectable by the gravitational wave observatory LISA \\citep{Stroeer06,Roelofs07}. The two systems with the shortest known periods are HM Cnc ($324$s, \\citealt{Roelofs10}) and V407 Vul ($569$s, \\citealt{Haberl95}). The exact nature of these systems is unclear: the two leading models are the unipolar induction (UI) model \\citep{Wu02} and the direct-impact accretion model \\citep{Marsh02,Roelofs10}. There are also 23 known systems with longer periods ranging from $620$ to $3906$s (the AM Canum Venaticorum stars, AM CVns; see \\citealt{Solheim10} for a recent review). The natures of these systems are better established: the spectra show an absence of hydrogen and the presence of helium lines, many of which are the double-peaked emission lines characteristic of sources accreting via a disc \\citep{Marsh99,MoralesRueda03}. These systems are clearly accreting, and are the helium equivalent of the cataclysmic variable (CV) stars. Gravitational radiation has a huge influence on these systems, driving the evolution and determining the orbital period distribution, luminosities and numbers. Measurements of the time derivative of the orbital period $P$ for V407 Vul and HM Cnc have shown the periods in these systems to be decreasing \\citep{Hakala03, Hakala04, Ramsay05, Strohmayer05}. This is consistent with the UI model, and contrary to what might be expected for accretion. In an accreting double-degenerate system, the gravitational radiation should drive the binary towards longer orbital periods. However, if the mass transfer in an accreting double-degenerate system is significantly lower than its equilibrium value, due to either some non-secular process \\citep{Marsh05} or the binary being in its mass transfer turn-on phase \\citep{Willems05,DAntona06,Deloye06}, then a decreasing binary period is possible. The counter-argument to this is that such phases are expected to be short, although \\citet{Deloye06} calculated that the early contact phase can last much longer ($10^3$ - $10^6$ yr) than originally thought. To date, no period derivative has been measured in any of the longer period AM CVn stars. Since these are unambiguously accreting binaries, then we would expect to detect an increasing period in these objects. Additionally, it is not certain that gravitational radiation is the dominant angular momentum loss mechanism for the ultra-compact binary stars: some other mechanism, such as the magnetic braking observed in cataclysmic variables, could be operating to drive the evolution at a higher rate than assumed in current models \\citep{Farmer10}. These astrophysical issues will in the future be important for the use of these systems for the verification of LISA as they can make the difference between detectability or not \\citep{Stroeer06}. In order to address this issue we have conducted a long-term timing study of the AM CVn star ES Ceti \\citep{Warner02}. This system shows a $620$s optical modulation which is complex in structure and varies on a night-by-night basis \\citep{Espaillat05}. Spectroscopic observations have confirmed that this is the orbital period, and also show the double-peaked helium emission lines which imply the presence of an accretion disc in this system (Steeghs, in prep). ES Cet is the ideal subject for a timing study, since it is the shortest period ultracompact binary population after V407 Vul and HM Cnc. It potentially connects these two systems to the rest of the AM CVn population: the period of ES Cet is only $51$s greater than V407 Vul, so ES Cet may also be in the mass transfer turn-on phase. If however ES Cet has already evolved on to the long-term and stable AM CVn path of lengthening period, then we would expect the period evolution to be the most rapid of all the systems, since gravitational radiation is strongest when masses are large and periods short. ", "conclusions": "\\label{sec:conclusions} We have obtained high time resolution lightcurves of the AM CVn star ES Cet with WHT+ULTRACAM and other smaller telescopes around the world. We observed the previously reported cyclical variability which has been confirmed to be orbital in origin. Our observations cover a baseline of $9$ years and we fit the lightcurves to obtain precise orbital timings over that period. We find that there is no evidence for a deviation from a linear ephemeris over this time period, although scatter in the timings is high due to flickering and large cycle-to-cycle variations in the light curve shape. There may be an underlying trend which is masked by this scatter. However, many predictions do suggest a lengthening orbital period which should be clearly detectable over this noise. Our findings suggest a low mass for the donor star, close or equal to the zero temperature mass of $0.062$\\Msun. This would make the double white dwarf formation channel the most likely scenario for this system. Alternatively, the accretion rate in this system may be significantly below the long-term, secular rate for an accreting AM CVn star. ES Cet is the shortest period AM CVn star and the only ultracompact binaries with shorter periods are HM Cnc and V407 Vul, both of which have been shown to have a decreasing period. ES Cet may be in an intermediate phase between these systems and the longer period AM CVns." }, "1101/1101.5431_arXiv.txt": { "abstract": "The 21cm forest -- HI absorption features in the spectra of high-redshift radio sources -- can potentially provide a unique probe of the largely neutral intergalactic medium (IGM) during the epoch of reionization. We present simulations of the 21cm forest due to the large scale structure of the reionization-era IGM, including a prescription for x-ray heating and the percolation of photoionization bubbles. We show that, if detected with future instruments such as the Square Kilometer Array (SKA), the 21cm forest can provide a significant constraint on the thermal history of the IGM. Detection will be aided by consideration of the sudden increase in signal variance at the onset of 21cm absorption. If radio foregrounds and the intrinsic source spectra are well understood, the flux decrement over wide bandwidths can also improve detection prospects. Our analysis accounts for the possibility of narrow absorption lines from intervening dense regions, but, unlike previous studies, our results do not depend on their properties. Assuming x-ray heating corresponding to a local stellar population, we estimate that a statistically significant detection of 21cm absorption could be made by SKA in less than a year of observing against a Cygnus A-type source at $z \\sim 9$, as opposed to nearly a decade for a significant detection of the detailed forest features. We discuss observational challenges due to uncertainties regarding the abundance of background sources and the strength of the 21cm absorption signal. ", "introduction": "\\label{s:intro} In both theoretical and observational astronomy today, a great deal of attention is being focused on improving our understanding of the epoch of reionization, when luminous sources began to ionize the intergalactic medium (IGM) for the first time after recombination. Understanding the process of reionization would lend insight into not only the nature of the first sources themselves, but also the evolution of the IGM, the build-up of large-scale structures, and the radiative transfer processes occurring at high redshifts. Currently, the most powerful data sets for constraining the ionization history of the IGM are the Gunn-Peterson absorption in the Lyman-$\\alpha$ forest \\citep{Bolton:2007} and the electron-scattering optical depth as observed in the cosmic microwave background (CMB) \\citep{Larson:2010}. The former places a lower limit on the fraction of hydrogen which is neutral (limited by the fact that the Gunn-Peterson trough saturates at a neutral fraction of $\\sim 10^{-5}$) and the latter measures the integrated level of ionization between the observer and the CMB. Thus we find a range of redshifts during which the universe went from substantially neutral to substantially ionized, but the detailed history of the reionization process is difficult to examine with these methods. The hyperfine transition of ground-state neutral hydrogen (with a rest-frame wavelength of 21cm) has the potential to offer much more detailed information about the epoch of reionization than either the Lyman-$\\alpha$ forest or the CMB optical depth measurement. Since 21cm radiation is the result of a low-energy transition at radio wavelengths, it can free-stream to the observer without being substantially absorbed in the intervening IGM. Also, most studies of the 21cm transition at high redshift focus on the signal in absorption or emission against the CMB, so the mean IGM itself rather than rare overdensities can be probed. There has been a great deal of recent progress in observation and instrumentation aimed at the search for the redshifted 21cm signal from the epoch of reionization. Several large-scale radio interferometers are currently being assembled to constrain the redshift of reionization as well as to produce all-sky brightness temperature measurements and a 21cm fluctuation power spectrum. These instruments include the Low Frequency Array (LOFAR)\\footnote{http://www.lofar.org/} in the Netherlands, the Giant Metrewave Radio Telescope (GMRT)\\footnote{http://gmrt.ncra.tifr.res.in/} in India, the Experiment to Detect the Global EoR Signal (EDGES)\\footnote{http://www.haystack.mit.edu/ast/arrays/Edges/}, the Precision Array to Probe the Epoch of Reionization (PAPER)\\footnote{\\citet{Parsons:2010}}, and the Murchison Widefield Array (MWA)\\footnote{http://www.MWAtelescope.org/} in Western Australia. Planning is also underway for the Square Kilometre Array (SKA)\\footnote{http://www.skatelescope.org/}. In this work, we discuss how studies of 21cm absorption against high-redshift radio sources can be used to study the evolution of the IGM around the epoch of reionization in what is known as the ``21cm forest'' \\citep{Carilli:2002,Furlanetto:2002}. In Section \\ref{s:od}, we discuss how observations of the 21cm forest have potential advantages over other probes. We describe in Section \\ref{s:numsims} our simulation of the expected signal using an IGM evolution code. In Section \\ref{s:fx}, we discuss our parameterization of the x-ray heating of the IGM, and in Section \\ref{s:sourcepop}, we discuss uncertainties regarding the population of high-redshift sources that could potentially be used as background sources for 21cm forest observations. Using a hypothetical radio-loud high-redshift source as an example, we present sample 21m forest spectra in Section \\ref{s:spectra}. In Section \\ref{s:stats}, we show how statistical detection methods can be used to isolate a faint signal in the case that radio-loud sources are rare at high redshift. We present our conclusions and outlook for future work in Section \\ref{s:conclusions}. ", "conclusions": "\\label{s:conclusions} We have shown that the magnitude and statistical significance of the 21cm forest signal is highly sensitive to assumptions about the thermal history of the IGM, and that the detectability of the signal also depends strongly on the redshift and radio-loudness of the background source. The strongest signal can be seen when x-ray heating is relatively inefficient (as parameterized by low values of the x-ray efficiency $f_X$, described in Sections \\ref{s:numsims} and \\ref{s:fx}) and when the source is very distant and radio-loud. For example, if the IGM is heated in a process for which $f_X=1$, a background radio source with flux density $S \\gtrsim 100$ mJy at redshift $z \\gtrsim 10$ would be needed for 21cm forest features to be significantly detected within a week of integration on the SKA. If x-ray heating is relatively inefficient, however, and $f_X \\approx 0.01$, the same signal-to-noise could be achieved with a source of 10 mJy at redshift $z \\approx 9$. In any case, it is clear that radio-loud sources at high redshift must be found for the 21cm forest signal to be observed. This conclusion is based on a relation [equation (\\ref{eq:Smin})] commonly used to quantify the detectability of the 21cm forest. However, we point out that the relation applies specifically to attempts to study the features of the 21cm forest at high signal-to-noise. In this study, our main interest is in using the 21cm forest to study the thermal history of the high-redshift IGM. We find this can be accomplished through the use of statistical methods that allow a high-significance detection of 21cm absorption against objects with lower flux densities or redshifts than those suggested by equation (\\ref{eq:Smin}). This suggests that the detection of 21cm absorption may be less challenging than previously thought. One of the statistical methods we discuss, which is applicable when strong absorption lines are not present, makes use of the increase in variance in the source spectrum at the frequency at which 21cm absorption begins. The passage of the radio signal alternately through regions of high ionization (HII regions produced by photoionization from bright UV sources) and through regions of relatively neutral gas (which can strongly absorb 21cm radiation) results in a pattern of high and low transmission and thus significantly increases the variance of the signal above the noise in certain cases. The measurement of this effect may allow us to detect 21cm absorption and place limits on the thermal history of the IGM even in cases where the detection of the forest features is not significant. A combination of long integration times and variance studies can therefore allow us to extract information from the 21cm forest even if the IGM is relatively hot and if bright high-redshift sources are rare. In the case that strong absorption lines are numerous, the lines dominate the spectrum's variance in the absorbed region, making the use of the variance method less straightforward. Although an increase in variance can signal that absorption is occurring \\citep{Carilli:2004}, the uncertainties in modelling the number and character of absorption lines make it difficult to directly link the level of variance increase with the thermal properties of the high-redshift IGM. A measurement of the average flux decrement in the absorbed region would similarly be strongly affected by the properties of the absorption lines. In Section \\ref{s:gauss}, we present a second method that can be used to extract thermal information about the IGM even if a spectrum contains many strong absorption lines. The method effectively removes absorption lines based on their narrowness and non-Gaussian nature, by fitting the source-subtracted flux distribution across a wide band to a Gaussian. We show that the flux decrement determined with this method can be detected at high significance even for observational scenarios far less favorable than those derived through a naive application of equation (\\ref{eq:Smin}). The estimates here should be considered a best-case scenario, due to the difficulty of accurately source-subtracting the spectrum in the presence of foregrounds. Equation (\\ref{eq:Smin}) suggests that an SKA-like observatory looking for absorption against a $z=9$ source with a radio flux similar to Cygnus A would require nearly a decade of integration time for a detection of the 21cm forest. By contrast, our method shows that a significant detection of the flux decrement -- and therefore a measurement of the optical depth and a constraint on $f_X$ -- could, in principle, be achieved with an integration time of less than 1 week for the same object and observatory. Even without perfect source subtraction, the width of the flux distribution over a wide band can be used to clearly distinguish amongst different values of $f_X$. Note that this does not depend strongly on properties of the absorption lines [in contrast to, e.g. \\citet{Xu:2011}]. This has implications for near-term reionization-era observational efforts. Our result shows that high-significance detections may be possible with instruments such as LOFAR or MWA, given a sufficiently radio-loud and distant source and a good understanding of foregrounds and source spectra. Looked at a different way, this method could make it possible for the 21cm forest to be detected against lower-redshift or dimmer sources with the SKA. In this work, we aimed to present the 21cm forest signal produced by the mean evolving IGM and the structure of large-scale photoionization along the line of sight. While this study presents an outline of the feasibility and major challenges of future 21cm forest measurements of the IGM, another potential extension of this work would be to use a more detailed simulation of the IGM properties and intervening absorption lines, perhaps from a large-scale numerical simulation including realistic dark matter and baryonic physics and radiative transfer, to more precisely simulate the expected signal in a high-redshift radio source. Planned future work also involves a more detailed treatment of the scope of this method in light of the capability of near-term observatories, and the possibility of fitting for the reionization model as part of the analysis. Another crucial element of any estimate of the utility of 21cm forest studies is an accurate estimate of the number of high-redshift radio-loud sources that will be available in the era of the SKA. We have given a brief overview in Section \\ref{s:sourcepop} of recent estimates, but updated theoretical studies and more observational follow-up of known radio-loud sources would both strongly impact the case for pursuing 21cm forest observations as a diagnostic tool for studying the reionization-era IGM." }, "1101/1101.0813_arXiv.txt": { "abstract": "The mass function $dN \\propto m^{-\\beta_0}dm$ of molecular clouds and clumps is shallower than the mass function $dN \\propto m^{-\\beta_\\star}dm$ of young star clusters, gas-embedded and gas-free alike, as their respective mass function indices are $\\beta_0 \\simeq 1.7$ and $\\beta_\\star \\simeq 2$. We demonstrate that such a difference can arise from different mass-radius relations for the embedded-clusters and the molecular clouds (clumps) hosting them. In particular, the formation of star clusters with a constant mean {\\it volume} density in the central regions of molecular clouds of constant mean {\\it surface} density steepens the mass function from clouds to embedded-clusters. This model is observationally supported since the mean surface density of molecular clouds is approximately constant, while there is a growing body of evidence, in both Galactic and extragalactic environments, that efficient star-formation requires a hydrogen molecule number density threshold of $n_{th} \\simeq 10^{4-5}\\,cm^{-3}$. Adopting power-law volume density profiles of index $p$ for spherically symmetric molecular clouds (clumps), we define two zones within each cloud (clump): a central cluster-forming region, actively forming stars by virtue of a local number density higher than $n_{th}$, and an outer envelope inert in terms of star formation. We map how much the slope of the cluster-forming region mass function differs from that of their host-clouds (clumps) as a function of their respective mass-radius relations and of the cloud (clump) density index. We find that for constant surface density clouds with density index $p \\simeq 1.9$, a cloud mass function of index $\\beta_0 = 1.7$ gives rise to a cluster-forming region mass function of index $\\beta \\simeq 2$. Our model equates with defining two distinct SFEs: a global mass-varying SFE averaged over the whole cloud (clump), and a local mass-independent SFE measured over the central cluster-forming region. While the global SFE relates the mass function of clouds to that of embedded-clusters, the local SFE rules cluster evolution after residual star-forming gas expulsion. That the cluster mass function slope does not change through early cluster evolution implies a mass-independent local SFE and, thus, the same mass function index for cluster-forming regions and embedded-clusters, that is, $\\beta = \\beta_\\star$. Our model can therefore reproduce the observed cluster mass function index $\\beta_\\star \\simeq 2$. For the same model parameters, the radius distribution also steepens from clouds (clumps) to embedded-clusters, which contributes to explaining observed cluster radius distributions. ", "introduction": "\\label{sec:intro} The star formation efficiency (SFE) achieved by star cluster gaseous precursors at the onset of residual star-forming gas expulsion is a crucial quantity since it influences the cluster dynamical response to gas expulsion significantly \\citep[the so-called violent relaxation;][]{hil80,gey01,bau07,pro09}. Specifically, the SFE is tightly related to whether the cluster survives violent relaxation and, if it survives, what mass fraction of its stars it retains. The SFE being the ratio between the stellar mass of embedded-clusters and the initial gas mass of their precursor molecular clouds, the comparison of the mass functions of young star clusters and of molecular clouds holds the potential of highlighting whether the SFE varies with molecular cloud mass. The mass function $dN \\propto m^{-\\beta_0}dm$ of giant molecular clouds (GMCs) in the Local Group of galaxies has an index $\\beta_0 \\simeq 1.6$-$1.7$ \\citep{ros05, bli06} \\citep[see also ][for the case of the GMC mass function in the Magellanic Clouds]{fuk08}. These GMCs, when compressed by the high pressure of violent star-forming environments, are expected to be the parent clouds of massive star clusters forming profusely in galaxy mergers and starbursts \\citep{jog92,jog96}. The same slope $\\beta_0 \\simeq 1.6$-$1.7$ is also found for the mass function of density enhancements contained by GMCs -- referred to as molecular clumps \\citep{lad91,kra98,won08}. In quiescent disc galaxies such as the Milky Way, those are observed to be the progenitors of open clusters \\citep[][and references therein]{hp94,ll03}. In contrast to molecular clouds and clumps, the mass function $dN \\propto m^{-\\beta_{\\star}}dm$ of embedded and young clusters is, in most cases, reported to have an index $\\beta_{\\star} \\simeq 2$ \\citep[e.g.][]{zf99,bik03,ll03,oey04}, which is steeper than the mass function of molecular structures. Given the uncertainties affecting both slopes, the significance of the $\\beta_{\\star} - \\beta_0$ difference remains uncertain. \\citet{elm96} suggest that error-free measurements of GMC masses may bring the mass function slopes of young star clusters and GMCs in agreement. Conversely, one can consider that the slope difference is significant, which is the approach we adopt in this paper. The question we set to answer is: what process of the physics of cluster-formation steepens the power-law mass function of molecular clouds and clumps from $\\beta_0=1.7$ to $\\beta_{\\star}=2$? The $\\beta_{\\star}-\\beta_0$ difference suggests that the SFE is a decreasing function of the cloud (clump) mass. Besides sounding counter-intuitive, a mass-varying SFE is necessarily conducive to mass-dependent cluster infant weight-loss since the fraction of stars remaining bound to clusters through violent relaxation is a sensitive function of the SFE \\citep[e.g. fig.~1 in][]{par07}. This does not seem to be supported by observations of young star clusters, as their mass function slope is reported to remain invariant with time over their first 100\\,Myr of evolution \\citep {ken89,mck97,ll03,zf99,oey04,dow08,cha10}. However, this contradiction is apparent only for it is worth keeping in mind that the SFE driving cluster violent relaxation is the mass fraction of gas turned into stars {\\it over the volume of gas forming stars}. And this volume of star-forming gas may not coincide with the entire cloud (clump). In what follows, we refer to it as the {\\it cluster-forming region (CFRg)}. Its SFE is the {\\it local} SFE and its mass function slope is $-\\beta$. The invariance of the young cluster mass function slope at early time suggested by many observations demands a mass-independent {\\it local} SFE. That is, the mass fraction of gas turned into stars by a CFRg is independent of its mass. This in turn implies that the slopes of the CFRg and embedded-cluster mass functions are identical: $\\beta = \\beta_{\\star}$. Therefore, understanding the slope difference $\\beta_\\star - \\beta_0$ between the star cluster and molecular cloud (clump) mass functions equates with understanding why the CFRg mass function is steeper than the mass function of their host clouds (clumps), i.e. $\\beta = \\beta_\\star \\simeq 2$ and $\\beta_0 \\simeq 1.7$. The $\\beta - \\beta_0$ difference suggests that the mass fraction of star-forming gas inside molecular clouds (clumps) is a decreasing function of the cloud (clump) mass. Besides, that the CFRg represents a fraction only of its host cloud (clump) allows us to define a {\\it global} SFE, namely, the ratio between the mass in stars formed inside a molecular cloud (clump) and its initial gas mass. The {\\it global} SFE is relevant to explaining the $\\beta_{\\star}-\\beta_0$ slope difference, but irrelevant for modelling cluster violent relaxation. What could be the origin of a mass-varying mass fraction of star-forming gas inside molecular clouds (clumps)? In other words, why should the {\\it global} SFE vary with the cloud (clump) mass such that $\\beta_\\star \\neq \\beta_0$? The mean {\\it surface} density of GMCs in our Galaxy is about constant \\citep[fig.~8 in][]{bli06}. This result is reminiscent of Larson's seminal study \\citep{lar81} showing that molecular clouds have approximately constant mean column densities \\citep[see also][]{lom10}. On the other hand, star-forming regions are observed to be systematically associated with dense molecular gas, namely, with {\\it number} densities of at least $n_{H2} \\simeq 10^4$-$10^5\\,cm^{-3}$ \\citep[][]{mue02,gao04,fau04,fon05,shi03,wu10}. See also fig.~1 and section 3 in \\citet{par10} for a discussion. Several studies have therefore suggested that star formation requires a gas {\\it volume} (or {\\it number}) density threshold \\citep[e.g.][]{eva08,wu10,lad10}. In this contribution, we develop a model for a spherically symmetric molecular cloud (clump) with a power-law density profile which forms a star cluster in its central region. We demonstrate that if the mass-radius relation of CFRgs differs from that of the host-clouds (clumps), then the slopes of their respective mass functions are different too (i.e. $\\beta \\neq \\beta_0$). This will be the case for molecular clouds of constant mean surface density hosting CFRgs of constant mean volume density. Applying the same model to the radius distribution, we will show that it can also contribute to explaining why the distribution of star cluster half-light radii is significantly steeper than the distribution of GMC sizes. Figure \\ref{fig:skMF} summarises the different mass functions encompassed through the paper, along with their respective index and the various mass ratios relating them. Note that this paper does not intend to explain the slope of the {\\it stellar} initial mass function. That issue is addressed in \\citet{sha10} whose model successfully reproduces the Salpeter slope of $-2.35$ for a population of pre-stellar cores exceeding a volume density threshold in a fractal cloud. \\begin{figure} \\includegraphics[width=\\linewidth]{sketch_MF.eps} \\caption{Illustration of the different mass functions tackled through the paper. From right to left: the clump (cloud) mass function $dN \\propto m_{clump}^{-\\beta_0}\\,dm_{clump}$, the CFRg mass function $dN \\propto m_{th}^{-\\beta}\\,dm_{th}$, the embedded-cluster mass function $dN \\propto m_{ecl}^{-\\beta_{\\star}}\\,dm_{ecl}$. For the sake of completeness, the cluster mass function at the end of violent relaxation is also shown as the leftmost straigthline. The horizontal arrows depict mass ratios relating pairs of mass functions. From right to left: the mass fraction $m_{th}/m_{clump}$ relates the CFRg and clump mass functions, the global SFE $\\epsilon_{global} = m_{ecl}/m_{clump}$ quantifies the embedded-cluster stellar mass contained within molecular clouds (clumps), the local SFE $\\epsilon_{loc} = m_{ecl}/m_{th}$ is the CFRg mass fraction turned into stars, the bound fraction $F_{bound}$ quantifies the mass fraction of stars which stays bound to a cluster when violent relaxation is over (i.e. after infant weight-loss). Vertical scaling is arbitrary. \\label{fig:skMF} } \\end{figure} The outline of the paper is as follows. Section \\ref{sec:evid} summarises two types of evidence supporting the hypothesis of a constant mean {\\it volume} density for CFRgs. One hinges on the early dynamical evolution of star clusters. The second is based on the mapping of star-forming regions with dense molecular gas tracers. In Section \\ref{sec:cc}, we build a model relating the power-law mass function of CFRgs to the power-law mass function of their host clouds (clumps). We map how the slope difference $\\beta - \\beta_0$ varies as a function of the mass-radius relation and density profile of molecular clouds (clumps). In Section \\ref{sec:implic}, we discuss the implications of our model. Specifically, we focus on the physically-motivated case of virialized pressure-bounded (i.e.~constant mean surface density) clouds (clumps). Section \\ref{sec:rdist} is the counterpart of Section \\ref{sec:cc} as it models the radius distribution of CFRgs in relation to that of their parent clouds (clumps). Our conclusions are presented in Section \\ref{sec:conclu}. ", "conclusions": "\\label{sec:conclu} It has long been recognized that the mass function $dN \\propto m^{-\\beta_0} dm$ of GMCs and of their molecular clumps mapped in CO-emission line is shallower than the `canonical' young cluster mass function $dN \\propto m^{-\\beta_\\star} dm$, i.e. $\\beta_0 \\simeq 1.7$ and $\\beta_\\star \\simeq 2$. This slope difference is puzzling since it seemingly implies an SFE varying with the GMC or clump mass, hence mass-dependent cluster infant weight-loss while the cluster responds to gas-expulsion. This is in contradiction with most young cluster mass function data gathered so far. In this contribution we bring an original solution to this problem by assuming that star formation requires a {\\it number} density threshold $n_{th} \\simeq 10^{4-5}\\,cm^{-3}$, equivalent to a {\\it volume} density threshold $\\rho_{th} \\simeq 700-7000\\,M_{\\sun}.pc^{-3}$. This hypothesis is supported by the tight association observed between star-formation and dense molecular gas (as evidenced by e.g. $H^{13}CO^+$ and $HCN$ tracers; see Section \\ref{sec:evid}). Our model builds on a spherically symmetric cloud (or clump) with a power-law density profile and forming a star cluster in its central region. The density threshold for star formation $\\rho_{th}$ is not necessarily achieved through the whole molecular cloud (clump), thereby implying that the mass and radius of the CFRg differ from those of the cloud (or clump) containing it. We refer to $\\beta$ as the mass function index of the spatially-limited CFRg where $n_{H2} \\geq n_{th}$ (see Fig.~\\ref{fig:sketch}). In that context, star formation can be quantified by two distinct efficiencies {\\it of different physical significances}. We refer to the {\\it global} SFE as the ratio between the embedded-cluster stellar mass at the onset of gas-expulsion and the initial gas mass of the {\\it clump (cloud)} hosting it. As such, the global SFE is relevant to understand the difference between the cloud (clump) mass function on the one hand, and the embedded-cluster mass function on the other hand. In contrast, the {\\it local} SFE quantifies the ratio between the embedded-cluster stellar mass and the initial gas mass of the CFRg, i.e. the gas mass with $n_{H2} \\geq n_{th}$. This is the local SFE -- {\\it not} the global one -- which is relevant to understand why cluster violent relaxation is mass-independent. Mass-independent infant weight-loss demonstrates that the local SFE is CFRg-mass-{\\it in}dependent hence that the slopes of the CFRg and embedded-cluster mass functions are identical ($\\beta = \\beta_\\star$). This does not prevent the global SFE from being {\\it clump/cloud} mass-dependent, as suggested by the difference in slope $\\beta_\\star - \\beta_0$ between the cloud (clump) and cluster mass functions. To adopt a volume density threshold for cluster formation immediately implies that CFRgs have a constant mean volume density (Eq.~\\ref{eq:av_rhoth}). Based on the conditions required for the tidal-field impact upon clusters responding to gas-expulsion to be mass-independent, this is also the conclusion reached by \\citet{par10}. Actually, not only does mass-independent violent relaxation demand mass-independent local SFE, it also requires mass-independent gas-expulsion time-scale \\citep{par08b}, and mass-independent tidal-field impact \\citep{par10}. We have shown that the difference in slope between the clump- (cloud-) and CFRg-mass functions is a sensitive function of the mass-radius relation and density index $p$ of clumps (clouds). Constant radius clumps result in the mass function of CFRgs (hence of embedded-clusters) being shallower than the mass function of their host clumps (clouds). This is due to more massive clumps being denser, thus containing a greater fraction of their mass above the number density threshold $n_{th}$. Equivalently, the global SFE increases with the clump mass. Conversely, the volume density of constant surface density clumps is a decreasing function of their mass, and so is their mass fraction of star-forming gas. This renders the mass function of CFRgs steeper than that of clumps/clouds ($\\beta > \\beta_0$). For constant volume density clumps/clouds, CFRg and cloud/clump mass function slopes are alike (Fig.~\\ref{fig:beta}). Given a cloud (clump) mass-radius relation, the slope difference $|\\beta-\\beta_0|$ depends on the density index $p$ of clumps (clouds): the shallower the clump/cloud density profile, the larger $|\\beta-\\beta_0|$ (Fig.~\\ref{fig:betap}). The steepening of the mass function $\\beta_0 \\simeq 1.7$ of molecular clumps and GMCs into that $\\beta_{\\star} \\simeq 2$ of young star clusters therefore requires molecular clouds and clumps to have a constant surface density (Fig.~\\ref{fig:MF}). This property is actually well-established for GMCs in the Milky Way \\citep{bli06, hey09}. Whether it also stands for molecular clumps, namely, the density enhancements -- birth sites of open clusters -- observed locally within Galactic GMCs, is less certain. Rather, molecular clumps show a constant volume density corresponding to that required to excite the molecular transition of relevance (middle panel of Fig.~\\ref{fig:isomth}). Based on $C^{18}O$ data, we speculate that the mass-radius relation of {\\it cluster-forming} molecular clumps is one of constant surface density, rather than of constant volume density (see Fig.~\\ref{fig:mcSF}). The transition from a narrow range in volume densities for all $C^{18}O$ clumps to a narrow range in surface densities for those with signs of star formation stems from excluding the lowest mass clumps. Those clumps contain a tiny mass only with $n_{H2} \\geq n_{th}$, which explains their failure at displaying signs of star formation (top panel of Fig.~\\ref{fig:isomth}). From their survey in dust-continuum emission of star-forming regions in the Galactic disc, \\citet{mue02} infer a mean density index $p \\simeq 1.8$. Interestingly, in that case, we find that the mass function slope $\\beta_0 \\simeq 1.7$ of clouds (clumps) steepens into a CFRg mass function slope $\\beta \\simeq 2$ (Fig.~\\ref{fig:betap}) hence $\\beta_{\\star} \\simeq 2$, in agreement with what is suggested by observations. Equivalently, the global SFE of molecular clouds (clumps) is a decreasing function of their mass (Fig.~\\ref{fig:sfe}). A natural outcome of our model is that as mapping of molecular clumps move inwards to their higher-density CFRgs, the inferred mass function is expected to steepen and to near $\\beta \\simeq 2$. This may be the reason why \\citet{shi03} find $\\beta \\simeq 1.9$ for CFRgs mapped in $CS (J~5 - 4)$ which, as they quote, is steeper than what is measured with tracers of lower density gas, and closer to the mass spectral index of OB associations (their fig.~20). In addition to the mass functions, we have also studied the radius distributions. Given constant mean surface density clouds (clumps), not only does our model steepen the mass function, it also steepens the radius distribution (top panel of Fig.~\\ref{fig:rdist}). The slope of the radius distribution of clouds (clumps) is determined by their mass-radius relation and mass function slopes (Eq.~\\ref{eq:debex}). Constant surface density clouds (clumps) ($\\delta=1/2$) with $\\beta_0 = 1.7$ have a radius distribution $dN \\propto r^{-x_0} dr_{clump}$ of index $x_0=2.4$. For a density index $p=1.8$, the slope of the radius distribution steepens to $x \\simeq 4$ for CFRgs hence embedded-clusters (top panel of Fig.~\\ref{fig:x}). Our model thus helps explain why the cluster radius distribution is significantly steeper than the size distribution of GMCs \\citep{sss85,sch07}." }, "1101/1101.2211_arXiv.txt": { "abstract": "{The chromosphere of the Sun is a temporally and spatially very varying medium for which the assumption of ionisation equilibrium is questionable. } {Our aim is to determine the dominant processes and timescales for the ionisation equilibrium of calcium under solar chromospheric conditions. } {The study is based on numerical simulations with the RADYN code, which combines hydrodynamics with a detailed solution of the radiative transfer equation. The calculations include a detailed non-equilibrium treatment of hydrogen, calcium, and helium. Next to an hour long simulation sequence, additional simulations are produced, for which the stratification is slightly perturbed so that a ionisation relaxation timescale can be determined. The simulations are characterised by upwards propagating shock waves, which cause strong temperature fluctuations and variations of the (non-equilibrium) ionisation degree of calcium. }{The passage of a hot shock front leads to a strong net ionisation of Ca~II, rapidly followed by net recombination. The relaxation timescale of the calcium ionisation state is found to be of the order of a few seconds at the top of the photosphere and 10 to 30\\,s in the upper chromosphere. At heights around 1\\,Mm, we find typical values around 60\\,s and in extreme cases up to $\\sim 150$\\,s. Generally, the timescales are significantly reduced in the wakes of ubiquitous hot shock fronts. The timescales can be reliably determined from a simple analysis of the eigenvalues of the transition rate matrix. The timescales are dominated by the radiative recombination from Ca\\,III into the metastable Ca\\,II energy levels of the 4d\\,$^2$D term. These transitions depend strongly on the density of free electrons and therefore on the (non-equilibrium) ionisation degree of hydrogen, which is the main electron donor. }{ The ionisation/recombination timescales derived here are too long for the assumption of an instantaneous ionisation equilibrium to be valid and, on the other hand, are not long enough to warrant an assumption of a constant ionisation fraction. Fortunately, the ionisation degree of \\ion{Ca}{II} remains small in the height range, where the cores of the H, K, and the infrared triplet lines are formed. We conclude that the difference due to a detailed treatment of Ca~ionisation has only negligible impact on the modelling of spectral lines of \\ion{Ca}{II} and the plasma properties under the conditions in the quiet solar chromosphere. } ", "introduction": "\\label{sec:intro} Modern observations unambiguously show that the chromosphere of the Sun is very dynamic and inhomogeneous \\citep[see the reviews by, e.g.,][]{2001ASPC..223..131S, 2006ASPC..354..259J, 2007ASPC..368...27R, 2009SSRv..144..317W}. The interpretation of chromospheric observations, which often involves the construction of numerical models, thus must take into account spatial and temporal variations. Unfortunately, many simplifying equilibrium assumptions that can be made for the low photosphere, are not applicable for the layers above. A realistic chromosphere model must therefore account for deviations from the equilibrium state in the context of an atmosphere that changes on short timescales and small spatial scales. A successful example is the explanation of the formation of ``calcium grains'' by \\citet{1997ApJ...481..500C}. Other examples of simulations with a time-dependent non-equilibrium treatment include the ionisation of hydrogen (\\citeauthor{2002ApJ...572..626C} \\citeyear{2002ApJ...572..626C}, hereafter referred to as Paper~I; \\citeauthor{2006A&A...460..301L} \\citeyear{2006A&A...460..301L}; \\citeauthor{2007A&A...473..625L} \\citeyear{2007A&A...473..625L}) and the concentration of carbon monoxide \\citep{asensio03, 2005A&A...438.1043W}. When an individual process, be it a chemical reaction or atomic transition, changes on finite timescales longer than the dynamic timescales of the atmosphere, then it becomes necessary to solve a system of rate equations instead of the much simpler statistical equilibrium equations. In paper~I, it was shown that the resulting variation of the hydrogen ionisation degree is significantly reduced due to slow recombination rates. The ionisation degree then depends not only on the local gas temperature, density, and the radiation field but also on the history of these properties. Already \\citet{1991A&A...250..212R} considered the non-equilibrium ionisation of magnesium (\\element{Mg}) and calcium (\\element{Ca}) in their ab-initio solar chromosphere models. Their numerical implementation included a non-local thermodynamic equilibrium (NLTE) approximation with very simplified two-level atoms for \\ion{Mg}{II}-{III} and \\ion{Ca}{II}-{III}. In this paper, we investigate the \\ion{Ca}{II}-{III} ionisation balance in detail. The study is based on RADYN simulations similar to the hydrogen case in Paper~I. The numerical simulations are described in Sect.~\\ref{sec:sim}. The results of this study are presented in Sect.~\\ref{sec:res}, followed by a discussion and conclusions in Sect.~\\ref{sec:disc}. ", "conclusions": "\\label{sec:disc} The ionisation/recombination timescales, which are here derived from 1-D RADYN simulations, are too long for the assumption of an instantaneous ionisation equilibrium to be valid. On the other hand, the timescales are not long enough to warrant an assumption of a constant ionisation fraction. We find noticeable deviations from the ionisation equilibrium in the middle model chromosphere but in general the \\ion{Ca}{II}-{III} ionisation fraction remains small. The error due to the often made simplifying assumption of statistical equilibrium is therefore negligible for most applications. The effect is barely visible in the synthesized intensity for the diagnostically important spectral lines of \\ion{Ca}{II}, i.e., the H and K lines and the infrared triplet. This finding is illustrated for the \\ion{Ca}{II}\\,K line in Fig.~\\ref{fig:intensity}. The differences in the emergent intensity between the statistical equilibrium approach ($I_\\mathrm{K, SE}$) and the time-dependent non-equilibrium simulation ($I_\\mathrm{K, TD}$) are very small and hardly discernible in the temporal evolution of the intensity at a fixed wavelength position in the line core (top row) and even smaller (order of 0.1\\,\\%) in the average spectrum (bottom row). The relative difference $(I_\\mathrm{K, TD} - I_\\mathrm{K, SE})/I_\\mathrm{K, SE}$ reveals peaks with values of up to a few percent for wavelengths close to the line core. This is true both at disk-centre (right column) and close to the limb (left column). The differences are even smaller away from the line core and the emission peaks. Noticeable deviations occur only in connection with shock fronts and thus sharp intensity increases. There, the effect is caused by a small temporal offset in the evolution of the level populations in the TD case with respect to the SE case, like it is seen for the ionisation fraction in Fig.~\\ref{fig:ionvstime}. It is safe to conclude that the time-dependent non-equilibrium treatment of the \\ion{Ca}{II}-{III} ionisation appears to be of minor importance for the lower atmosphere in quiet Sun regions. One restriction of the simulations is the use of only one spatial dimension, which could cause the shock wave profiles to be too extreme with potentially too high peak temperatures. In 3-D, the shock fronts are expected to be weaker owing to the larger number of degrees of freedom. The consequences of the deviations from the \\ion{Ca}{II}-{III} ionisation equilibrium would be even less important in that case. Furthermore, we neglect the effect of incident radiation from the corona on the photoionisation of Ca. Most important in that respect would be the Lyman alpha line at $\\lambda = 121.5$\\,nm. However, the Lyman $\\alpha$ photons have too little energy to ionise \\ion{Ca}{II} from the ground state (11.88\\,eV or $\\lambda = 104.4$\\,nm). The photoionisation edges for levels 2 and 3, on the other hand, both lie close to the Lyman alpha line. The incident radiation certainly matters most at large heights close to the transition region. There, however, photoionisation from the ground level dominates, which is obviously not influenced by Lyman alpha photons. We conclude that the RADYN simulations employed here are sufficiently realistic for an evaluation of the importance of non-equilibrium effects for the ionisation of calcium in the context of a strongly varying chromosphere in quiet Sun regions." }, "1101/1101.1354_arXiv.txt": { "abstract": "We present the results of astrometric observations with VERA toward the H$_2$O maser sources in IRAS 05137+3919, which is thought to be located in the far outer Galaxy. We have derived the parallax of $\\pi = 0.086\\pm 0.027$ mas, which corresponds to the source distance of $D=11.6^{+5.3}_{-2.8}$ kpc. Although the parallax measurement is only 3-$\\sigma$ level and thus the distance uncertainty is considerably large, we can strongly constrain the minimum distance to this source, locating the source at the distance from the Sun greater than 8.3 kpc (or 16.7 kpc from the Galaxy's center) at 90\\% confidence level. Our results provide an astrometric confirmation that this source is located in the far outer Galaxy beyond 15 kpc from the Galaxy center, indicating that IRAS 05137+3919 is one of the most distant star-forming regions from the Galaxy center. ", "introduction": "Star-forming regions located in the far outer region of the Galaxy are interesting targets for astronomical studies in terms of both the Galactic structure and star formations. For examples, star-forming regions in the extremely outer Galaxy can be used to trace the extent of the Galaxy's stellar disk as well as the spiral structure in the outer regions. Also, star-forming regions in the far outer Galaxy provide unique laboratories to investigate how stars form in an extreme environment: in the far outer Galaxy, the metallicity is much lower than that in the Solar neighborhood (e.g., Smartt \\& Rolleston 1997; Rudolph et al. 2006), and hence the star-formation in the outer Galaxy at present could represent the star formation process in the early phase of galaxy evolutions, which are difficult to investigate through direct observations. Currently several star-forming regions are expected to be located in the far outer region (here we define the far outer Galaxy as the region with the galacto-centric radius $R_{\\rm GC}$ greater than 15 kpc). These star-forming regions have been mainly discovered based on extensive CO surveys toward the outer Galaxy (e.g., Wouterloot \\& Brand 1989; Digel et al. 1994), and the distances were estimated based on the kinematic distances obtained from observed radial velocities and assumed rotation curves. These surveys provided handful candidates of star-forming regions in the far outer regions, and in fact, some of them are intensively observed to study the star formation process in an extreme environment (e.g., Ruffle et al. 2007; Kobayashi et al. 2008). However, since the distance estimates were based on the kinematic distances, there exist large distance uncertainties and hence the locations of such sources are yet to be confirmed with better accuracy. For instance, the star-forming region W3(OH), located in the Perseus arm, (though this is not the source in the far outer Galaxy) had a large discrepancy between the photometric distance and the kinematic distance, and accurate astrometries with phase-referencing VLBI (Xu et al. 2006; Hachisuka et al. 2006) revealed that the kinematic distance was an overestimation by a factor of two. The case for W3(OH) clearly demonstrated that the distance estimates based only on the kinematic distances are inadequate to conclude the locations of the star forming regions in the Galaxy. Fortunately, phase-referencing VLBI astrometry (such as using VERA and VLBA) are now powerful enough to determine accurate distances even beyond 5 kpc (e.g., Honma et al. 2007; Reid et al. 2009a), and hence astrometric measurements of star forming regions with VLBI will have great impact on the research of star-formation in the far outer Galaxy. IRAS 05137+3919 is one of such star-forming regions which is thought to be located in the far outer Galaxy. The CO line was detected toward this source by Wouterloot \\& Brand (1989), catalogued as WB 621. The systemic velocity of CO was found to be $-25.9$ km s$^{-1}$, which provided a kinematic distance of 12 kpc. HCO$^+$ line is associated with this source, with a systemic velocity of $-26$ km s$^{-1}$ (Molinari et al. 2002). In HCO$^+$ line profile, a wing with a full width of 80 km s$^{-1}$ was also detected, indicating molecular outflows from the forming star(s). Also associated with this source are infrared sources, dust continuum emission, compact HII regions, H$_2$CO absorption, OH, CH$_3$OH and H$_2$O masers (e.g., Brand et al. 1994; Molinari et al. 2002; 2008; Edris et al 2007; Araya et al. 2007; Sunada et al. 2007; Xu et al. 2008; Faustini 2009). Molinari et al.(2008) conducted a spectrum fitting to IRAS 05137+3919 from mm/sub-mm wave continuum taken with SIMBA at SEST and SCUBA at JCMT to infrared obtained with IRAS and MSX, and reported that the exciting source is likely to be an O8 ZAMS star with $L=2.5\\times 10^5 L_\\odot$. Thus, if the kinematic distance is correct, IRAS 05137+3919 is a massive star-forming region located at $R_{\\rm GC}\\sim 20$ kpc, which makes this source most interesting target to study the star formation in the extreme environment. However, since the source is located in the anti-Galactic center region, the kinematic distance, which is solely based on the radial velocity, would be highly uncertain. In order to measure the distance of IRAS 05137+3919 by means of trigonometric parallax, we have conducted the astrometry of H$_2$O maser and here report the results. ", "conclusions": "\\subsection{Distance comparisons and location in the Galaxy} Previously the distance to IRAS 05137+3919 was estimated based on the kinematic distance. For instance, Wouterloot \\& Brand (1989) estimated the distance of $D=12$ kpc assuming nearly flat rotation curve with $R_0=8.5$ kpc and $\\Theta_0=220$ km s$^{-1}$, and a similar value was adopted in recent studies (e.g, Molinari et al. 2008, $D=11.5$ kpc). On the other hand, our astrometric measurements provide a source distance of $D=11.6^{+5.3}_{-2.8}$ kpc, and the minimum distance $D_{\\rm min}$ of 8.3 kpc at 90\\%-confidence level. Although the uncertainty of our parallax distance is considerably large, there is no indication of systematic difference between the kinematic distances and the astromteric distance for this source, implying that the kinematic distances previously estimated were fairly reasonable. This is in contrast to distance over-estimations using kinematic distance, found in such as W3(OH) in Perseus arm (e.g., Xu et al. 2006). In the galacto-centric coordinate, the $1-\\sigma$ distance of $D=11.6^{+5.3}_{-2.8}$ kpc from the Sun corresponds to $R_{\\rm GC}=20.0_{-2.8}^{+5.3}$ kpc, and the minimum distance $D_{\\rm min}$ of 8.3 kpc corresponds to the minimum Galacto-centric radius of 16.7 kpc (in both cases the IAU standard $R_0=8.5$ kpc is adopted). Therefore, the star-forming region IRAS 05137+3919 is most likely to be located beyond 15 kpc from the Galaxy and thus is in the far outer Galaxy. Hence, our results provide the first astrometric confirmation that there exist star-forming regions in the far outer regions of the Galaxy, demonstrating the existence of star formation activities in such region. Currently the farthest spiral arm confirmed by astrometry is so-called Outer Arm. Recently VLBI astrometries of star forming regions in the Outer Arm have been carried out (e.g., Honma et al. 2007; Hachisuka et al. 2009), and the distance from the Sun to the Outer Arm in the direction of the Galactic anti-center is found to be 5 -- 6 kpc (or $\\sim$14 -- 15 kpc in Galacto-centric radius). We note that IRAS 05137+3919's location in the Galaxy is far beyond the Outer arm, and hence the location of this star-forming region in the Galaxy impose a question on how molecular clouds are formed in the far outer regions: whether there exists another (rather faint) spiral arm beyond the Outer arm, or there exist another mechanism to form a molecular cloud in such an extreme region. \\subsection{Maser structure} Figure 2 shows the maser spot distributions in IRAS 05137+3919 with respect to the tracking center position. The maser spots are basically located in two maser features, one in the north-east (feature 1) and the other in the south-west (feature 2), with a separation of $\\sim$140 mas. In figure 2, maser proper motions with respect to the position reference source are also plotted (black arrows). Note that these proper motions include Galactic rotation of IRAS 05137+3919 itself as well as its systematic deviation from Galactic rotation (i.e., non-circular motion). To remove such effects and to see the internal motions of maser spots, we estimated the systematic motion as follows: first we average the maser motions in each feature (5 spots in the north-east feature and 1 spot in the south-west feature), which were found to be ($\\bar{\\mu}_X$, $\\bar{\\mu}_Y$)= (0.571, $-0.165$) mas yr$^{-1}$ and ($0.022$, $-1.619$) mas yr$^{-1}$, respectively. Then we obtained the systematic proper motion of IRAS 05137+3919 by taking the mean of the averaged motions of the two maser features, yielding ($\\mu_{X,0}$, $\\mu_{Y,0}$)= (0.297, $-0.892$) mas yr$^{-1}$. The grey vectors in figure 2 are the internal maser motions obtained by subtracting the systematic proper motion. Interestingly, the directions of estimated internal motions are fairly close to the direction of elongation of the two maser features, and the internal maser motions shows that the separation of the two maser features are increasing. These results are consistent with a picture in which a bipolar outflow/jet from an exciting source (a proto-star in IRAS 05137+3919) forms two shock regions, where the maser emissions are observed. The amplitude of the internal proper motions obtained above is about 0.78 mas yr$^{-1}$, corresponding to 43 km s$^{-1}$ using the distance of 11.6 kpc obtained in the present paper. This is significantly larger than the width of the radial velocity of H$_2$O maser, which is only $\\sim$2 km s$^{-1}$. However, we note that the observations of HCO$^+$ by Molinari et al.(2002) reported a velocity wing with a width of $\\sim$80 km s$^{-1}$ (most-likely due to a large-scale outflow), which is fairly comparable with internal maser motions obtained here. Hence, the internal proper motion of 43 km s$^{-1}$ is not unlikely for this source. If the H$_2$O maser in IRAS 05137+3919 indeed traces the outflow/jet motion, then the orientation of outflow/jet axis is nearly perpendicular to the line of sight. We note that the large-scale outflow traced with HCO$^+$ has a different orientation with the small-scale outflow traced with H$_2$O maser emissions, because the high velocity wing of HCO$^+$ is seen in the radial velocity profile, indicating that the large scale outflow cannot be perpendicular to the line of sight, in contrast to the maser outflow. This may imply that the outflow orientation changes in different scales (due to, e.g., interaction with ambient matters and/or precession of outflow axis), or that there exist two different outflows from two independent exciting sources. Note that the maser spot identification in the present studies was strictly done to ensure high precision in astrometry and hence some maser spots were removed from our analyses based on the four criteria described in section 3. In case that we loosen the criteria 3 and 4 in section 3 to include more spots, we found a few additional spots in the two features shown in figure 2, and also found one spot which is about 300 mas east of the two features shown in figure 2, which appears to be unassociated with the features in figure 2. Therefore, even if the selection criteria are modified, the basic maser structure presented in this chapter is not drastically changed. \\subsection{Galactic Rotation in the far-outer Galaxy} Using the systematic motion of IRAS 05137+3919 obtained in section 5.2, we can constrain the Galactic rotation velocity in the far outer region. As we described above, by taking the mean of maser spot motions, we estimated the systematic motion of the source as ($\\mu_{X,0}$, $\\mu_{Y,0}$)=(0.297, $-0.892$) mas yr$^{-1}$. By correcting for the standard solar motion of IAU 1985 with ($U_\\odot$, $V_\\odot$, $W_\\odot$)=(10.0, 15.4, 7.8) km s$^{-1}$ (Kerr \\& Lynden-Bell 1986), the motions of IRAS 05137+3919 with respect to the Local Standard of Rest are converted to be ($\\mu_l$, $\\mu_b$)=(0.588, $-0.130$) mas yr$^{-1}$ in the Galactic coordinate. The proper motion along with the Galactic plane ($\\mu_l$) is considerably large, since the source is located toward the anti-center region ($l$=168.1$^\\circ$) and hence the proper motions of the source and the LSR should mostly cancel out. For instance, if we assume a flat rotation curve with $\\Theta_0=220$ km s$^{-1}$, the expected relative proper motion of IRAS 05137+3919 with respect to the LSR along Galactic plane is $\\mu_l=-0.071$ mas yr$^{-1}$ (at $D=11.6$ kpc). Therefore, our results suggest that the Galactic rotation of IRAS 05137+3919 is smaller than that of a flat rotation curve by $\\Delta v=$ 36 km s$^{-1}$, which is obtained by using $D=11.6$ kpc and the proper motion difference of $0.589-(-0.071)=0.660$ mas yr$^{-1}$ (note that this result is not significantly changed with different values of $\\Theta_0$). Reid et al.(2009b) suggested that the massive star-forming region could rotate around the Galaxy slower than the Galactic rotation velocity by $\\sim 15$ km s$^{-1}$. Our results for IRAS 05137+3919, rotation speed slower by 36 km s$^{-1}$, may be partly explained such a slow rotation of star forming regions. Also, the value of $\\Delta v$ is dependent of the source distance $D$, and if a smaller distance is adopted, the deviation from the flat rotation curve becomes small. For instance, if $D_{\\rm min}$ of 8.3 kpc is adopted, the difference from the flat rotation reduces to $\\Delta v=$22 km s$^{-1}$. However, even if this is the case, the discrepancy between the observed proper motion and the proper motion expected from the flat rotation curve still remains. This may suggest the Galactic rotation itself is slower in the far outer regions. However, since we have only one source in the far outer region and since the proper motion of this source could also be largely affected by modeled internal proper motions of maser spots ($\\sim$43 km s$^{-1}$ as discussed in section 5.2), at this moment we cannot reach at a decisive conclusion. For further conclusion, we have to increase the number of sources in the far outer region for which accurate astrometry is done, which should be definitely one of the important future works of the VERA project. \\bigskip One of the authors (MH) acknowledges financial support from grant-in-aid (No.21244019) from the Ministry of Education, Culture, Sports, Science and Technology (MEXT). Authors also would like to thank all the staffs at Mizusawa VLBI observatory and at Kagoshima University for supporting observations." }, "1101/1101.1953_arXiv.txt": { "abstract": "This Letter presents the first distance measurement to the massive, semi-detached, eclipsing binary LMC-SC1-105, located in the LH~81 association of the Large Magellanic Cloud (LMC). Previously determined parameters of the system are combined with new near-infrared photometry and a new temperature analysis to constrain the reddening toward the system, and determine a distance of $50.6\\pm1.6$ kpc (corresponding to a distance modulus of $18.52\\pm0.07$ mag), in agreement with previous eclipsing binary measurements. This is the sixth distance measurement to an eclipsing binary in the LMC, although the first to an O-type system. We thus demonstrate the suitability of O-type eclipsing binaries (EBs) as distance indicators. We suggest using bright, early-type EBs to measure distances along different sight lines, as an independent way to map the depth of the LMC and resolve the controversy about its three-dimensional structure. ", "introduction": "\\label{section:intro} As one of the nearest galaxies to the Milky Way, the Large Magellanic Cloud (LMC) has naturally been an attractive first rung for the Extragalactic Distance Scale. The {\\it HST} Key Project \\citep{Freedman01} adopted a distance modulus $\\mu=18.50\\pm0.10$ mag (corresponding to a distance of 50.1$\\pm2.4$ kpc) to the LMC, which has since become the consensus in the community. \\citet{Schaefer08} pointed out that overestimation of error bars and band-wagon effects are present in the literature, with pre-2001 LMC distance measurements yielding values between 18.1 and 18.8 mag \\citep[see][]{Benedict02}, and post-2001 values clustering around the Key Project value. Given that different systematic errors accompany each method, a careful comparison of the distances resulting from different methods is necessary to characterize them. Furthermore, there is increasing evidence for substantial and complex vertical structure in the disk of the LMC \\citep[see review by][]{vanderMarel06} from studies of red clump stars \\citep{Olsen02,Subramanian10}, Cepheid variables \\citep{Nikolaev04} and RR Lyrae stars \\citep{Pejcha09}, which demands further exploration. The only direct, geometrical method available for measuring distances to stars in the LMC is with eclipsing binaries (EBs). In particular, the light curve provides the fractional radii of the components, the radial velocity semi-amplitudes determine the masses and size of the orbit, which together with the effective temperature determination (e.g.\\ by comparison with synthetic spectra), yield luminosities and therefore distances \\citep[see reviews by][]{Andersen91,Torres10}. The EB distance method has so far been applied to four early-B type systems \\citep{Guinan98, Ribas02, Fitzpatrick02, Fitzpatrick03} and one G-type giant system \\citep{Pietrzynski09} in the LMC, with individual uncertainties ranging from 1.2 to 2.2 kpc. Four of these systems are located within the bar of the LMC and their individual distances are consistent with the quoted uncertainties, yielding an error-weighted mean value of $49.4\\pm1.1$~kpc. A fifth system, located several degrees away in the north-east quadrant of the disk of the LMC, gives a $3\\sigma$ shorter distance of $43.2\\pm1.8$~kpc. \\begin{figure}[ht] \\includegraphics[width=8.5cm]{fig1.eps} \\caption{Spatial distribution of known EBs from OGLE II and MACHO (blue circles) on the {\\it Spitzer} 3.6$\\mu$m image of the LMC. EBs with measured distances are labeled. Yellow circles mark the most suitable detached EBs for distance determination \\citep{Michalska05}; red circles mark the OGLE II binaries we plan to measure distances to next. The \\ion{H}{1} kinematic center (white ``x'') from \\citet{Kim98} and the dynamical center or center of the bar (green ``x'') from \\citet{vanderMarel02} are labeled; the solid line corresponds to the line of nodes \\citep{vanderMarel02}. Coordinates are given for J2000.} \\label{map} \\end{figure} Figure~\\ref{map} shows the spatial distribution of all known EBs from the OGLE II \\citep{Wyrzykowski03} and MACHO \\citep{Derekas07, Faccioli07} microlensing surveys of the LMC, and the systems with measured distances, overlaid onto the {\\it Spitzer} SAGE image in the IRAC 3.6 $\\mu$m band \\citep{Meixner06}. A magnitude cut ($V<17$ mag) and period cut ($>1.5$ days) were both applied to the EB catalogs to reject foreground systems and faint systems whose immediate follow up is unrealistic or impossible. The detached EBs selected by \\citet{Michalska05} among the OGLE II systems as being most suitable for distance determination are also shown. Both the \\ion{H}{1} kinematic center \\citep{Kim98} and the dynamical center \\citep[or center of the bar;][]{vanderMarel02} are overplotted, as is the line of nodes \\citep[$\\Theta=129.\\!^{\\circ}9\\pm6.\\!^{\\circ}0\\deg$;][]{vanderMarel02}. Motivated by the evidence for vertical structure in the LMC and the one discrepant EB distance, we proceed to compute the distance to LMC-SC1-105\\footnote{Or OGLE J053448.26-694236.4 = MACHO 81.8881.21 = LH~81-72.}. LMC-SC1-105 is a massive, semi-detached, short period ($P=4.25$ days) O-type system, with component masses of $\\rm M_{1}=30.9\\pm1.0\\;\\msun$, $\\rm M_{2}=13.0\\pm0.7\\;\\msun$, and radii of $\\rm R_{1}=15.1\\pm0.2\\;\\rsun$, $\\rm R_{2}=11.9\\pm0.2\\;\\rsun$ \\citep[determined by][]{Bonanos09}. The very accurate measurement of the radii ($<2\\%$) renders the system suitable for a distance determination, given that EB distances are independent of the usual distance ladder and therefore important checks for other methods. However, accurate radii are not sufficient for an accurate distance. Accurate fluxes (i.e.\\ effective temperatures) and extinction estimates are also needed, therefore this Letter sets out to determine these quantities and obtain the distance. Specifically, Section 2 presents new near-infrared photometry of LMC-SC1-105, Section 3 an analysis of the spectra with state-of-the-art model atmospheres, Section 4 the distance determination, and finally, Section 5 a discussion of our results. ", "conclusions": "LMC-SC1-105 is located in the LH 81 association \\citep{Massey00}, near the center of the LMC bar. It contains two early O-type stars and three Wolf-Rayet systems, one of which was recently found to be an EB \\citep{Szczygiel10}. Furthermore, this association resides in the superbubble N 154 \\citep{Henize56} = DEM 246 \\citep{Davies76}. We have determined a large value of $R_V=5.8\\pm0.4$ toward LMC-SC1-105, however, such high values are not uncommon. \\citet{Cardelli89} find $55$ for 12 out of the 328 stars in their sample. Large values of $R_V$ simply imply larger dust grain sizes, which are expected to occur in dense regions of the interstellar medium due to accretion and coagulation of grains. We therefore conclude that the environment in which LMC-SC1-105 resides has large dust grains. In this Letter, we have determined the distance to LMC-SC1-105 and consequently the LMC bar to be $50.6\\pm1.6$ kpc ($\\mu=18.52\\pm0.07$ mag). The agreement we find with previous EB distances to systems in the bar with different spectral types testifies to the robustness of the EB method and its potential as a powerful, independent distance indicator. Furthermore, it confirms that O-type (and semi-detached) EBs are suitable for distance determination, i.e.\\ that the fluxes predicted by FASTWIND are indeed accurate. EB-based distance determinations to M31 \\citep{Ribas05, Vilardell10} and M33 \\citep{Bonanos06} can therefore provide an independent absolute calibration of the Extragalactic Distance Scale. Future distance determinations to EBs in the LMC (e.g.\\ those marked in Figure~\\ref{map}), will additionally provide $R_V$ values in different environments of the LMC. Finally, we suggest using bright, early-type EBs to measure distances along different sight lines to the LMC, as an independent way to map its depth and resolve the controversy about its vertical structure." }, "1101/1101.1448_arXiv.txt": { "abstract": "IceCube has become the first neutrino telescope with a sensitivity below the TeV neutrino flux predicted from gamma-ray bursts if GRBs are responsible for the observed cosmic-ray flux above $10^{18}$ eV. Two separate analyses using the half-complete IceCube detector, one a dedicated search for neutrinos from $p \\gamma$-interactions in the prompt phase of the GRB fireball, and the other a generic search for any neutrino emission from these sources over a wide range of energies and emission times, produced no evidence for neutrino emission, excluding prevailing models at 90\\% confidence. ", "introduction": " ", "conclusions": "" }, "1101/1101.4938_arXiv.txt": { "abstract": "The most massive elliptical galaxies apparently formed the fastest, because the ratio of $\\alpha$ elements (such as oxygen) to iron is the smallest. In fact, iron is mainly produced from type Ia supernovae on a timescale of $\\sim 0.1-1$ billion years, while the $\\alpha$ elements come from massive stars on timescales of a few tens of million years (Matteucci 1994). Reproducing such a $\\alpha$/Fe correlation has long been a severe problem for cosmological theories of galaxy formation, which envisage massive galaxies to assemble gradually from smaller progenitors, and to be characterized by a star formation history too much extended towards late cosmic times. While it has recently become clear that feedback from Active Galactic Nuclei (AGN) activity play a role in the late quenching of star formation (e.g. Cattaneo et al. 2009), and that early star formation history in the galaxy progenitors affect the $\\alpha$/Fe ratio (Calura \\& Menci 2009), major mergers alone cannot enhance the star formation in the high-redshift progenitors to the levels required to match the steepness of the observed $\\alpha$/Fe correlation (Spolaor et al. 2010). Here we report that the inclusion of the effects of fly-by 'harassments', that trigger lower level starbursts, combined with the AGN quenching of the starburst activity, considerably enhances the capability to account for the observed $\\alpha$/Fe ratio in ellipticals within cosmological galaxy formation models . The critical difference between the earlier work and the present result is the effect of starbursts driven by fly-by encounters that would have been very common amongst the high-redshift progenitors of massive galaxies and which would have boosted star formation in the first 2 billion years after the Big Bang, combined with quenching of the burst activity within the first 3-4 Gyr. ", "introduction": "The correlation between stellar $\\alpha$/Fe and velocity dispersion observed in local ellipticals (Trager et al. 2000; Thomas et al. 2005; Bernardi et al. 2006; Graves et al. 2007; S\\'anchez-Bl\\'azquez et al. 2006, Spolaor et al. 2010; Zhu et al. 2010) is generally referred to as one of the evidences of the ``downsizing'' pattern of local galaxies, a pattern which is observable up to high redshift (e.g. Cowie et al. 2006; Kodama et al. 2004; Clements et al. 2008) indicating that the more massive the galaxy, the shorter the star formation timescale (Matteucci 1994; Renzini 2006; Pipino et al. 2010; Rogers et al. 2010). \\\\ The failure of cosmological galaxy formation models (Thomas et al. 2005; Nagashima et al. 2005; Spolaor et al. 2010) in accounting for the above correlation is in fact due to such a ''downsizing'' aspect. In fact, the precision measurements of the power spectrum of the initial density perturbations (Percifal et al. 2007) provide evidence for large-scale fluctuations to be smaller, on average, compared to small-scale ones - in accord to what expected in a Cold Dark Matter (CDM) scenario; this implies a hierarchical build-up of dark matter (DM) haloes, where small-mass galaxies collapse first, and later assemble to form massive ellipticals (De Lucia et al. 2006). The deep physical origin of the disagreement between the observed and the predicted $\\alpha/Fe-\\sigma$ correlation has been often considered as an evidence for a fundamental flaw of cosmological CDM models of galaxy formation (Thomas et al. 2005; Baugh 2006). To produce downsizing within standard galaxy formation models, some physical mechanisms are required which can enhance star formation in massive galaxies at early times, possibly quenching it at late times. In cosmological galaxy formation, environmental effects are likely to play a major role in producing such a scenario. The progenitors of massive galaxies formed from biased regions of the density field, and fly-by events and merging in such regions may trigger starbursts at early cosmic times ($t \\leq 2.5$ Gyr), when the Universe was denser and the rate of galaxy encounters was at least 5-10 times higher than the present. A fundamental mechanism to quench star formation in massive galaxies is linked to their Active Galactic Nuclei (AGN). Feedback from powerful AGN may switch off star formation when the progenitors assembled into the main progenitor of the final elliptical (Cattaneo et al. 2009). These processes may concur in yielding stellar populations in massive galaxies characterized by a prompt star formation at early epochs followed by a quenching, in agreement with what suggested by the observed $\\alpha/Fe-\\sigma$ correlation. Thus, testing whether the latter processes may in fact explain the above correlation is crucial for assessing the consistency of CDM models, and requires developed modelling which includes all the physical processed detailed above. \\\\ In a previous paper (Calura \\& Menci 2009, hereinafter CM09), we investigated the chemical properties of local galaxies within a cosmological hierarchical clustering scenario through a semi-analytic model (SAM) of galaxy formation. We used a hierarchical semi-analytic model where chemical evolution was computed by taking into account the stellar lifetimes, a significant step forward with respect to the instantaneous recycling approximation. However, chemical abundances were computed by considering the total (summed over all progenitors) star formation history of each model galaxy and assumed an average interstellar $\\alpha$/Fe. This treatment was suited to the study of interstellar abundances, comparable to the ones observed today in stars in the solar neighbourhood, but in many cases it lead us to underestimate the integrated stellar $\\alpha$/Fe abundances measured in local early type galaxies. \\\\ In this paper, by means of the same SAM, we perform a thorough study of the the average stellar $\\alpha$/Fe in early type galaxies, taking into account the contribution from all the single progenitors. Although computationally more expensive, this constitutes a step forward with respect to our earlier work (CM09), since thanks to the inclusion of the contribution of all the progenitors, our estimates of the integrated abundances are more accurate. A second enhancement with respect to our previous work is the calculation of the luminosity-weighted abundances (against the mass-weighted ones used in the first paper), another aspect which improves the comparison of our results to local observations. \\\\ The aim of this paper is to show the relative roles of the various processes in shaping the $\\alpha/Fe-\\sigma$ correlation, stressing the importance of some processes which have been neglected in previous studies, such as interaction-triggered fly-by events and merging-triggered starbursts. \\\\ This paper is organized as follows. In Sect. 2, our model is briefly described. In Sec. 3, we present our main results, which are then discussed in a broader context in Sect. 4. ", "conclusions": "Our results show that a standard CDM galaxy formation model including ITS, consisting in both merger-triggered and fly-by triggered starbursts, and AGN feedback can naturally lead to shorter star formation timescales in larger galaxies, i.e. to star formation histories reversed with respect to their mass-assembly histories.\\\\ The absolute novel ingredient of our model consists in the simultaneous inclusion of both the starbursts triggered by fy-by events, absent in any previous investigation of this fundamental property of early-type galaxies, and the AGN feedback related to the active Quasar phase of AGNs.\\\\ The general problem of quenching the star formation histories in massive galaxies have been investigated in several studies (e.g. Gabor et al. 2010) and is of great interest. Fig. 2 is useful to better understand how the features of our model act directly on the star formation histories of single galaxies. In each panel of this figure, we show the SFH of three galaxies drawn from the sample including ITS and AGN, including AGN but not ITS and from the one including ITS but not the effects of AGNs. The three galaxies shown in each panel present the same merger histories. It is worth to stress that in each case, the final stellar masses are not the same, since the inclusion of physical ingredient such as ITS and AGN can considerably alter the star formation hisory and the feedback history of the systems. The final stellar mass of each model is reported in Fig. 2 (see caption for further details).\\\\ It is important to note that the suppression of ITS has effects also on the efficiency of the AGN feedback, since, according to the model used here and described in Menci et al. (2008), the cold gas which causes the interact-triggered starbursts feeds also the nuclear activity. In fact, in every panel, both models without ITS and AGN present higher initial star fromation rates with respect to the standard case. The explanation of this is that in both cases (no ITS and no AGN) the AGN is inefficient in heating and ejecting gas from the galaxies, resulting in higher amount of cold gas when the star formation has already turned on.\\\\ At the present time, the models which include ITS but without AGN feedback show SFR values larger than those shown by the ones drawn from the standard sample. Globally, their SFHs are less skewed towards early times with respect to the ones of the model with both AGN and ITS. The effects of AGNs in quenching the star formation activity within the first 3-4 Gyr is evident from this figure. Finally, the model with AGN but without ITS presents very extended SFHs, with rather flat and smooth behaviours and very high present-day values. Without the effect of ITS, gas consumption at early times is not very efficient and at later times, after a few Gyrs, the large amounts of cold star-forming gas prevent AGNs from being effective and quenching the star formation. It is therefore clear that, in order to have realistic star formation histories for early type galaxies, both effects of ITS and AGN feedback are necessary.\\\\ Both effects of AGN feedback and ITS seem to play an important role in determining galactic downsizing. However, the effect of the former seems slightly dominant, and this is supported by the fact that the predicted $\\alpha/Fe$-$\\sigma$ relation without ITS is slightly shallower than the one without AGN. \\\\ Our study of the $\\alpha/Fe$-$\\sigma$ relation tells us that our galaxy formation model can naturally account for one of the most important aspects of the ``downsizing'' character of early type galaxies. Downsizing (i.e. oldest stars are in most massive galaxies, star formation shifts to lower masses at late times) in hierarchical models will arise if there is a minimum mass required to have star formation and a maximum mass above which feedback from AGN suppresses further star formation (Sheth 2003; Croton \\& Farrar 2008). Together, these will also lead to an [$\\alpha$/Fe]-sigma relation, since in principle, in such a scenario, the distribution of the formation timescales will be narrower for the stars forming in the most massive systems. In this framework, ITS help because it allows to use up more of the gas earlier, thus leaving less work to be done by the AGN feedback, and leading to a steeper [alpha/Fe]-sigma relation. A detailed calculation of chemical abundances in more hirarchical models (such as the ones above) will be of further help in understanding the relevance of these processes with respect to the downsizing character of local galaxies. \\\\ It is important to stress that the observations can be reproduced without the modification of any fundamental chemical evolution parameter, such as the stellar IMF. However, it is worth to stress that such ingredient may play some role in determing the shape of the $\\alpha/Fe$-$\\sigma$ relation. In fact, Although our unprecedented rendition of this relation, the predicted slope obtained with our standard model is slightly shallower than the observed one. This aspect will be investigated in our future work, and may have its explanation in effects such as a possible IMF dependence on the SFR and/or a flatter IMF in starbursts (Recchi et al. 2009; Haas \\& Anders 2010; Calura et al. 2010). Some direct evidences in favour of the latter hypotesis come from observational studies of nuclear star clusters in the Milky Way (Bartko et al. 2010). Also dynamical investigations of early-type galaxies seem to require stronger baryonic feedback at the epoch of the formation of these systems, fully consistent with an IMF skewed towards more massive stars (Napolitano et al. 2010), or from the study of local ultra-compact galaxies (Dabringhausen et al. 2009). A top-heavy IMF at early times seems also required in order to reproduce the abundance ratios observed in the hot intracluster medium (Matteucci \\& Gibson 1995; Loewenstein \\& Mushotzky 1996).\\\\ In our previous work (C09), the average stellar abundance ratios have been calculated without taking into account the contribution of the single progenitors, i.e. by considering the total star formation history of the selected galaxy, given by the sum of the star formation rates of all the progenitors, and considering an average interstellar $\\alpha$/Fe. In that paper, the observed $\\alpha$-Fe-$\\sigma$ relation could be partially explained by assuming a top-heavy IMF in the most massive galaxies, which mimics the effect of a higher star-formation efficiency (or, in other words, a shorter star formation timescale) in largest galaxies (Matteucci 1994; Ferreras \\& Silk 2003; Matteucci et al. 1998).\\\\ The largest galaxies have the most massive progenitors, which form all their stars in the shortest timescales and dominate the integral of Eq.~\\ref{afe}. This is the main reason why here, by considering the SFRs and the abundance ratios of each single progenitor, instead of a total SFR and an abundance averaged over all the progenitors, we can successfully reproduce the observed $\\alpha/Fe$-$\\sigma$ relation without any modification of the main chemical evolution parameters. \\\\ \\begin{figure*} \\centering \\vspace{0.001cm} \\epsfig{file=afe.ps,height=8cm,width=9cm} \\caption{Stellar average $\\alpha/Fe$ ratio vs velocity dispersion compared to local observations. The colour code represents the predicted number of galaxies with a given $\\alpha/Fe$ and a velocity dispersion $\\sigma$, normalised to the total number of galaxies with that velocity dispersion. The solid lines represent linear-regression fits to the model. The dashed line and dash-dotted line represent linear-regression fits to the observational $\\alpha/Fe$-$\\sigma$ relations, as compiled by Spolaor et al. (2010; open diamonds) and Arrigoni et al. (2009; solid circles) respectively. \\emph{Top-left panel}: our results have been computed by means of our standard assumptions, i.e. by taking into account both interaction-triggered starbusts at high redshift and AGN feedback. \\emph{Top-right panel}: no interaction-triggered starbursts and fly-by interactions. \\emph{Bottom-left panel}: no AGN feedback. \\emph{Bottom-right panel}: no AGN feedback, no fly-by interactions. } \\label{fig1} \\end{figure*} \\begin{figure*} \\centering \\vspace{0.001cm} \\epsfig{file=sfr.eps,height=8cm,width=9cm} \\caption{Star Formation histories of a few early-type galaxies of our model. In each panel, we show the SFH of three galaxies drawn from the sample including ITS and AGN (solid black lines), including AGN but not ITS (cyan dotted lines) and including ITS but not AGNs (dashed red slines). In each panel, the three systems present the same merging history. The final stellar masses of the models are indicated in each panel. From top to bottom, the indicated masses refer to the standard case, the No ITS case and the No AGN case. } \\label{fig2} \\end{figure*}" }, "1101/1101.4066.txt": { "abstract": "We present a multi-wavelength study of NGC 4330, a highly-inclined spiral galaxy in the Virgo Cluster which is a clear example of strong, ongoing ICM-ISM ram pressure stripping. The HI has been removed from well within the undisturbed old stellar disk, to 50\\% - 65\\% of R$_{25}$. Multi-wavelength data (WIYN\\footnote{The WIYN Observatory is a joint facility of the University of Wisconsin-Madison, Indiana University, Yale University, and the National Optical Astronomy Observatory. } BVR-H$\\alpha$, VLA\\footnote{The VLA is operated by the National Radio Astronomy Observatory, which is a facility of the National Science Foundation (NSF), operated under cooperative agreement by Associated Universities, Inc.} 21-cm HI and radio continuum, and GALEX\\footnote{This work is based in part on observations made with the NASA Galaxy Evolution Explorer. GALEX is operated for NASA by the California Institute of Technology under NASA contract NAS5-98034.} NUV and FUV) reveal several one-sided extraplanar features likely caused by ram pressure at an intermediate disk-wind angle. At the leading edge of the interaction, the H$\\alpha$ and dust extinction curve sharply out of the disk in a remarkable and distinctive ``upturn\" feature that may be generally useful as a diagnostic indicator of active ram pressure. On the trailing side, the ISM is stretched out in a long tail which contains 10\\% of the galaxy's total HI emission, 6 - 9\\% of its NUV-FUV emission, but only 2\\% of the H$\\alpha$. The centroid of the HI tail is downwind of the UV/H$\\alpha$ tail, suggesting that the ICM wind has shifted most of the ISM downwind over the course of the past 10 - 300 Myr. Along the major axis, the disk is highly asymmetric in the UV, but more symmetric in H$\\alpha$ and HI, also implying recent changes in the distributions of gas and star formation. The UV-optical colors indicate very different star formation histories for the leading and trailing sides of the galaxy. On the leading side, a strong gradient in the UV-optical colors of the gas-stripped disk suggests that it has taken 200-400 Myr to strip the gas from a radius of $>$8 to 5 kpc, but on the trailing side there is no age gradient. All our data suggest a scenario in which NGC~4330 is falling into cluster center for first time and has experienced a significant increase in ram pressure over the last 200-400 Myr. ", "introduction": "There is ample evidence that cluster galaxies are strongly affected by their environment over time. Denser environments in the nearby universe have a higher fraction of elliptical and S0 galaxies (the density-morphology relationship of Dressler 1980) and the spirals are systematically HI-deficient (Haynes \\& Giovanelli 1986, Solanes et al. 2001), clusters at z$>$0.1 have a higher fraction of blue, star-forming galaxies than nearby clusters (Butcher \\& Oemler 1978, 1984), and at z$\\sim$0.5, the fraction of spirals increases while the fraction of S0's decreases relative to that in the local universe (Dressler et al. 1997). This suggests that one or more cluster processes act to transform spiral galaxies into earlier-type spirals and S0's. Many effects have been presented as possible contributors, including galaxy-galaxy harrassment (Bekki 1999, Moore et al. 1998), galaxy-galaxy collisions (Kenney et al. 2008), galaxy interactions with the cluster potential (Moore et al. 1999), ram pressure stripping (Gunn \\& Gott 1972, Abadi et al. 1999, Poggianti et al. 1999), turbulent viscous stripping (Nulsen, 1982), and starvation (Larson et al. 1980). There remains significant debate about the relative importance of these processes. Several recent studies of clusters out to a redshift of z$\\sim$0.8 support the importance of ICM-ISM interactions as a driver of cluster galaxy evolution. Many studies have found cluster populations of disky, quiescent galaxies with structure similar to spiral galaxies but with little or no star formation (e.g., van den Bergh 1976, Balogh et al. 1998, van der Wel et al. 2010). Such galaxies are consistent with gas stripping. Moran et al. (2007) present evidence from two z$\\sim$0.5 clusters that there are multiple pathways for star formation to be quenched in spirals, depending on the properties of the cluster, with ICM-ISM stripping emerging as a major contributing factor in situations where the ICM is dense enough for stripping to be effective. From a spectral analysis of galaxies in various environments at z = 0.4 - 0.8, Poggianti et al. (2009) find that that clusters are the preferred location of k+a galaxies. Such galaxies have had their star formation was quenched between 5$\\times10^7$ and 1.5$\\times10^9$ years prior to observation, making them a likely transition phase between actively star-forming spirals and passively evolving S0 and Sa galaxies. Since the star formation quenching efficiency increases with cluster velocity dispersion, Poggianti et al. (2009) propose this transformation is likely the result of ICM-ISM interactions. These studies document how key galaxy properties differ for large samples of cluster galaxies, but do not by themselves offer direct evidence of the mechanism transforming cluster galaxies. %Several recent studies of clusters out to a redshift of z$\\sim$0.8 support the importance of ICM-ISM interactions as a driver of cluster galaxy evolution. van der Wel et al. (2010) find that at z=0.04 - 0.08, there is a population of disky, quiescent galaxies. They conclude that the cluster galaxy population changes over time due to ISM stripping, since the quiescent galaxies retain similar structure to spiral galaxies but lack star formation. Moran et al. (2007) present evidence from two z$\\sim$0.5 clusters that there are multiple pathways for star formation to be quenched in spirals, depending on the properties of the cluster, with ICM-ISM stripping emerging as a major contributing factor in situations where the ICM is dense enough for stripping to be effective. Poggianti et al. (2009) find a correlation between cluster mass and k+a galaxy fraction in galaxies at z = 0.4 - 0.8. The spectra of k+a galaxies show that their star formation was quenched between 5$\\times10^7$ and 1.5$\\times10^9$ years prior to observation, making them a likely transition phase between actively star-forming spirals and passively evolving S0 and Sa galaxies. Poggianti et al. (2009) present evidence that this transformation is likely the result of ICM-ISM interactions, due in part to the fact that the star formation quenching efficiency increases with cluster velocity dispersion. In order to understand the impact of ram pressure stripping on galaxy evolution, we need to be able to identify galaxies which are clearly experiencing stripping or did at some point in the past, and determine which of a given galaxy's characteristics are due to ram pressure stripping. In particular, we are interested in finding out how the multi-phase ISM behaves during stripping, how long it takes to strip a galaxy, and how ram pressure and the associated stripping affects the rate and distribution of star formation. Through detailed studies of the best cases of ram pressure stripping in the nearby Virgo cluster, we hope to identify key parameters of the ram pressure stripping interaction, such as the current strength of ram pressure, the angle between the ICM wind and the disk, the time evolution of these quantities, and the evolutionary stage of the stripping event. If we can quantify key stripping parameters for enough galaxies, we can understand the impact of stripping on galaxy evolution in different environments. There have been many ram pressure stripping simulations over the past few years (e.g., Quilis et al. 2000, Schulz \\& Struck 2001, Roediger \\& Bruggen 2006, Tonnesen et al. 2007, Tonnesen \\& Bryan 2008, 2009, 2010, Kronberger et al. 2008, Vollmer 2009, Jachym et al. 2009; also see review by E. Roediger 2009). They are all in general agreement that in most situations, ram pressure stripping is well-approximated by the simple Gunn \\& Gott (1972) formula, which fails only in cases of nearly edge-on stripping and short, impulsive episodes of ram pressure (Jachym et al. 2009). All of the simulations predict that the gas stripped from the disk will form a tail downwind of the galaxy, and all of the groups predict a lack of star formation in the stripped outer disks. Kronberger et al. (2008) and Tonnesen \\& Bryan (2009) both predict enhanced star formation rates due to ICM pressure in certain cases. Progress has been made in modeling the effects of the stripping process on the multi-phase ISM. For example, Tonnesen \\& Bryan (2009, 2010) model the effects of ram pressure on a clumpy ISM with self-gravity and cooling, and also account for the effects of heating. However, more work remains to be done. Many complicated physical processes relevant to the ISM, particularly star formation, are simulated using simple prescriptions, and observations are necessary to determine what actually happens. There have also been spectacular examples of individual galaxies being stripped in rich clusters between z$\\sim$0.2 and 0.5 (Owen et al. 2006, Cortese et al. 2007, Sun et al. 2007), but at such distances we lack the resolution to see what really happens. The Virgo cluster is the nearest moderately rich cluster, and provides the best resolution to observe cluster processes at work. For this reason, we have carried out the VIVA (VLA Imaging of Virgo spirals in Atomic gas) survey, which presents high-quality HI imaging of 53 spirals in the Virgo cluster (Chung et al. 2009). Using data from the VIVA survey, Chung et al. (2007) noted the presence of HI tails in at least 25\\% of large spiral galaxies, a total of 7 galaxies, near the cluster's virial radius, providing striking evidence of galaxies actively losing their gas as they plummet towards the cluster center. One of these galaxies was NGC 4330, and we have followed up with a detailed multi-wavelength study. In this paper, we show that NGC 4330 is one of the best examples of ram pressure stripping actively transforming a spiral galaxy in the nearby Virgo cluster. There is a wealth of evidence that it is being actively stripped, and several interesting phenomena can be clearly seen, some for the first time. In this multi-wavelength study, we examine the effects of stripping on the galaxy's current star formation, recent star formation history, and the stripping of the HI and dust components of its ISM. In Section \\ref{thegal}, we introduce the galaxy NGC 4330. In Section \\ref{observations}, we describe the observations taken in optical broad-band and H$\\alpha$ narrow-band filters, HI, and the UV. In Section \\ref{results}, we describe the galaxy's stellar distribution and optical dust extinction in the BVR bands, the HI distribution and kinematics, and the H$\\alpha$ and UV distributions. In Section \\ref{discussion} we discuss the interpretation of the galaxy's properties in the context of an ongoing ICM-ISM interaction. In Section \\ref{conclusion}, we summarize our findings. Throughout this paper, we assume a distance to the Virgo Cluster of 16 Mpc (Yasuda et al. 1997). ", "conclusions": "\\label{conclusion} NGC 4330 provides an outstanding opportunity to study the effects of active ram pressure stripping on a nearby galaxy. The wind angle, evolutionary stage, and galaxy orientation make it possible to clearly observe distinct behavior at the leading edge (the upturn) and the trailing side (the tail) of the ICM-ISM interaction. Its stellar disk is undisturbed, with no warps, bridges, strong asymmetries, or other evidence of gravitational interaction. Ram pressure stripping is the only plausible explanation for its radially truncated, one-sided extraplanar HI distribution and undisturbed old stellar disk. We conclude by describing the morphological features related to stripping, how we think they were created, and the derived evolutionary timescales for changes in the galaxy over the past $\\sim$400 Myr. Finally we provide a self-consistent description of the timeline of the galaxy's stripping history which takes into account these morphological features and their evolutionary timescales. \\begin{enumerate} \\item \\textbf{Morphological Features Related to Stripping and Derived Stripping Interaction Parameters:} \\begin{itemize} \\item ISM truncation along the major axis: The galaxy's ISM has been stripped to well within the undisturbed old stellar disk. Its dust extinction and HI are truncated at roughly the same radii, $\\sim 50\\% - 65\\%$ of R$_{25}$ on both the leading and trailing sides of the galaxy. \\item Extraplanar HI: The galaxy's HI distribution is strongly asymmetric as a result of having been removed from the disk by the ICM - 32\\% of the emission is extraplanar on the side we identify as downwind of the major axis. This includes 10\\% of the galaxy's HI in a long tail. \\item UV major-axis asymmetry: Recent star formation as traced by UV is quite asymmetric about the galaxy's nucleus - the UV emission strength drops rapidly at 0.5 $R_{25}$ on the upwind side, but gradually tapers off until R$_{25}$ on the downwind side. The UV emission extends radially further on the upwind side than other tracers of the ISM and star formation - ongoing star formation as traced by H$\\alpha$ is truncated at $\\sim0.5 R_{25}$, roughly the same radii as the other ISM tracers, and is relatively symmetric about the galaxy's nucleus. The different extents of H$\\alpha$ and UV emission indicate that star formation extended to larger radii when the UV-bright stellar populations formed a few hundred Myr ago, but was much more radially truncated when the H$\\alpha$-bright regions formed $<$10 Myr ago. Based on the UV color gradient in the outer disk on the galaxy's leading side, we estimate that it has taken 200-400 Myr to strip the disk from $>$8 kpc to 5 kpc. However, there is no color/age gradient on the trailing side, so the effects of stripping have probably been more uniform and recent on the trailing side. \\item The ``upturn\": At the leading edge of the ICM-ISM interaction, the galaxy's H$\\alpha$ emission veers abruptly out of the disk midplane, and a large, significantly obscuring dust cloud extends beyond the H$\\alpha$ upturn emission at a projected angle almost perpendicular to the disk. This striking ``upturn\" feature is near the HI truncation boundary, and the UV emission extends radially beyond it, suggesting that the ISM in the disk beyond the upturn was removed only recently. The upturn's H$\\alpha$ morphology indicates that most of the ISM in this region has been pushed out of the disk by ram pressure. The galaxy also displays a strong, localized radio deficit region near the upturn, characteristic of galaxies experiencing strong ram pressure. \\item The tail: NGC 4330 has a long tail on its trailing side which is visible in HI, UV, and H$\\alpha$. The HI tail is offset downwind from the the UV/H$\\alpha$ tail, suggesting that much of the ISM has moved from the location of the UV tail to the current location of the HI tail. There is significant UV emission from the tail but the H$\\alpha$ is sparse, indicating that there was significant star formation in the past $\\sim$10 - 200 Myr but there has been little in the past $\\sim$5 Myr. We estimate that it has taken the ISM tail $\\sim$10 - 300 Myr to move downwind from the position of the UV-bright tail to the current position of the HI tail. This compares with the galaxy's rotational period of $\\sim$400 Myr at a radius of 5 kpc. \\item Extraplanar UV-bright regions: We detect 9 regions downwind of the galaxy but outside the tail that are bright in the UV, but for the most part lack H$\\alpha$ emission. A comparison of their UV/optical colors with stellar population models indicates that most of the regions have ages $<$ 100 Myr, though the oldest is $\\sim$350 Myr. Based on the ages and distances of the most distant regions from the galactic plane and the kinematics of the associated HI, it is likely that star formation in these regions was initiated after the gas left the disk. \\item The projected ICM wind angle: We discuss three methods to constrain the projected ICM wind angle from observations: the morphology of the large scale stripped gas tail, the orientation of elongated small-scale dust features, and the shape of the radio deficit region. We give quantitative estimates for all three techniques in NGC 4330, and constrain its disk-ICM wind angle to be 30-70$\\degree$. \\end{itemize} %\\item \\textbf{Derived Evolutionary Timescales:} %\\begin{itemize} %\\item Leading edge stripping: Based on the UV color gradient in the outer disk on the galaxy's leading side, we estimate that it has taken 200-400 Myr to strip the disk from $>$8 kpc to 5 kpc. However, there is no color/age gradient on the trailing side, so it may not have been as strongly affected by the ICM wind. %\\item Tail evolution: We estimate that it has taken the ISM tail $\\sim$10 - 300 Myr to move downwind from the position of the UV-bright tail to the current position of the HI tail. This compares with the galaxy's rotational period of $\\sim$400 Myr at a radius of 5 kpc. %\\item Extraplanar UV region star formation: By comparing their UV colors with stellar population models, we have determined that the extraplanar UV-bright regions of recent star formation downwind of the galaxy are all $<$ 400 Myr old. %\\end{itemize} \\item \\textbf{Stripping History:} The observations are consistent with the following scenario. Ram pressure increased rapidly $\\sim200-400$ Myr ago, stripping gas from the disk on the leading edge side from $>$8 to 5 kpc. The disk was stripped asymmetrically, and for a while the gas disk was stripped more deeply on the leading side than the trailing side. After $\\sim$1 rotation period, ($\\sim$400 Myr) at this higher pressure, the gas disk has become more symmetric, as currently observed. Within the last $\\sim$400 Myr, as the gas disk has gone from highly asymmetric to fairly symmetric, gas on the trailing side has left the disk to form the base of the tail, and the tail has continued to move from the disk. It is not yet clear whether smooth or abrupt changes in ram pressure are responsible for this galaxy's recent, dramatic evolution. \\end{enumerate}" }, "1101/1101.3942_arXiv.txt": { "abstract": "{}{}{}{}{} \\abstract {SN 2008D, a core collapse supernova at a distance of 27 Mpc, was serendipitously discovered by the \\textit{Swift} satellite through an associated X-ray flash. Core collapse supernovae have been observed in association with long gamma-ray bursts and X-ray flashes and a physical connection is widely assumed. This connection could imply that some core collapse supernovae possess mildly relativistic jets in which high-energy neutrinos are produced through proton-proton collisions. The predicted neutrino spectra would be detectable by Cherenkov neutrino detectors like IceCube. } {A search for a neutrino signal in temporal and spatial correlation with the observed X-ray flash of SN 2008D was conducted using data taken in 2007-2008 with 22 strings of the IceCube detector. } {Events were selected based on a boosted decision tree classifier trained with simulated signal and experimental background data. The classifier was optimized to the position and a ``soft jet'' neutrino spectrum assumed for SN 2008D. Using three search windows placed around the X-ray peak, emission time scales from $100 - 10000$ s were probed.} {No events passing the cuts were observed in agreement with the signal expectation of 0.13 events. Upper limits on the muon neutrino flux from core collapse supernovae were derived for different emission time scales and the principal model parameters were constrained. } {While no meaningful limits can be given in the case of an isotropic neutrino emission, the parameter space for a jetted emission can be constrained. Future analyses with the full 86 string IceCube detector could detect up to $\\sim$100 events for a core-collapse supernova at 10 Mpc according to the soft jet model. } ", "introduction": "Observations in recent years have given rise to the idea that core collapse supernovae (SNe) and long duration gamma-ray bursts (GRB) have a common origin or may even be two different aspects of the same physical phenomenon, the death of a massive star with $M >8\\,M_{\\odot}$ (for a review, see Woosley, Bloom \\cite{woosley}). Like GRBs, SNe could produce jets, though less energetic and collimated and possibly ``choked'' within the stellar envelope. Observed associations of supernovae with XRFs, short X-ray flashes with similar characteristics to long GRBs, suggest including XRFs in the SN-GRB connection as well. Although XRFs are considered a separate observational category from GRBs, a common origin and a continuous sequence connecting them have been suggested (Lamb et al. \\cite{lamb}, Yamazakia et al. \\cite{yamazakia}). XRF could be long GRBs with very weak jets or simply long GRBs observed off-axis. Several XRFs or or long duration, soft-spectrum GRBs have been observed in coincidence with core collapse SNe thus far: SN 1998bw (Galama et al. \\cite{galama}), SN 2003lw (Malesani et al. \\cite{malesani}), SN 2003dh (Hjorth et al. \\cite{hjorth}), SN 2006aj (Pian et al. \\cite{pian}), and of course SN 2008D (Soderberg et al. \\cite{soderberg}, Modjaz et al. \\cite{modjaz}, Mazzali et al. \\cite{mazzali}). For SN 2007gr (Paragi et al. \\cite{paragi}) and SN 2009bb (Soderberg et al. \\cite{soderberg2009bb}), two core collapse SNe not associated with an XRF or GRB, recent radio observations provide strong evidence for jets with bulk Lorentz factors of $\\,\\Gamma > 1$. If some core collapse SNe indeed form such \"soft'' jets, protons accelerated within the jet could produce TeV neutrinos in collisions with protons of the stellar envelope (Razzaque et al. \\cite{rmw}, Ando \\& Beacom \\cite{ab}). The soft jet scenario for core collapse SNe can be probed with high-energy neutrinos even if the predicted jets stall within the stellar envelope and are undetectable in electromagnetic observations. On January 9, 2008, the X-ray telescope aboard the SWIFT satellite serendipitously discovered a bright X-ray flash during a pre-scheduled observation of NGC 2770. Optical follow-up observations were immediately triggered and recorded the optical signature of SN 2008D, a core collapse supernova of type Ib at right ascension $\\alpha = 09\\,\\mathrm{h}\\,\\,09\\,\\mathrm{m}\\,\\,30.70\\,\\mathrm{s}\\,$ and declination $\\delta = 33^{\\circ}\\,08'\\,19.1\"\\,$ (Soderberg et al. 2008). SN 2008D offers a realistic chance to detect high-energy supernova neutrinos for the first time since the observed X-ray peak provides the most precise timing information ever available to such a search. Whether or not the existence of jets in aspherical explosions is evidenced in the spectroscopic data for SN 2008D remains highly debated. While Soderberg et al. (\\cite{soderberg}) ``firmly rule out'' any asphericity and Chevalier and Fransson (\\cite{chevalierfransson}) speak of a purely spherical shock-breakout emission, Mazzali et al. (\\cite{mazzali}) and Tanaka et al. (\\cite{tanaka}) find evidence that SN 2008D possessed jets which have been observed significantly off-axis. The IceCube neutrino detector, currently under construction at the South Pole and scheduled for completion in 2011, is capable of detecting high-energy neutrinos ($E_{\\nu} \\gtrsim 100\\,\\mathrm{GeV}$) of cosmic origin by measuring the Cherenkov light emitted by secondary muons with an array of Digital Optical Modules (DOMs) positioned in the transparent deep ice along vertical strings (J. Ahrens et al. \\cite{icecube}). The full detector will comprise 4,800 DOMs deployed on 80 strings between 1.5 and 2.5 km deep within the ice, a surface array (IceTop) for observing extensive air showers of cosmic rays, and an additional dense subarray (DeepCore) in the detector center for enhanced low-energy sensitivity. Each DOM consists of a 25 cm diameter Hamamatsu photo-multiplier tube (PMT, see Abbasi et al. \\cite{pmt}), electronics for waveform digitization (Abbasi et al. \\cite{daq}), high voltage generation, and a spherical, pressure-resistant glass housing. The DOMs detect Cherenkov photons emitted by relativistic charged particles passing through the ice. In particular, the directions of muons (either from cosmic ray showers above the surface or neutrino interactions within the ice or bedrock) can be well reconstructed from the track-like pattern and timing of hit DOMs. Identification of neutrino-induced muon events in IceCube has been demonstrated in Achterberg et al. (\\cite{ic9}) using atmospheric neutrinos as a calibration tool. Sources in the northern sky, like SN 2008D, can be observed with very little background since contamination by atmospheric muon tracks is eliminated by the shielding effect of the Earth. When SN 2008D was discovered, the installation of IceCube was about one quarter completed and the detector was taking data with 22 strings. As shown above, a search for cosmic neutrinos from core collapse SNe is motivated by both observational evidence and theoretical predictions. While analyses using catalogs of SNe/GRBs with timing uncertainties $\\sim$1 d as the signal hypothesis have been performed on archived AMANDA/IceCube data (see Lennarz \\cite{lennarz} for SNe and Abbasi et al. \\cite{grbapj} for GRBs), the unprecedentedly precise timing information available for SN 2008D suggests a designated study of this event. While electromagnetic observations provide no conclusive evidence for the existence of highly relativistic jets, soft, hidden jets could be revealed by high energy neutrinos, assuming sufficient alignment with the line of sight. The paper is organized as follows: Section 2 discusses the assumed model for neutrino production. Section 3 describes the experimental and simulated data used for the analyis. The selection criteria used to separate signal events from background are detailed in Section 4. Section 5 presents the results of the search and constraints derived therefrom. Finally, the analysis is summarized in Section 6. ", "conclusions": "We have searched for high-energy muon neutrinos in coincidence with SN 2008D using data from the IceCube 22 string detector. Using a blind analysis optimized with experimental background and simulated signal data, we observed no events which passed the cuts. From the non-observation, we have derived first constraints on the soft jet model for core collapse SNe under the condition that the predicted jet was pointing in the direction of the Earth. Given the strong dependence of the signal expectation on the model parameters, the non-detection of neutrinos places significant constraints on the principal model parameters. A two dimensional parameter scan in $\\Gamma_{b}$ and $E_j$ shows that the jet Lorentz factor is generally constrained to $\\Gamma_b<4$ for jet energies $E_{j}>10^{51}\\,\\mathrm{erg}$. As mentioned above, the constraints quoted here only hold if the assumed jet of SN 2008D was pointing towards Earth. IceCube is now operating in an additional mode, scanning online data for neutrino bursts, i.e. two nearly collinear neutrinos within 100 s, in real time. If a burst is detected, IceCube triggers optical follow-up observations searching for a SN in the corresponding direction (Franckowiak et al. \\cite{ofu}). Constantly monitoring the entire northern sky, this approach has the potential to generalize the constraints obtained from studying individual objects." }, "1101/1101.3459_arXiv.txt": { "abstract": "The massive clump G10.6-0.4 is an OB cluster forming region, in which multiple UC H\\textsc{ii} regions have been identified. In the present study, we report arcsecond resolution observations of the CS (1--0) transition, the NH$_{3}$ (3,3) main hyperfine inversion transition, the CH$_{3}$OH J=5 transitions, and the centimeter free--free continuum emissions in this region. The comparisons of the molecular line emissions with the free--free continuum emissions reveal a 0.5 pc scale massive molecular envelope which is being partially dispersed by the dynamically--expanding bipolar ionized cavity. The massive envelope is rotationally flattened and has an enhanced molecular density in the mid--plane. In the center of this massive clump lies a compact ($<$0.1 pc) hot ($\\gtrsim$100 K) toroid, in which a cluster of O--type stars has formed. This overall geometry is analogous to the standard core collapse picture in the low--mass star forming region, with a central (proto--)stellar object, a disk, an envelope, and a bipolar outflow and outflow cavity. However, G10.6-0.4 has a much larger physical size scale ($\\le$0.1 pc for typical low--mass star forming core). Based on the observations, we propose a schematic picture of the OB cluster forming region, which incorporates the various physical mechanisms. This model will be tested with the observations of other embedded OB clusters, and with numerical simulations. \\vspace{1cm} ", "introduction": "\\label{chap_introduction} A M$\\gtrsim$10 M$_{\\odot}$ massive star has a short Kelvin--Helmholtz (cooling) timescale, and starts the nuclear burning while it is still deeply embedded in the dense molecular clump. During this earliest evolutionary stage, the massive stars can continue to increase their stellar mass via accretion of more dense molecular gas (Kahn 1974; Garay and Lizano 1999). Meanwhile, these massive stars produce large amounts of strong ionizing photons, and create the \\textbf{Ultracompact (UC)} \\textbf{H}\\textsc{ii} \\textbf{Regions}. The UV radiation (Krumholz, Mckee, Klein 2005; Peters et al. 2010), the pressure of the ionized gas (Keto 2002b; Krumholz et al. 2009; Peters et al. 2010) , and the strong stellar wind directly interact with the surrounding molecular gas. In addition, as massive stars typically form in clusters, their environments may be also influenced by their interactions both radiatively and dynamically. These interactions potentially can disperse the accretion flow, unless the system is confined or trapped by the high density molecular gas with the right physical profiles (example solutions see Keto 2002b, 2003, 2007; also see Galv{\\'a}n-Madrid et al. 2008 for an observational evidence of accretion). This is the greatest contrast to the formation of the low--mass stars, where accretion only occurs before the pre--main sequence phase when the surrounding gas is neutral. To understand how massive stars finally attain a high stellar mass (e.g. $>$20 M$_{\\odot}$), it is necessary to understand the accretion after the formation of an UC \\textsc{Hii}. \\textit{One key factor in determining the relative importance between these feedback mechanisms as well as gravity, and hence the subsequent evolution of the dense molecular core/envelope, is the overall geometry/morphology of the entire system.} For example, the density distribution of the molecular gas around massive stars determines how the molecular structures are self--shielded from the ionizing photons; and the overall geometry of the system determines how the radiation and the ionized photons can leak out of the system. Owing to their typical large distances (a few kpc), the molecular structures around the UC H\\textsc{ii} regions, are not well resolved in previous observations. How the accretion flows around individual stars and in the entire cluster is unknown. Hence how the OB stars gain mass via accretion after nuclear burning begins, is also unknown. We perform new interferometric observations toward a well--studied O--type cluster forming region G10.6--0.4, in order to resolve the molecular structures. The high column density and high brightness temperature of this source allow us to follow the detailed structures of the molecular gas over an extended region, at arcsecond resolutions. Here we emphasize that the subject under consideration is the molecular and the ionized gas flow around a cluster of OB stars, rather than the accretion around a single massive (proto--)stellar object. This is due to the still insufficient spatial resolution. The UC H\\textsc{ii} region complex G10.6-0.4 is at a distance of 6 kpc (Caswell et al. 1975; Downes et al. 1980), and is still deeply embedded in a dense molecular core/envelope. The brightest UC H\\textsc{ii} region is extremely luminous (integrated flux is 2.6 Jy at 23 GHz) (Ho \\& Haschick 1981; Sollins et al. 2005; Sollins \\& Ho 2005), and contains an OB cluster within a 0.05 pc region. Detections of water and OH masers (Genzel \\& Downes 1977; Ho \\& Haschick 1981; Ho et al. 1983; Hofner \\& Churchwell 1996; Fish et al. 2005 ) indicate that the formation of massive stars is still ongoing. The early NH$_{3}$, CS and C$^{18}$O line observations (Ho \\& Haschick 1986; Keto, Ho, \\& Haschick 1987; Keto, Ho, \\& Haschick 1988; Omodaka et al. 1992; Ho, Terebey, \\& Turner 1994 ) unveiled a 0.5 pc scale molecular envelope. These observations also suggested that the dynamics of the majority of the molecular gas have not yet been strongly disturbed by the YSOs or stellar activities. At the 0.5 pc scale, the general motion of the molecular gas appears to be dominated by gravity, showing rotation along a flattened geometry (Liu et al. 2010). The inner part at the 0.1 pc scale appears to rotate faster while infalling towards the UC H\\textsc{ii} region (Keto, Ho, \\& Haschick 1987; Keto, Ho, \\& Haschick 1988; Liu et al. 2010; Beltr{\\'a}n et al. 2011). Recombination line studies further suggest that the accretion flow continues across the H\\textsc{ii} boundary, while evidence of outflow can be seen in the motions and structures of the H\\textsc{ii} region at the sub--arcsecond angular scale (Keto 2002; Keto \\& Wood 2006). Structures found in projection against the UC H\\textsc{ii} continuum emission from NH$_{3}$ optical depth studies, suggest accretion via a very clumpy and geometrically thick rotating core at the 0.1 pc scale (Sollins \\& Ho 2005). In this paper, we present the new arcsecond resolution (1$''$ corresponds to 0.03 pc at the distance of G10.6-0.4) molecular line (CH$_{3}$OH, NH$_{3}$, CS) data, together with very sensitive and 0$''$.5 resolution 3.6 centimeter continuum observations. The linear scale of the synthesized beams is smaller than the natural scale for self-gravity (Jeans length $\\sim$0.1 pc for temperature $\\sim$20--30 K and H$_{2}$ number density $\\sim$10$^{5}$ cm$^{-3}$). We see an unprecedented detailed morphology of the massive accretion flow which clarifies the geometrical relation between the hot core and the envelope. With the aid of the new results and the discussions on this target source in previous publications, a more thorough understanding is achieved. In massive star forming regions, the NH$_{3}$ inversion transitions are usually regarded as reliable molecular tracers for 20--100 K gas. The main hyperfine components of these NH$_{3}$ lines are optically thick and may mask the detailed density structures. The satellite hyperfine components are more optically thin (see Ho \\& Townes 1983 for a review). However, in most cases, the satellite hyperfine components of the NH$_{3}$ (J, K) = (1,1) inversion transitions are blended with the main hyperfine component; the satellite hyperfine components of the higher (J,K) level transitions are much weaker and are difficult to be observed with high resolutions. With a factor of $\\sim$100 higher Einstein A coefficient, the selected CH$_{3}$OH transitions potentially allow us to probe the denser regions. We simultaneously observed multiple CH$_{3}$OH transitions which have comparable Einstein A--coefficients (and therefore critical densities), but trace a broad range of upper--level--energy (E$_{up}$). These observations allow us to distinguish the hottest molecular gas directly associated with the UC H\\textsc{ii} region from the cooler envelope. The low E$_{up}$ and the modest critical density of the CS (1-0) transition provide sufficient optical depth, which allows us to trace the molecular gas to a much more extended region. However, in the case of G10.6-0.4, significant local structures can still be identified without being severely affected by the sidelobes of the dirty beam and the missing flux. The high optical depth of CS (1-0) also provides better signal--to--noise ratio to the geometrically thinner expansional signatures, and provides information for the on--going feedback processes. The observational results are presented in Section \\ref{chap_result}. In Section \\ref{chap_discussion}, we discuss the physical implication of the overall geometrical picture of the region. In addition, we estimate the feedback from various physical processes (radiative pressure, ionized gas pressure, stellar wind) and compare with our observational measurements. A brief summary of the results are provided in Section \\ref{chap_summary}. ", "conclusions": "\\label{chap_discussion} \\subsection{The Overall Picture of the O--Type Cluster Forming Region} \\label{chap_overall} We summarize our series of high resolution observations on G10.6--0.4 by a schematic model, which is shown in Figure \\ref{fig_shematic}. The observational results enable the discussions of the importance of the radiation and the pressure force of the ionized gas. In Sections \\ref{chap_radiation} and \\ref{chap_timescale}, we provide the order--of--magnitude estimates of the radiative pressure and the pressure forces of the ionized gas and the stellar wind. We discuss their implications for the dynamical evolution of the system in Section \\ref{chap_dp}. We are looking forward to comparing our results with numerical hydrodynamical simulations in the future. The dominant molecular structure in this model is a 0.5 pc scale massive envelope. The massive envelope rotates and contracts globally. A flattened overdensity formed in the mid--plane of the massive envelope (this is supported by the CS (1--0) observations, and also the $^{13}$CS (5--4) observations in Liu et al. 2010). Meanwhile, the molecular gas also contracts locally, and forms intermediate--mass or B--type massive (proto--)stars, and some low--mass (proto--)stars. The protostellar objects eject molecular outflows with high momentum and high energy into the ambient environment, which may affect the global contraction of the massive envelope (Li \\& Nakamura 2006; Nakamura \\& Li 2007; Carroll et al. 2009; Wang et al. 2010). Those molecular outflows can be revealed as highly blueshifted and redshifted $^{12}$CO emissions ($|v|$$\\ge$20 kms$^{-1}$) if the protostellar objects are not located in the mid--plane (Liu, Ho, \\& Zhang 2010). In the mid--plane, the outflows are impeded and decelerated by the high density ambient gas to have lower terminal velocities, which have to be diagnosed from the molecular outflow/shock tracers (Liu, Ho, \\& Zhang 2010). The regulated global contraction may lead to the formation of the O--type cluster in the center (Li \\& Nakamura 2006), which creates an UC H\\textsc{ii} region. The ionized gas and the ionizing photons leak out from the H\\textsc{ii} region mainly in the bipolar direction, either because of a configuration where the rotationally flattened system has a low gas density in the bipolar region, or because of a strong protostellar MHD wind which created the biconical cavities in that region. The outflowing molecular/ionized gas may flush through the ambient molecular gas on a timescale comparable to or much shorter than the global dynamical timescale\\footnote{The lower limit of the global dynamical timescale can be estimated by the free--fall timescale of $\\sim$10$^{5}$ years.}, and leave behind the elongated relics of interactions. This may explain the detected large filamentary structures by CS (1--0) in the northeast (Section \\ref{chap_csdis}). Due to the short expansional timescale, the filamentary structures were not yet relaxed by the global rotational/infall motion. The pressure of the ionized gas can also push the cavity--wall to a quick dynamical expansion (Section \\ref{chap_pv}). However, in the rotationally flattened system, most of the molecular mass are concentrated in the mid--plane. Owing to the self--gravitational or the Rayleigh--Taylor instabilities, the molecular gas becomes clumpy (Sollins \\& Ho 2005), and additionally reduces the effective solid angle as presented to the central stars. In such cases, the momentum impulse exerted by the ionized gas would be much smaller because of this smaller solid angle. The inertia of the molecular accretion flow in the mid--plane can easily overcome the impulse exerted by the stellar radiative pressure (also see Section \\ref{chap_discussion}) and the pressure of the ionized gas. This is consistent with the lack of a clear expansional motion in the PV diagram of the molecular lines (Figure \\ref{fig_ch3ohpv}). When the molecular gas gets sufficiently close to the central OB cluster, it is heated by the stellar radiation to a significantly higher temperature, and is further ionized to a temperature of 10$^{4}$ K (also see discussions in Keto 2002; Keto \\& Wood 2006). This overall scenario allows the O--type stars to continue accreting after the nuclear burning begins. \\subsection{The Role of the Biconical Cavity in Enhancing the Molecular Accretion} \\label{chap_radiation} The biconical cavity provides a low opacity channel for the photons to leak out of the envelope. This mechanism can significantly reduce the radiative pressure on the molecular gas, and can therefore enhance the molecular accretion flow. To quantitatively examine this effect, we first compare the relative importance of radiative pressure and the gravitational force in a system without the biconical cavity. Following the derivations in Jijina \\& Adams (1996), we assume the radiation field is nearly isotropic; and the stellar mass predominantly contributes to the gravitational force in the relevant scale. We assume the temperature distribution follows the power law $T(r)=T_{0}(r/r_{0})^{-1/2}$ (Keto, Ho, \\& Haschick 1987), and quote the resulting effective potential from Jijina \\& Adams (1996): \\begin{equation} V_{eff} = \\frac{GM}{r}\\left\\{ \\alpha r^{-1/2}-1\\right\\}, \\label{eq1} \\end{equation} where $M$ is the mass of the embedded stellar cluster. The parameter $\\alpha$ is defined by \\begin{equation} \\alpha\\equiv\\frac{L\\kappa_{P}(T_{0})}{6\\pi GMc}\\sqrt{r_{0}}, \\end{equation} where $L$ is the bolometric luminosity of the embedded stellar cluster, $\\kappa_{P}(T_{0})$ is the Planck mean opacity at the fiducial temperature $T_{0}$, and $r_{0}$ is the radius where $T=T_{0}$. The first term in equation \\ref{eq1} represents the radiative pressure, and the second term represents the gravitational force. From equation \\ref{eq1}, we see that when $\\alpha^{2}/r\\sim1$, the radiative pressure is comparably important with the gravity. The radiative pressure becomes dominant at smaller radii. In G10.6-0.4, the molecular hot toroid with temperature about 300 K is detected inward of the 0.05 pc radius\\footnote{The averaged temperature might be overestimated. A lower value of 87 K is reported by Beltr{\\'a}n et al. (2011) by observing optically thinner lines. We adopt the previous higher value to provide an upper limit of the radiation.}. To match the measured molecular gas temperature, we adopt\\footnote{We express the temperature power law with a parameter $T_{d}$, which is the dust sublimation temperature. If the temperature is higher than $T_{d}$, the dust is sublimated and the radiative pressure is reduced. The temperature power law of $T(r)=T_{0}(r/r_{0})^{-1/2}$ potentially overestimates the temperature at the 0.3 pc radius by a factor of 2 since the temperature decay faster with radius in regions with lower opacity, and has to be fitted with a more negative power law index. However, our discussions about the radiative pressure focus on the inner region which has high opacity.} $T_{0}=T_{d}\\sim$ 2300 K ($\\kappa(T_{0})$$\\sim$30 cm$^{2}$g$^{-1}$) and $r_{0} =$ 8.5$\\cdot$10$^{-4}$ pc. Given the G10.6-0.4 bolometric luminosity $L\\sim$10$^{6}$ L$_{\\odot}$ and stellar mass $M\\sim$200 M$_{\\odot}$, we obtain $\\alpha^{2}\\sim$0.07 pc (Ho \\& Haschick 1981; Keto 2002; Sollins \\& Ho 2005; Keto \\& Wood 2006). In such a case, the radiative pressure can significantly affect the dynamics at the scale of the hot toroid, or even reverse an inflow. However, based on the numerical radiation transfer calculations of a set of systems with different geometries, Krumholz, McKee and Klein (2005) suggested that the presence of the biconical outflow cavity allows the photons to leak out, and therefore potentially reduces the radiative pressure force by a factor of $\\sim$10. This effectively reduce $L$ by a factor of 10, and therefore reduce $\\alpha^2$ by a factor of 100. Assuming this conclusion can be generally applied, with the presence of the biconical cavity, the value of $\\alpha^{2}$ in the case of G10.6--0.4 might be effectively reduced to 7$\\cdot$10$^{-4}$ pc (140 AU), which is much smaller than the radius of the UC H\\textsc{ii} region in G10.6--0.4 ($\\sim$0.03 pc). Qualitatively this estimate is consistent with the case in Krumholz, McKee \\& Klein (2005), suggesting that without considering the ionization, the radiative pressure will only be important on the scale of a few hundred AU. We therefore suggest that in G10.6--0.4, the radiative pressure is unimportant for the dynamics of the molecular accretion flow. We note that we may overestimate the reduction factor of the radiative pressure owing to the non-self-consistent treatment of the temperature distribution. However, since $\\alpha^{2}$ scales as the square of the reduction factor, it can easily become smaller than the radius of the H\\textsc{ii} region with a small reduction factor. The presence of the biconical cavity can also lead to a temperature gradient steeper than our assumption, and may makes $\\alpha^{2}$ shrink to an even smaller radius. The dusty disk around the stars may redirect the radiation (Yorke \\& Sonnhalter 2002; Kuiper et al. 2010), and potentially play the role of reducing the radiative pressure on the accretion flow. However, it is still uncertain whether the disks can stably exist in UC H\\textsc{ii} regions, where multiple O--type stars are embedded. This has to be examined in future observations. \\subsection{The Ionized Gas Pressure and the Stellar Wind Feedback} \\label{chap_timescale} The large molecular mass in the contracting massive envelope contains enormous inertia. From the PV diagrams (Figure \\ref{fig_ch3ohpv}), we find that the rotational velocity at the 0.25 pc radius is about 4--5 kms$^{-1}$. This rotational motion is gravitationally bound by an enclosed mass of 1000--1600 M$_{\\odot}$. While the embedded OB stars contribute $\\sim$200 M$_{\\odot}$, the rest of the binding mass is dominantly contributed by the molecular gas. For a spherical geometry, this molecular mass corresponds to a mean molecular density $\\bar{n_{H_{2}}}$ of (3--5)$\\cdot$10$^{5}$ cm$^{-3}$. The momentum flux in the molecular gas flow can be estimated by $\\mu$$\\cdot$$\\bar{n_{H_{2}}}$$\\cdot$$v^{2}$, where $v$ is the infall velocity of the molecular gas, which can be larger than 1 kms$^{-1}$ (Ho \\& Haschick 1986; Keto, Ho \\& Haschick 1987, 1988; Keto 1990; Klaassen \\& Wilson 2008), and $\\mu$ is the mean molecular weight. In the embedded UC H\\textsc{ii} region, the momentum flux in the ionized gas can be estimated by $\\mu_{i}$$\\cdot$$n_{i}$$\\cdot$$v_{i}^{2}$, where $\\mu_{i}$ is the ion mass, $n_{i}$ is the ion density, and $v_{i}^{2}$ is the thermal velocity of the ions, which has the order of magnitude of 10 kms$^{-1}$. We assume that the ion density $n_{i}$ is equal to the electron density $n_{e}$, which is constrained to be 10$^{3}$--10$^{4}$ in previous observations (Ho and Haschick 1981; for a more sophisticated model fitting see Keto, Zhang, \\& Kurtz 2008). Assuming $\\mu$$\\sim$2$\\mu_{i}$, we find that the momentum flux in the molecular gas flow (6--10$\\cdot$10$^{5}$$\\cdot$$\\mu_{i}$ M$_{\\odot}$kms$^{-1}$cm$^{-2}$s$^{-1}$) and the momentum flux in the ionized gas (1--10$\\cdot$10$^{5}$$\\cdot$$\\mu_{i}$ M$_{\\odot}$kms$^{-1}$cm$^{-2}$s$^{-1}$) have the same order of magnitudes. This suggests that the two components are in an approximate dynamical equilibrium. In the real case, the mass is anisotropically distributed. The dynamics of the accretion flow in the denser region will not be significantly affected by the pressure of the ionized gas, while the lower density region can undergo a dynamical expansion powered by the pressure of the ionized gas. We estimate the momentum budget of the ionized gas and the stellar wind in this section. In the next section, we will suggest a simple geometric model for the density distribution, and estimate the velocity of the dynamical expansion, to compare with the observations. The embedded massive stars in UC H\\textsc{ii} regions B and C have similar spectral type (Ho \\& Haschick 1981) with the embedded OB stars in UC H\\textsc{ii} region A. Some of the energetic relations estimated in UC H\\textsc{ii} region A can also be applied in UC H\\textsc{ii} region B and C. \\paragraph{Expansional Time Scale} The expansional signatures of the biconical cavity have a characteristic radius $r_{c}$ of about 5$''$. Assuming the fastest expanding front of the cavity--wall uniformly expands with the terminal velocity $v_{e}$ of 5 kms$^{-1}$, the total expanding time $t_{e}$ is about 3$\\cdot$10$^{4}$ years, which is much shorter than the global dynamical timescale ($>$10$^{5}$ years). The timescale estimation constrains the total momentum budget from the stellar wind and the ionized gas. \\paragraph{The Stellar Wind Feedback} The UC H\\textsc{ii} region A contains a few O6--O9 stars with a total stellar mass M$_{*}$ of $\\sim$200 M$_{\\odot}$. The stellar wind from each of these massive stars has a mass loss rate of the order of magnitude of 10$^{-6}$ M$_{\\odot}$yr$^{-1}$, and terminal velocity of $\\sim$2000 kms$^{-1}$ (Tout et al. 1996). Given the number of the embedded O stars $N_{*}$ ($\\sim$4; Ho \\& Haschick 1981), the wind momentum feedback rate of the massive cluster $f_{w}$ is $N_{*}$$\\cdot$2$\\cdot$10$^{-3}$ M$_{\\odot}$kms$^{-1}$yr$^{-1}$. Assuming the embedded stellar cluster starts to feed back the stellar wind by the beginning of the expansion of the biconical cavity, the total wind momentum budget $I_{w}$ is $N_{*}$$\\cdot$60 M$_{\\odot}$kms$^{-1}$ isotropically spread in the entire 4$\\pi$ solid angle. The exerted pressure force has a $r^{-2}$ radial dependence. \\paragraph{The Ionized Gas Pressure} We adopt the characteristic radius $r$ of 0.15 pc, and adopt the electron density $n_{e}$ = 10$^{3}$--10$^{4}$ cm$^{-3}$ according to the measurements in Ho \\& Haschick (1981), to estimate the thermal pressure force of the ionized gas. Assuming the ion density equals to the electron density, and has the thermal velocity of $\\sim$10 kms$^{-1}$, the momentum feedback rate $f_{i}$ can be estimated by $\\Omega\\cdot r^{2}\\cdot(\\rho\\cdot v\\cdot v)$, where $\\Omega$ is the characteristic solid angle. As an order of magnitude estimation, we adopt $\\Omega$=4$\\pi$, which leads to $f_{i}$$\\sim$6$\\cdot$10$^{-4}$--6$\\cdot$10$^{-3}$ M$_{\\odot}$kms$^{-1}$yr$^{-1}$. The total ionized gas momentum feedback $I_{i}=\\int_{0}^{t_{e}} f_{i}(r(t)) dt$ depends on the expansion history, and has the order of magnitude of 10$^{1}$--10$^{2}$ M$_{\\odot}$kms$^{-1}$. The ionized gas feedback differs from the stellar wind feedback in the sense that it is a pressure effect (see also Figure \\ref{fig_shematic}), which does not exert force in the radial direction, but in the direction perpendicular to the cavity--wall. In addition, the exerted pressure force depends only on the temperature and density of the ion, which is determined by the ionization, the recombination, and the pressure balance. \\subsection{The Dynamical Process} \\label{chap_dp} From the observations of the 1.3 mm continuum emission, and the measurements of the rotational velocity of the molecular gas, we constrained the molecular mass inward of the 0.15 pc radius around the UC H\\textsc{ii} region A to have the order of magnitude of 400 M$_{\\odot}$. We assume a simple geometry in this region, and perform order--of--magnitude estimates of the energetic relation. \\subsubsection{The Geometrical Model} Assume the molecular gas has an initial mass density distribution $\\rho_{0}$ before the creation and the dynamical expansion of the biconical cavity. Defining $\\eta$ to be the angular separation from the plane of rotation, if the system is approximately axisymmetric, $\\rho_{0}$ can be represented by $\\rho_{0}\\equiv\\rho_{0}(r, \\eta)$. If $\\rho_{0}(r)$ can be represented as a polynomial of $\\cos(\\eta)$, for a rotationally flattened system, we expect the leading order dependence to be $\\cos(\\eta)$. This suggests that about 70\\% of molecular mass (280 M$_{\\odot}$) are concentrated in the region with $|\\eta|<$45$^{o}$. After the creation and the dynamical expansion of the biconical cavity, the molecular mass initially distributed in the region $|\\eta|>$45$^{o}$ (120 M$_{\\odot}$) are accumulated on the geometrically thin cavity--wall, which has an opening angle of 90$^{o}$. The assumption of the mass accumulation on the cavity--wall is consistent with the fast expansion (see Section \\ref{chap_timescale} for the estimate of timescale). The molecular gas initially distributed in the region $|\\eta|<$45$^{o}$ forms a dense flattened structure with density distribution $\\rho$. We argue that the opening angle of 90$^{o}$ is a reasonable value while comparing with the geometrical picture revealed by the CH$_{3}$OH images (Figure \\ref{fig_nh3}, \\ref{fig_ionout}). The cavity--wall has two components, W$_{\\perp}$ and W$_{\\parallel}$, of which the surface area are perpendicular and parallel to the radial direction, respectively. Given the opening angle, the ratio of the surface area of these two components is approximately 1; we assume the ratio of the mass accumulated on these two components is also 1. \\subsubsection{The Dynamical Expansion of Cavity Wall} \\paragraph{W$_{\\parallel}$} This component has a cone shape, with the total mass of $\\sim$60 M$_{\\odot}$. For it to expand at 2.5--5 kms$^{-1}$, it requires a momentum of 150--300 M$_{\\odot}$kms$^{-1}$. From the estimations in the previous section, we see that this required momentum has the same order of magnitude of the feedback from the ionized gas pressure. The gravitational force acts in the radial direction and does not retard the expansion in the $\\hat{\\theta}$ direction. We expect the stellar wind to have weak effects on the dynamics of W$_{\\parallel}$ since the wind is parallel to the cavity--wall. \\paragraph{W$_{\\perp}$} This part is directly driven by the stellar radiation and the stellar wind, which are competing with gravity in the radial direction. At the radius of 0.15 pc, the stellar wind and the ionized gas pressure can have comparable effect on W$_{\\perp}$. Assuming the mass is 60 M$_{\\odot}$, the gravitational force has the order of magnitude of 2$\\cdot$10$^{-3}$ M$_{\\odot}$kms$^{-1}$yr$^{-1}$, which is also comparable to the force of the stellar wind and the ionized gas. Observationally this part of cavity--wall is not as prominently detected as W$_{\\parallel}$. Since the embedded stellar cluster in the UC H\\textsc{ii} region A has a high ionizing photon emission rate ($S$ = 10$^{49}$ s$^{-1}$; Ho \\& Haschick 1981), which is able to ionize on the order of 10$^{3}$ M$_{\\odot}$ of gas in the biconical region. Even if the recombination rate marginally balances the ionization rate, the stellar ionization can still explain why the W$_{\\perp}$ is not clearly detected. Some local high density clumps which have higher recombination rate may be self-shielded from the stellar ionization and have long lifetime. Those high density clumps occupy small solid angles, and are less accelerated by the stellar wind and the pressure of the ionized gas. Some of those clumps can still fall toward the OB cluster owing to the gravitational attraction. The clumpy 3.6 cm free-free continuum emissions in the UC H\\textsc{ii} region A, especially, northeastern to the emission peak, may be explained by those clumps externally ionized by the central OB cluster. We note that we may overestimate the dynamical expansion timescale of the biconical cavity since the cavity can start expansion from a finite size owing to the initial ionization, which leads to the overestimation of the momentum feedback. However, those ionized gases are not gravitationally bound and are free to leak out. Therefore, the accumulated mass on the cavity--walls W$_{\\parallel}$ and W$_{\\perp}$ is reduced if the initial ionization is efficient, and requires smaller momentum feedback to reach the final state velocity. Similar arguments are valid if the biconical cavity is initially created by the massive bipolar outflow (an example see G240.31+0.07: Qiu et al. 2009) and the molecular gas is evacuated. \\paragraph{The Dense Flattened Structure} This structure has the mass of $\\sim$280 M$_{\\odot}$, and continues from the outer of $\\sim$0.15 pc radius to the inner $\\sim$0.03 pc (Liu et al. 2010). The total gravitational force on this structure depends on the embedded mass. A lower limit of 9.6$\\cdot$10$^{-3}$ M$_{\\odot}$kms$^{-1}$yr$^{-1}$ can be given assuming that the majority of mass is distributed at the 0.15 pc radius. This lower limit is already on the same order of magnitude as the feedback from the stellar wind and the ionized gas pressure; the actual value of the gravitational force should be much larger. Depending on the effective solid angle it occupies, the stellar wind and the ionized gas pressure may have negligible effect on the radial motion of the flattened structure. With the biconical cavity structure which leads to the photon leakage, the radiative pressure is also negligible. We therefore suggest that the dynamics in this region is dominated by gravity and rotation. \\paragraph{UC HII region B and C} The stellar wind and ionized gas pressure feedback in UC H\\textsc{ii} region B and C have the same order of magnitude as those in the UC H\\textsc{ii} region A. However, these two regions have lower initial gas density, which leads to a smaller accumulated mass on the expansional shell. The observed expansional velocities in these two regions are comparable to but smaller than the thermal velocity of the ionized gas, suggesting that the ionized gas pressure is an important driving source of the dynamical expansion. From the estimations in the previous section, the stellar wind feedback can be comparably important, depending on the detailed geometry of these system. We present high resolution observations of molecular lines and free-free continuum for the UC H\\textsc{ii} region G10.6--0.4. The resolved projected distributions of the molecular gas and the ionized gas suggest an overall picture consisting of an extended ($\\sim$0.5 pc) envelope, a single compact ($\\sim$0.1 pc) hot rotating toroid, and a biconical molecular cavity filled with ionized gas. This overall geometry resembles the standard envelope--disk and protostellar outflow model for the low--mass star forming region. With the presence of the biconical cavity, we suggest that the radiative pressure can be significantly reduced.% In the plane of rotation, at the scale larger than 1$''$ (0.03 pc), we see that the rotational motion of the dense gas is not severely disturbed, which consistently suggests that the radiation is not yet important for the molecular accretion flow. The stellar radiation may increase and play a more important role as the stellar mass is increased. From the observations of CS (1--0), we suggest that the biconical cavity around the UC H\\textsc{ii} region A is undergoing an expansional motion, with velocity of 2.5--5 kms$^{-1}$. We perform simple order of magnitude estimates, and suggest that the feedback from the ionized gas pressure can account for the required momentum of this dynamical expansion. The expansional signatures are also detected in the UC H\\textsc{ii} region B and C, which are driven by the ionized gas pressure and the stellar wind. The expansional motions of the UC H\\textsc{ii} regions significantly disturb the local dynamics of the molecular gas, inject the energy, and may induce the non-uniformity in the molecular accretion flow, which is an important feedback mechanism in the massive molecular clump." }, "1101/1101.4086_arXiv.txt": { "abstract": "The availability of vector-magnetogram sequences with sufficient accuracy and cadence to estimate the temporal derivative of the magnetic field allows us to use Faraday's law to find an approximate solution for the electric field in the photosphere, using a Poloidal--Toroidal Decomposition (PTD) of the magnetic field and its partial time derivative. Without additional information, however, the electric field found from this technique is under-determined -- Faraday's law provides no information about the electric field that can be derived the gradient of a scalar potential. Here, we show how additional information in the form of line-of-sight Doppler-flow measurements, and motions transverse to the line-of-sight determined with \\textit{ad-hoc} methods such as local correlation tracking, can be combined with the PTD solutions to provide much more accurate solutions for the solar electric field, and therefore the Poynting flux of electromagnetic energy in the solar photosphere. Reliable, accurate maps of the Poynting flux are essential for quantitative studies of the buildup of magnetic energy before flares and coronal mass ejections. ", "introduction": "\\label{Introduction} The launch of SDO, with its ability to measure the Sun's vector magnetic field anywhere on the disk with a high temporal cadence, promises to usher in a new era of solar astronomy. This new era of measurement demands new approaches for the analysis and use of this data. We show in this article how the vector magnetic field and Doppler-flow measurements that can now be made with HMI (\\opencite{Scherrer2005}) lead to new methods for determining the electric field vector, and the Poynting Flux vector \\be \\vecS = {1 \\over 4 \\pi} c \\vecE \\times \\vecB \\label{equation:sdef} \\ee at the solar photosphere. The Poynting flux measures the flow of electromagnetic energy at the layers where the magnetic field is determined. Quantitative observational studies of how energy flows into the corona depend on deriving accurate estimates of the Poynting flux. Most work estimating the Sun's electric field or Poynting flux either explicitly or implicitly assumes that the electric field is determined by ideal MHD processes, and therefore the problem can be reduced to determining a velocity field associated with the observed magnetic-field evolution. One class of velocity estimation techniques are ``Local Correlation Tracking'' (LCT) methods, which essentially capture pattern motions of the line-of-sight magnetic field or white-light intensity. This approach was pioneered by \\inlinecite{November1988}. Other implementations include the Lockheed--Martin LCT code (\\opencite{Title1995}; \\opencite{Hurlburt1995}), ``Balltracking'' (\\opencite{Potts2004}), and the FLCT code (\\opencite{Fisher2008}). Another class of velocity-estimation methods incorporate solutions of the vertical component of the magnetic induction equation into determinations of the velocity field (\\opencite{Kusano2002}; \\opencite{Welsch2004}; \\opencite{Longcope2004}; Schuck, 2006, 2008; \\opencite{Chae2008}). The work we present in this article incorporates solutions of the three-dimensional magnetic induction equation, using the electric field as the fundamental variable, rather than the velocity field. The temporal evolution of the Sun's magnetic field is governed by Faraday's law, \\be { \\partial \\vecB \\over \\partial t} = - \\grad \\times c \\vecE\\ . \\label{equation:faraday} \\ee If one can can make a map on the photosphere of $\\partial \\vecB / \\partial t$, can one determine $\\vecE$ by uncurling this equation? Addressing this question was the focus of \\inlinecite{Fisher2010a}, in which a poloidal--toroidal decomposition (PTD) of the temporal derivative of the magnetic field was used to invert Faraday's law to find $\\vecE$. \\inlinecite{Fisher2010a} found that one could indeed find solutions for $\\vecE$ that solve all three components of Faraday's law, but the solutions are not unique: the gradient of a scalar function can be added to the PTD solutions for $\\vecE$ without affecting $\\grad \\times \\vecE$. \\inlinecite{Fisher2010a} explored two different methods for determining the scalar function using \\textit{ad-hoc} and variational methods, both of which enforced the assumption, from ideal MHD, that $\\vecE$ must be normal to $\\vecB$. Unfortunately, the agreement with a test case from an MHD solution, while better than conventional correlation-tracking methods, was still disappointing. The authors concluded that including additional information from other observed data was one possible approach for improving the electric field inversions. In this article, we use the same MHD simulation test case used in \\inlinecite{Welsch2007} and \\inlinecite{Fisher2010a} to show that using Doppler-flow measurements to determine the electric scalar potential, especially in regions where the magnetic field is primarily horizontal, can dramatically improve the inversion for the electric field and the Poynting flux. In Section \\ref{section:ptd} we review the PTD formalism that describes how one can derive the purely inductive part of the electric field from measurements that estimate the time derivative of $\\vecB$, and the technique of Section 3.2 of \\inlinecite{Fisher2010a}, showing how one can derive a potential electric field, which, when added to the inductive part of the electric field, is normal to the magnetic field. This is useful in generating electric-field solutions that are both consistent with Faraday's law and with ideal MHD, which is generally believed to be a good approximation in the solar photosphere. Section \\ref{section:doppler} argues from physical grounds why magnetic-flux emergence may make a large contribution to the part of the electric field attributable to a potential function. Then, starting from this argument, we derive a Poisson equation for an electric-field potential function that is determined primarily from knowledge of the vertical velocity field, as determined from Doppler measurements, and the horizontal magnetic field near polarity inversion lines where the field is nearly horizontal. The electric field from this contribution is then added to that determined from the PTD solutions. We then apply this technique to the MHD simulation test data, to compare the electric field from the simulation with that from PTD alone, and with that from combining PTD with Doppler measurements. In Section \\ref{section:flct}, we try a similar approach, but instead of using contributions to the horizontal electric field from Doppler measurements, we use non-inductive contributions to the electric field determined from the FLCT correlation-tracking technique, applicable in regions where the magnetic field is mainly vertical. This technique is essentially the three-dimensional analogue of the ILCT technique described by \\inlinecite{Welsch2004}. We also try combining PTD with contributions from both the Doppler measurements and those from FLCT, and compare with the simulation data. Our results are summarized in Section \\ref{section:conclusions}, along with a discussion of where additional work is needed. ", "conclusions": "\\label{section:conclusions} We have reviewed how the PTD solutions of Faraday's law for $\\vecE$ can be found using temporal sequences of vector magnetograms that can be obtained with the HMI instrument on NASA's SDO mission. We discussed why these solutions are under-determined, and the importance of determining the contributions to the electric field that can be derived from a scalar potential. We demonstrate, using simulation data where the true electric field is known, that knowledge of the vertical-velocity field (obtainable by Doppler measurements) can provide important information about the electric field. When this information is combined with the PTD solutions of Faraday's law, dramatically more accurate recovery of the true electric field is possible. We find that additional information about flows from local correlation-tracking methods can also be combined with the PTD solutions, but the additional information is signficantly less important than that from the Doppler measurements. We are able to quantitatively reconstruct the electromagnetic Poynting flux in the simulations by using our combination of the PTD solutions and those from Doppler measurements. This ``proof-of-concept'' demonstration argues strongly for the development of electric-field and Poynting-flux tools to be used routinely in the analysis of HMI vector magnetic-field measurements. Routinely available Poynting-flux maps will be useful for scientific studies of flare-energy buildup, understanding the flow of magnetic energy in the solar atmosphere prior to CME initiation, and will aid in understanding the flow of energy that heats the corona. Further, the PTD formalism for the magnetic field itself (Equation (\\ref{equation:ptdstatichz})) allows for a straight-forward decomposition of the Poynting flux into changes in the potential-field energy, and the flux of free magnetic energy (see \\opencite{Welsch2006} and the end of Section 2.1 of \\opencite{Fisher2010a}). The flux of free magnetic energy is especially important in determining how energy builds up in flare-productive active regions. To find solutions for $\\vecE$ and the Poynting flux $\\vecS$ using the PTD formalism plus Doppler measurements requires only the solution of three two-dimensional Poisson equations. While real vector-magnetogram patches will not have periodic boundary conditions (as were employed in this article), straightforward numerical techniques exist to solve these equations routinely. Preliminary investigations also indicate that generalizing the PTD solutions and Doppler measurements to cases of non-normal viewing angle will be straightforward. In our opinion, the major obstacle that remains before such solutions can be routinely applied to the HMI data, is a detailed understanding of how measurement errors and disambiguation errors in the vector magnetograms will affect the solutions, and how the effects of these errors are best ameliorated. \\begin{acks} This research was funded by the NASA Heliophysics Theory Program (grant NNX08AI56G), the NASA Living-With-a-Star TR\\&T Program (grant NNX08AQ30G), by the NSF SHINE program (grants ATM0551084 and ATM0752597), and support from NSF's AGS Program (grant ATM0641303) for our participation in the University of Michigan's CCHM Project. The authors are grateful to US taxpayers for providing the funds necessary to perform this work. The authors wish to acknowledge Dick Canfield for his pioneering work in the use of vector magnetograms in solar physics. The inspiration for the work described here can be traced to a Solar MURI workshop held at UC Berkeley in 2002, in which Dick Canfield played a major role in defining long-term research goals for the use of vector magnetograms in quantitative models of the Sun's atmosphere. \\end{acks}" }, "1101/1101.4072.txt": { "abstract": "We mapped the MSX dark cloud G084.81$-$01.09 in the NH$_3$~(1,1) - (4,4) lines and in the $J$ = 1-0 transitions of $^{12}$CO, $^{13}$CO, C$^{18}$O and HCO$^+$ in order to study the physical properties of infrared dark clouds, and to better understand the initial conditions for massive star formation. Six ammonia cores are identified with masses ranging from 60 to 250~$M_\\sun$, a kinetic temperature of 12~K, and a molecular hydrogen number density $n({\\rm H_2})\\sim10^5$~cm$^{-3}$. In our high mass cores, the ammonia line width of 1~km~s$^{-1}$ is larger than those found in lower mass cores but narrower than the more evolved massive ones. We detected self-reversed profiles in HCO$^+$ across the northern part of our cloud and velocity gradients in different molecules. These indicate an expanding motion in the outer layer and more complex motions of the clumps more inside our cloud. We also discuss the millimeter wave continuum from the dust. These properties indicate that our cloud is a potential site of massive star formation but is still in a very early evolutionary stage. ", "introduction": "Infrared dark clouds (IRDCs) are identified as small areas with high extinction against the bright diffuse mid-infrared background along the Galactic Plane. These ``flecks'' in the sky were first discovered by the 15~$\\mu$m ISOGAL survey (the inner Galactic disk survey of the Infrared Space Observatory, \\citealt{per96}) and later cataloged and studied using the 8.3~$\\mu$m survey data of the Mid-course Space Experiment (MSX, \\citealt{ega98}). \\citeauthor{ega98} found about 2000 IRDCs and note that these objects are cold and dense molecular cores. These results were refined by later molecular line observations \\citep{car98,tey02} and radio continuum observations \\citep{car00}. These authors concluded that these clouds have low kinetic temperatures of 10-20~K and a gas density over $10^5$~cm$^{-3}$. Their small line widths of 1 - 3~km~s$^{-1}$ can be related to a very early phase of star formation. Therefore, the IRDCs are good candidates for pre-protostellar cores and may be the primary sites for future massive star formation, which plays an important role in the evolution of our galaxy. However, IRDCs remain mysterious in at least some aspects. A controversial issue being discussed is whether the cloud evolution is quasi-static or whether it is undergoing a dynamical process caused by lack of equilibrium \\citep{ber07}. Understanding the role of supersonic turbulence in dark clouds is important in settling this argument. Evidence from current observations of IRDCs seems to support both views. Most IRDC cores are supported by turbulent pressure but appear to be virialised \\citep{pil06}. Furthermore, the origin of the turbulence in the molecular clouds and the transition from the turbulent cloud to the quiescent core regime is still unclear. Also, more systematic studies need to be done to compare them with other relatively nearby massive star-forming clouds. The answers will provide a critical element in our general understanding of star formation. In this study, we have obtained spectroscopic data in several molecules towards the dark cloud G048.81--01.09, a source in the catalog of \\citet{sim06}. The cloud has an elongated morphology (N-S) with an extent of $\\sim 15$ arcmin, and is situated in a region of high extinction ($A_v > 20$, \\citealt{cam02}) of the dark cloud LDN935 \\citep{lyn62} which separates the W80 HII region into the North America (NGC 7000) and Pelican (IC 5070) nebulae \\citep{com05}. \\citet{com05} proposed that the ionizing source of the complex is 2MASS~J205551.25+435224.6, an O5V star located behind LDN935. The IRDC is almost starless with no associated IRAS or MSX point sources. Previous molecular line observations indicated the presence of large amounts of gas with different velocity components and signs of ongoing star formation in LDN935 around our dark cloud \\citep{fel93}. A single pointed CO observation towards our cloud conducted by \\citet{mil75} suggests however that the gas is not appreciably heated. There is no direct distance measurement for G084.81$-$01.09, but it is associated with dark cloud LDN935 for which there are a number of distance estimates. A radio continuum study carried out by \\citet{wen83} gave a distance estimate of 500~pc. \\citet{str93} obtained a consistent result of 580~pc on the basis of photometric results of 564 stars toward three areas near the North America and Pelican Nebulae complexes. New measurements done by \\citet{cer07} derived a distance of about 0.7~kpc, which was obtained from radio recombination line observations at a frequency around 1.4~GHz and by applying a flat rotation-curve model. \\citet{com05} obtained a distance of 610~pc to the ionizing star, 2MASS~J205551.25+435224.6. Here we adopt a distance of 580~pc in our calculation of physical parameters. ", "conclusions": "\\label{sec:ana} \\subsection{Evidence of massive star formation} In this section, we will compare our results with those from other low-mass star forming regions, and discuss the evidence for massive star formation and the evolutionary state of the cloud. The mean intrinsic line width for NH$_3$~(1,1) in the samples given by \\citet{jij99} (most of the objects are low-mass cores) is 0.5~km~s$^{-1}$. A recent work by \\citet{fos09} also reported a small observed line width of 0.3~km~s$^{-1}$ for low mass NH$_3$ cores in the Perseus Molecular cloud. \\citet{cra05} undertook a survey toward low-mass starless cores and also reported a narrow line width of 0.5~km~s$^{-1}$ for C$^{18}$O~(1-0). A broader line width of 1.2~km~s$^{-1}$ for C$^{18}$O~(1-0) is derived from a JCMT observation by \\citet{jor02}, which is still less than those given in our observed region. Meanwhile, the dominant non-thermal line widths in our cloud indicate more non-thermal support than those in low-mass star forming regions. Although line widths can be broadened by active outflows in low-mass protostellar regions to 2~km~s$^{-1}$ for C$^{18}$O~(1-0) \\citep{swi08} and 1.2~km~s$^{-1}$ for NH$_3$~(1,1 ) \\citep{rud01}, at this stage, an active outflow has not developed yet. The NH$_3$ line widths here are comparable to those IRDCs of \\citet{pil06}. Moreover, our derived column densities of NH$_3$, a few times 10$^{15}$~cm$^{-2}$, are higher than those of the low-mass cores ($\\sim10^{13}-10^{15}$~cm$^{-2}$) given by \\citet{suz92}, and are also comparable to those of the massive IRDCs given by \\citet{pil06}. In addition, the supersonic motions of the envelope and the large velocity gradients in our cloud also indicate a more dynamic motion compared to low-mass cores. Therefore, we suggest that MSXDC G084.81$-$01.09 is potentially at an early evolutionary state of massive star formation. \\subsection{Gravitational stability of the cores}\\label{sec:sta} The expanding envelope detected from the line profiles described in \\S\\ref{sec:lin} lead us to consider the gravitational stability of MSXDC G084.81$-$01.09. By using the molecular mass, $M_{Molecular}$, derived from C$^{18}$O~(1-0) and the core radius $R$ determined from the NH$_3$~(1,1) maps, we obtain an escape velocity ($\\sqrt{2GM/R}$) which ranges from $\\sim 1.9$~km~s$^{-1}$ in core 6 to a maximum of $\\sim 3.2$~km~s$^{-1}$ in core 2 (Table~\\ref{tbl:core}). The three-dimensional velocity dispersion of the cores provided by the NH$_3$ line width lies in the range of 1.2 -- 2.4~km~s$^{-1}$ which is less than the escape velocity in each core. However, the C$^{18}$O velocity dispersions (2.5 - 4~km~s$^{-1}$) exceed the escape velocity in some regions of the cloud, especially for core 6, which has a small escape velocity. This may lead to instability for core 6. Additionally, an expansion speed of 1.0~km~s$^{-1}$ in the envelope of the core (cf. \\S\\ref{sec:lin}), though significantly less than the escape velocity, may aggravate the situation. To examine the gravitational stability of the cores, it is useful to calculate the virial mass of the cores as well as the virial parameter, the ratio between the virial mass and core mass. The virial parameter given by \\citet{ber92} can be expressed as $\\alpha=5\\sigma^2R/GM$, where R is the radius of the core, $\\sigma=\\sqrt{3/(8\\ln2)}\\times FWHM$, and $M$ is the core mass. The virial mass and virial parameter can be found in Table~\\ref{tbl:core}. The average virial parameter is about 1.1 which suggests that most of the cores are virialised. However, the virial parameters of cores 2 and 6 deviate significantly from unity. This can be caused by a number of factors. One explanation is the potential for the presence of spatially unresolved stellar clusters in the cores, which can be identified from the color temperature of the dust continuum map. For $M_{\\rm Virial}>M_{\\rm Molecular}$, beam filling factors smaller than unity or streaming motions may adversely affect the estimate and may lead to an underestimated column density or an overestimated line width. Alternatively, it is possible that the cores may not actually be in virial equilibrium, and may be transient entities \\citep{bal06}, which is probably the case in core 6. We derive a Jeans mass greater than 170~$M_\\sun$ for each core using the equation $M_{Jeans} = 17.3~{T_{kin}}^{1.5} n^{-0.5} M_\\sun$, where $T_{kin}$ is the kinetic temperature derived from NH$_3$ data, and $n$ is the average density as listed in Table.~\\ref{tbl:core}. We note that the masses of cores 1 and 2 are comparable to their Jeans masses, but the other cores are significantly less massive. However, considering the clumping in cores 1 and 2, it is still possible for them to form several individual protostars. Therefore, we suggest that cores 1 - 5 are gravitationally bound, among which cores 1 and 2 are marginally stable against collapse. However, core 6 is probably transient. \\subsection{Comparison with more evolved clouds} Our dark cloud has a lower kinetic temperature than other evolved clouds. \\citet{wu06} reported a mean temperature of 19~K in massive water maser sources excluding known HII regions, and in ultra-compact HII (UC~HII) regions. \\citet{chu90} give a significantly higher temperature spread from 15~K to $>$60~K for UC~HII regions. The temperature of the cloud is a good tracer to determine its evolutionary stage in the context of massive star formation. The low temperatures in our cloud suggests that it is not an active high-mass star forming region yet. The C$^{18}$O line widths are typically a factor of two to five times smaller than the line widths in more evolved massive star forming regions traced by other molecules with similar densities. The NH$_3$ line widths we find are comparable to the massive cores reported by \\citet{pil06} and \\citet{wu06}, but are smaller than those associated with water masers or are in proximity to UC~HII regions. More specifically, the line widths of our cloud are similar to those of the most quiescent NH$_3$ cores with no methanol maser or 24~GHz continuum emission described by \\citet{lon07} and smaller than those in more evolved cores with maser or continuum emission. Our cloud also presents similar NH$_3$ line width to those bright-rimmed clouds which are undergoing recently initiated star formation and being subjected to intense levels of ionizing radiation \\citep{mor10}. After removing the thermal broadening contributions to the line widths, the velocity dispersions associated with turbulence are supersonic and lie between those of the sources triggered by the radiatively driven implosion and the non-triggered ones. For the more evolved clouds in the UC~HII phase, \\citet{chu90} reported significantly larger line widths of $\\sim 3$~km~s$^{-1}$. The large line widths in more evolved stages are likely due to a combination of warmer temperatures and broadening from dynamics such as outflows. The HCO$^+$~(1$-$0) spectra associated with hyper-compact H~II regions have a considerably larger line width (15-60~km~s$^{-1}$) reported by \\citet{chu10}. For our cloud, our HCO$^+$ line profiles are similar to the samples of \\citet{bra01} and \\citet{ces99}, which are all in the pre-UC~HII stage of star formation, though some sources in \\citet{ces99} are proven to be more evolved than those in \\citet{bra01}. We find similar line widths and profiles as theirs, with a self-absorption dip between the blue and red peak. However, our cloud is colder than the pre-UC~HII sources and do not have any associated IRAS point source. The optical depths of NH$_3$~(1,1) are comparable to the cores given by \\citet{pil06} and \\citet{lon07}. Our core 4, 5 and 6 have a lower optical depth than the NH$_3$ sources ($\\sim 2.7 - 2.9$) associated with methanol masers or 24-GHz continuum emission in \\citet{lon07}, while core 1, 2 and 3 near the cloud center gives higher optical depth but still lower than those samples ($\\sim 3.1$) with only NH$_3$ association. However, due to the large uncertainty, without further observation, it would be premature to classify these sources into more detailed evolutionary stages based on optical depth. Our cloud density is about an order of magnitude lower than the typical value in more evolved massive star-forming regions: an average density of $10^6$~cm$^{-3}$ is given by \\citet{beu02} using LVG calculations in clouds prior to building up an UC~HII region, and a similar value is given by \\citet{pil07} in clouds harboring a UC~HII region. Our NH$_3$~(1,1) core sizes lie between the mean value of 0.28~pc given by \\citet{lon07} and 0.57~pc given by \\citet{pil06}, but are significantly smaller than the mean size of 1.6~pc reported by \\citet{wu06}. The small values of \\citet{lon07} probably result from their small beam size (11\\arcsec), but the large values of \\citet{wu06} are likely a result of evolution. All these aspects suggest that our cloud is in an early evolutionary stage of massive star formation. These quiescent cores are likely the candidates for pre-protostellar cores. \\subsection{Rotation in the cores}\\label{sec:rot} One of the possible generators of the velocity gradients mentioned in \\S\\ref{sec:vel} is rotation in the cores. \\citet{goo93} has discussed the relation between rotation and the geometric properties of the cores. They found no causal relation between velocity gradient direction and core elongation, nor any relationship between the magnitude of the gradient and core shape, on the size scale of $10^{17}$~cm. An examination of the gradients in our clouds also reveals no obvious relation of this kind. In view of the early evolutionary state and the dominant role of non-thermal line broadening in the cloud, we believe that the velocity gradients in the cloud are more affected by stochastic processes, such as fragmentation, collisions, and nonuniform magnetic fields rather than ordered motions such as rotation. We use the parameter $\\beta$ as defined by \\citet{goo93} to compare the rotational kinetic energy to the gravitational energy. Thus $\\beta$ can be written as \\begin{displaymath} \\beta={(1/2)I\\omega^2 \\over qGM^2/R}={1 \\over 2}{p \\over q}{\\omega^2R^3 \\over GM}, \\end{displaymath} where $I$ is the moment of inertia given by $I=pMR^2$, $qGM^2/R$ represents the gravitational potential energy of the mass $M$ within a radius $R$, and $\\omega=\\mathscr{G}/\\sin i$, where $i$ is the inclination of the core along the line of sight. We assume $p/q=0.22$ as for a sphere with an $r^{-2}$ density profile and $\\sin i=1$. Our $\\beta$ values as listed in Table~\\ref{tbl:vgradient} are consistent with the result of \\citet{goo93} that most clouds have $\\beta\\leq 0.05$. The small values of $\\beta$ show that the effect of rotation is not significant in maintaining the overall dynamical stability for the cloud. The results also indicate that these cores are unlikely to experience instabilities driven by rotation (e.g., bars, fission, or rings), if the magnetic fields are not taken into account." }, "1101/1101.5878_arXiv.txt": { "abstract": "Recently, the Fermi-LAT collaboration reported upper limits on the GeV gamma-ray flux from nearby clusters of galaxies. Motivated by these limits, we study corresponding constraints on gamma-ray emissions from two specific decaying dark matter models, one via grand unification scale suppressed operators and the other via R-parity violating operators. Both can account for the PAMELA and Fermi-LAT excesses of $e^{\\pm}$. For GUT decaying dark matter, the gamma-rays from the M49 and Fornax clusters, with energy in the range of $1$ to $10$ GeV, lead to the most stringent constraints to date. As a result, this dark matter is disfavored with conventional model of $e^{\\pm}$ background. In addition, it is likely that some tension exists between the Fermi-LAT $e^\\pm$ excess and the gamma-ray constraints for any decaying dark matter model, provided conventional model of $e^{\\pm}$ background is adopted. Nevertheless, the GUT decaying dark matter can still solely account for the PAMELA positron fraction excess without violating the gamma-ray constraints. For the gravitino dark matter model with R-parity violation, cluster observations do not give tight constraints. This is because a different $e^{\\pm}$ background has been adopted which leads to relatively light dark matter mass around 200 GeV. ", "introduction": "It is now believed that the dominant matter in the universe should be non-baryonic dark matter (DM) instead of visible ones. In addition, DM should not be composed of any known Standard Model (SM) particles. There are three experimental avenues to prove the existence of DM: direct detection, indirect detection and collider searches. In this paper, we will focus on the indirect detection of DM through gamma-rays and electrons/positrons in cosmic-rays. The Pamela collaboration has reported significant excess in the positron fraction between $1$ and $100$ GeV \\cite{Adriani:2008zr}. Later on the Fermi-LAT collaboration also observed a clear feature in the spectrum of electrons and positrons from $20$ GeV to $1$ TeV \\cite{Abdo:2009zk}, which is harder than predictions of conventional models of cosmic ray propagation. These excesses of electrons and positrons could be attributed by the annihilation or decay of DM. If so, one might also detect gamma-rays from these DM annihilation/decay. Recently, the Fermi-LAT collaboration has reported measurements of the GeV gamma-ray from nearby clusters of galaxies \\cite{Ackermann:2010qj}. In an observation period of 18-month, there was no direct observational evidence. The upper limits on gamma-ray flux are given with 95\\% confidence-level. As clusters of galaxies are DM dominated and DM annihilation/decay would almost inevitably emit gamma-rays, these upper limits may lead to stringent constraints on DM model parameters. Notice that based on observations of the first 11-month of Fermi-LAT, there already existed some discussions about DM models along this way \\cite{Ackermann:2010rg,Yuan:2010gn,Dugger:2010ys}. To interpret the PAMELA and Fermi-LAT excesses in terms of DM annihilation, a large boost factor of order $100-1000$ is required for the theory to be consistent with the relic abundance measured by the WMAP \\cite{Bergstrom:2008gr,Cholis:2008hb,Cirelli:2008pk}. Moreover, gamma-ray fluxes produced by the DM annihilation are anticipated to be large at locations with very high density of DM, because they are proportional to the local DM density squared. There are plenty of experimental measurements of gamma-rays from inside/outside the Galactic halo. Accordingly, the annihilating DM scenarios have been strongly constrained by the observations \\cite{Nardi:2008ix,Bertone:2008xr,Cirelli:2009dv,Papucci:2009gd,Ackermann:2010rg,Yuan:2010gn}. So it is not easy to build models of annihilating DM at TeV scale satisfying all constraints.\\footnote{There are certainly ways out of this problem, see for example a recent attempt in \\cite{Liu:2011aa}% } Decaying DM offers an alternative explanation to the $e^{\\pm}$ excesses. In decaying DM scenarios, constraints from gamma-ray fluxes are relatively easier to be satisfied since the gamma-ray fluxes are linearly proportional to the local DM density. However to fit the PAMELA and Fermi-LAT excesses, a lifetime of the order $10^{26}$'s for DM is needed, which is even much longer than that of the Universe. From the perspective of particle physics, there are several frameworks which can naturally provide such a long lifetime. For instance, DM may decay via operators suppressed by the grand unification theory (GUT) scale $10^{16}$ GeV, which makes it sufficiently long-living. DM may also decay via a very weak R-parity violation process: R-parity conservation makes the lightest supersymmetric particle (LSP) stable to be DM candidate while R-parity violating operators make DM decay sufficiently slow. There are of course other possibilities, for example DM may decay via instanton-induced operators \\cite{Carone:2010ha}. Recently it has also been shown that such a lifetime can arise naturally from the goldstino decay \\cite{Cheng:2010mw}. For decaying DM with $\\mu^{+}\\mu^{-}$ and $b\\bar{b}$ final states, the impact of the Fermi 11-month observation of gamma-rays from nearby clusters has been discussed in \\cite{Dugger:2010ys}. It was found that these gamma-rays from clusters considerably improves previous constraints on the lifetime and mass of the DM. Especially, the Fornax cluster provides the strongest constraint to date. This motivated us to investigate similar constraints on some (arguably) theoretically better-motivated models of decaying DM. Specifically, we will focus on two different decaying DM scenarios mentioned above: decaying via GUT scale suppressed operators or decaying via R-parity violating operators. As shown before, each scenario could give a reasonable fit to the $e^{\\pm}$ excesses. In addition, both scenarios are consistent with measurements of Galactic and extragalactic gamma-rays. We expect that the gamma-rays from nearby clusters should lead to more stringent constraints than before. This paper is organized as follows. We discuss the gamma-ray fluxes from galaxy clusters in general in Section II. In Section III, DM gamma-ray signals from nearby clusters are discussed in a GUT framework. Section IV is devoted to the study of a decaying DM scenario with R-parity violating operators. We conclude with a summary in section V. ", "conclusions": "In this paper, we have studied the gamma-ray signals from six nearby clusters of galaxies in decaying DM scenarios. Specifically, we have concentrated on two decaying DM models which are physically well motivated. One is decaying DM in the framework of GUT and the other is decaying DM with R-parity violation. As discussed in previous works, both of them are able to account for the observational $e^{\\pm}$ excesses and consistent with various measurements of gamma-rays from inside/outside the Galactic halo. Recently, the Fermi-LAT collaboration has reported new upper limits on GeV gamma-ray fluxes from clusters of galaxies. We use these limits to further constrain decaying DM scenarios. For GUT decaying DM, we find that the DM induced gamma-ray signals are comparable in magnitude to the experimental upper limits for all six clusters. Too much gamma-ray is predicted to come from the M49 and Fornax clusters in the range $1-10$ GeV, which are in disagreement with Fermi-LAT observations. This conclusion remains unchanged even if we include the uncertainties of the total DM mass and the starlight in the cluster. In this energy range most contributions come from the ICS process, which are model-independent to a large extent. With the conventional astrophysical background of electrons and positrons, it seems unlikely for any decaying DM models to account for the PAMELA and Fermi-LAT excesses under the constraints of GeV gamma-ray flux upper limits from nearby clusters. However, it is possible that the Fermi-LAT electrons excess is due to, for instance, unidentified astrophysical sources instead of DM decay. In this case, the GUT decaying DM could interpret the PAMELA positron fraction excess and induce gamma-ray fluxes consistent with observations from nearby clusters. Unlike the GUT decaying DM model, the gamma-ray fluxes from nearby clusters are predicted to be consistent with Fermi-LAT observations for decaying DM scenario with R-parity violation. This is because the DM mass here is just around $200$ GeV. With such a relatively light DM, the ICS process hardly emits gamma-rays above $0.1$ GeV due to kinematics. The contribution to GeV gamma-ray flux comes dominantly from the FSR process. It should be kept in mind that in order to account for the Fermi-LAT $e^{\\pm}$ excess in this model, astrophysical $e^{\\pm}$ background harder than conventional one has to be adopted. As a result, such interpretation would require the PAMELA positron excess to disappear at energies above $100$ GeV, which could be checked in the near future." }, "1101/1101.5563_arXiv.txt": { "abstract": "{}{We present new spectropolarimetric observations of the chromospheric \\ion{He}{i} 10830~{\\AA} multiplet observed in a filament during its phase of activity.}{The data were recorded with the new Tenerife Infrared Polarimeter (TIP-II) at the German Vacuum Tower Telescope (VTT) on 2005 May 18. We inverted the He Stokes profiles using multiple atmospheric components.}{The observed He Stokes profiles display a remarkably wide variety of shapes. Most of the profiles show very broad Stokes $I$ absorptions and complex and spatially variable Stokes $V$ signatures. The inversion of the profiles shows evidence of different atmospheric blue- and redshifted components of the \\ion{He}{i} lines within the resolution element ($\\sim 1$~arcsec), with supersonic velocities of up to $\\sim 100$~km/s. Up to five different atmospheric components are found in the same profile. We show that even these complex profiles can be reliably inverted.}{} ", "introduction": "Filaments are typically elongated dark structures on the solar disk, characterized by additional absorption in strong chromospheric lines, like the \\ion{H}{$\\alpha$} line. They are the on-disc counterparts of prominences seen above the solar limb and represent dense and cool chromospheric gas suspended in the hot and thin corona. The strength and structure of the magnetic field plays a key role in keeping the prominence material from flowing down to normal chromospheric levels. It is therefore of considerable interest to determine the magnetic structure associated with a prominence. A number of observations of the magnetic vector in prominences have been carried out \\citep[e.g.,][]{leroy,casini,lin,Trujillo_nature,merenda,kuckein}. All these investigations, except for the last one, have been restricted to quiescent prominences. \\\\ In this paper we introduce spectropolarimetric observations in the \\ion{He}{i} 10830~{\\AA} triplet of a filament located in an active region that has been activated by a flare. Filaments also appear as dark absorption features in the \\ion{He}{i} 10830~{\\AA} line. \\\\ Spectropolarimetry in the \\ion{He}{i} triplet at 10830~{\\AA} is an important tool to determine the magnetic field vector in the solar chromosphere \\citep{Trujillo_nature,solanki_nature}. It is well suited for the measurement of the magnetic vector near the base of the solar corona because these lines are often narrow, nearly optically thin, have a reasonable effective Land\\'e factor (see Table~\\ref{tab:righe}) and are easily observed \\citep{ruedi}. Measurements of the polarization of the \\ion{He}{i} 10830~{\\AA} radiation are also indicated as a useful new tool for the diagnostics of the magnetic field in filaments \\citep{lin,Trujillo_nature}. Recently, \\citet{kuckein} studied the vector magnetic field of an active region filament by analyzing spectropolarimetric data in the \\ion{He}{i} 10830~{\\AA} lines with three different methods, one of these being, as in our case, Milne--Eddington inversions. They find the highest field strengths measured in filaments so far, around 600-700~G, and conclude that strong transverse magnetic fields are present in active region filaments. \\\\ The \\ion{He}{i} 10830~{\\AA} multiplet originates between the atomic levels $2$ $^3S_1$ and $2$ $^3P_{2,1,0}$. It comprises a component at 10829.0911~{\\AA} with $J_u=0$ (hereafter referred to as Tr1), and two components at 10830.2501~{\\AA} with $J_u=1$ (Tr2) and at 10830.3397~{\\AA} with $J_u=2$ (Tr3) which are blended at solar chromospheric temperatures. In the subsequent discussion we call this blended absorption the Tr2,3 component. \\\\ We concentrate on describing the observations and the inversion of these profiles. The He Stokes profiles in the activated filament reveal a remarkable level of complexity, requiring up to five independent magnetic components to reproduce. In a following paper the obtained maps of the magnetic and velocity fields in the filament will be presented, critically analyzed and interpreted in terms of prominence models. ", "conclusions": "We focus on the analysis and the inversion of Stokes profiles of the \\ion{He}{i} 10830~{\\AA} triplet observed in an active region filament during its phase of activity. We showed that even the most complex observed profiles can be well reproduced by multi-component atmospheres. Although we cannot completely exclude that the profiles could be simply strongly turbulently broadened, the fits based on multiple components are of clearly better quality (see below). We tested the response of the numerical code H{\\footnotesize E}LI{\\footnotesize X} for the inversion of Stokes profiles to work with a high number of atmospheric components and free parameters. This was necessary to fit not only the different observed He components, but also the photospheric and telluric lines in the wavelength range. We were able to retrieve the atmospheric parameters for the individual He components at different levels of accuracy. The values of the magnetic field strength, $B$, retrieved from the fit presented here are in the range $100-250$~G. These values are lower than those found by \\citet{kuckein} for another active prominence. This difference could be caused by the different methods employed, because unlike \\citet{kuckein} we inverted our profiles including a scattering polarization correction (see Sec.~\\ref{sec:inversion}) and the magnetic field strength we retrieved is a result of a coupling between many He components. The difference could also be merely a reflection of the different strengths of the two filaments' magnetic fields, or the analyzed profiles possibly do not sample the highest field strength region in the observed filament. The $v_{LOS}$ is well retrieved for all profiles with a small error. The inclination angle $\\gamma$ is retrieved with a somewhat larger error, but the inclinations of the different components of the magnetic field can often be distinguished. The error bars on the azimuth angle, $\\chi$, are instead quite large, and it is often impossible to distinguish between the azimuths of the different magnetic components. This is mainly because of the complexity and the noise in the $Q$ and $U$ Stokes profiles. The information in the observations for retrieving the azimuth is limited. We therefore refrain from conclusions about the azimuthal direction of the magnetic field in the filament. There is also insufficient information in the profiles to distinguish between the field strength $B$ of the different components. Note that because Stokes $V$ is generally much stronger than $Q$ and $U$, $B\\cos\\gamma$ is probably the most reliably determined quantity, although we do not explicitly tabulate it. \\\\ The analysis of the observed polarization of the \\ion{He}{i} 10830~{\\AA} multiplet in the filament, carried out by inverting the Stokes profiles, reveals different unresolved atmospheric components of the He lines, coexisting within each spatial resolution element ($\\sim1$~arcsec). The components are distinguished by their Doppler shifts, which generally differ by more than the sound speed. For more complex analyzed profiles we also tried to find a solution considering a broadened profile with some emission contributions distributed at the appropriate places, but the fits to the observed profiles were not as good as the ones we present in the paper. The main problem is reproducing the $I$ and $V$ profiles at the same time. \\\\ In this active filament we find profiles requiring up to five atmospheric components for a reasonable fit. Multiple unresolved magnetic components are found not only in filaments, but also above pores and elsewhere in active regions \\citep{aznar,lagg1}, but so far in other data sets never more than three components were needed. As \\citet{lagg1} pointed out, the geometrical interpretation of these multiple downflow components is not clear-cut in the case of a line such as \\ion{He}{i} 10830~{\\AA}, which is generally optically thin. Thus the different components could be located at different horizontal positions within the resolution element, implying considerable fine structure in the filament. This fine structure in filaments at a sub-arc s scale has been reported earlier. See, e. g., \\citet{tandberg} or \\citet{heinzel}, for reviews. These fine structures show a random motion with velocities of 5-10~km/s, as deduced mainly from observations in the \\ion{H}{$\\alpha$} line. \\\\ An alternative explanation is that the various components are located at different heights along the LOS. It may be argued that the \\ion{He}{i} line gets very strong in the filament and can hardly be considered to be optically thin, so that one should not be able to see through different layers so easily. We note, however, that the line becomes optically thick only at wavelengths at which the absorption is large, whereas it remains optically thin at neighboring wavelengths. Consequently, if the different layers absorb at wavelengths that are shifted by more than a Doppler width relative to each other, this explanation remains possible even if individual layers produce optically thick absorption. \\\\ Because in many cases the profiles have some optical thickness (they are formed over a range of heights), it is possible that gradients in velocity along the LOS can affect their shape. The influence of these gradients cannot be treated in the simple model employed here. A more general treatment with gradients might lead to good fits with a reduced number of components. \\\\ The inversions support the idea that the He components correspond to plasma trapped in different magnetic field lines, which may well be pointing in different directions along the LOS. Indeed, the coupling of the magnetic field vector between the individual components results in worse fits to the observed profiles, which supports this interpretation. Note though that some He profiles do not show a significant signal in Stokes $V$. They describe intensity features without a polarization counterpart.\\\\ Multiple magnetic components are fairly common in most parts of the observed filament and are often associated with strong blue- or redshifts corresponding to supersonic velocities. We measured downflow velocities of up to 100~km/s and upflows of up to 60~km/s along the line-of-sight. These supersonic up- and downflows always coexist with a He atmospheric component almost at rest ($-10100{\\rm MeV}} \\sim 10^{-5}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$. We also discuss the origin of the break observed in the flaring spectra of 4C\\,+21.35. We show that, in principle, the model involving annihilation of the GeV photons on the He\\,{\\tt II} Lyman recombination continuum and line emission of BLR clouds \\citep[as proposed by][]{pou10} may account for such breaks. However, we also discuss the additional constraint provided by the surprising detection of 4C\\,+21.35 at $\\lesssim 1$\\,TeV energies by the MAGIC telescope, which coincided with the second GeV flare of the source. We argue that this sub-TeV emission (as well as the GeV emission of the source, if co-spatial), is not likely to be produced inside the region of highest ionization of BLR ($\\sim 10^{17}$\\,cm), because of the expected large opacity for $\\gamma$-rays related to photon-photon annihilation on the infrared photon field provided by the accretion disk, BLR, and nuclear dust. Instead, it seems to originate further away from the active center, most likely as far as the characteristic scale of the hot dusty torus surrounding the 4C\\,+21.35 nucleus ($\\sim 10^{19}$\\,cm). After completing this work, we learned that our basic estimates regarding the TeV opacity are supported by the most recent analysis of the {\\it Spitzer Space Telescope} data by \\citeauthor{malmrose} (2011), who found a prominent infrared excess in the observed spectrum of 4C\\,+21.35, and showed that the excess is well modeled by the $\\sim 10^3$\\,K dust emission. The extremely high isotropic luminosity of the hot dust estimated by \\citeauthor{malmrose} indicates that the fraction of the disk radiation reprocessed in the inner regions of the dusty torus in 4C\\,+21.35 is as large as $\\xi_{\\rm HDR} \\sim 1$. The exact spectral shape and the luminosity of the infrared continuum of the source available in a near future, as well as of the detail of the TeV spectrum soon reported by the MAGIC Collaboration (Aleksic et al., 2011, to be submitted), will allow us to re-examine the conclusions and estimates presented in our paper." }, "1101/1101.0200_arXiv.txt": { "abstract": "Water masers have been observed in several high redshift active galactic nuclei, including the gravitationally lensed quasar MG 0414+0534. This quasar is lensed into four images, and the water maser is detected in two of them. The broadening of the maser emission line and its velocity offset are consistent with a group of masers associated with a quasar jet. If the maser group is microlensed we can probe its structure and size by observing its microlensing behaviour over time. We present results of a high resolution numerical analysis of microlensing of the maser in MG 0414+0534, using several physically motivated maser models covering a range of sizes and emission profiles. Time-varying spectra of the microlensed maser are generated, displayed, and analysed, and the behaviour of the different models compared. The observed maser line in MG 0414+0534 is consistent with maser spots as in other quasar jets, provided substructure is de-magnified or currently lost in noise; otherwise smooth extended maser models are also candidates to generate the observed spectrum. Using measures of spectral variability we find that if the maser has small substructure of $\\sim 0.002$ pc then a variation of \\shl 0.12 mag in flux and 2.0 km s$^{-1}$ in velocity centroid \\ehl of the maser line could be observed within 2 decades. For the smallest maser model in this study a magnification of $>$ 35 is possible 22\\% of the time, which is of significance in the search for other lensed masers. ", "introduction": "Gravitational lensing occurs when light travelling to Earth from distant sources is deflected by the gravitational influence of an intervening massive object \\citep{schneider,wambsganss_lr}. The light source may be within the Galaxy and lensed by planetary systems, which is useful in the search for planets \\citep[e.g.][]{gaudi}; or it may be from a very distant galaxy and lensed by another galaxy \\citep[e.g.][]{willis} or galaxy cluster \\citep[e.g.][]{sand}. In this paper we are concerned with a quasar being lensed by a galaxy, many instances of which have been discovered \\citep[e.g.][]{walsh,huchra,turner,myers,inada,kayo}. In such systems, multiple magnified images of the quasar are produced, with the image properties determined by the source emission profile, the lensing galaxy's mass distribution, and the relative position of the source and lens. Assuming a smooth mass distribution for the lensing galaxy is sometimes enough to model the observed image configurations, but in other cases the effect of the compact objects within the galaxy (stars, planets, etc.) must also be considered. Due to the intricate nature of the light paths through a galaxy of myriad objects, a slight change in the location of the background source can produce a change in the image magnifications (independent of any intrinsic source variability). This effect of the compact structure is called \\emph{microlensing} \\citep{chang,young,paczynski,wambsganss_thesis,wambsganss_saas}, and identifiable microlensing fluctuations of high amplitude are called \\emph{high magnification events}. These events can last for relatively short durations -- months or even weeks \\citep{shalyapin} -- and usually occur with a frequency of decades. The angular size of the source is important for microlensing, as larger sources will ``wash out'' the location-induced variability in magnification \\citep{lewis1,bate}, so it is small sources that produce the most significant high magnification events. \\shlA The shape of the source emission profile also affects microlensing \\citep{mortonson}. \\ehlA Therefore, observations and modelling of flux variability can provide an estimate of source size and flux emission profile. Apparent chromatic effects are also seen \\hl{\\citep{lewis2, eigenbrod}} because if an extended source emits different frequencies from separate regions, the frequencies will be magnified differently, and the source spectrum will be altered. Like the magnification, the spectrum may also change over time, and studying spectral variability can lead to possible spectral emission profiles. In these (and other) ways microlensing becomes a useful tool for astronomical research \\citep{Kayser,wambsganss_anp,kochanek,wambsganssr}. \\begin{figure*} \\centering % \\subfigure[Grainy Patch Schematic]{\\includegraphics[width=55.3mm]{Figures/spots_sch.eps}} \\subfigure[Grainy Ring Schematic]{\\includegraphics[width=55.3mm]{Figures/ring_sch.eps}} \\subfigure[Hollow \\& Solid Jets Schematic (1 jet shown)]{\\includegraphics[width=57.41mm]{Figures/jet_sch.eps}} \\subfigure[Grainy Patch: a patch of maser spots]{\\includegraphics[width=55.3mm]{Figures/spots.eps}} \\subfigure[Grainy Ring: a ring of maser spots that can be oriented]{\\includegraphics[width=55.3mm]{Figures/ring.eps}} \\subfigure[Hollow Jets: a bi-conical jet structure of smooth masers that can be oriented. The Solid Jet is the same but filled with maser material.]{\\includegraphics[width=57.41mm]{Figures/jet_side.eps}} \\caption{Schematics (a-c) and example images (d-f) of the three complex maser models (the Single Spot and Smooth Patch are not shown). Figures (a) and (d) show the \\shl Grainy Patch;\\ehl\\ (b) and (e) the \\shl Grainy Ring; \\ehl (c) and (f) an example of the Hollow Jets, which is bi-conical as in (f), but since the jets are mirror images of each other only one schematic is shown in (c). The inclination of the jets in (f) is $i = 30\\,^{\\circ}$ South-East, relative to the magnification map as described in the text. The elements of the schematics are not to scale.} \\label{schemes} \\end{figure*} In this paper we present results of numerical analyses of the microlensing of a water maser in a known microlensing system, MG 0414+0534. We use several physically-motivated models to represent the maser source; each model is characterised by its shape, size, emission and velocity profile. A model is then subjected to numerical microlensing so that the alteration in the spectrum and flux after lensing can be found. The structure of the paper is as follows: Section \\ref{Background} discusses the maser in MG 0414+0534, and mechanisms used for numerical microlensing analysis. In Section \\ref{Method} we describe the parameters characterizing the MG 0414+0534 lens, we specify the maser source models, and explain how the models are used. Section \\ref{Results} presents the results of microlensing of the maser by examining which models induce the most variability, what time scales may be expected for this, what forms the spectra take and how they vary over time, and whether the models match the current observation. ", "conclusions": "Observations of jet masers in active galaxies indicate that the masers typically exist in a group with a velocity range of the order 100 km s$^{-1}$. The maser in MG 0414+0534 consists of a single peak. This is not consistent with a group of maser spots that produces multiple peaks, unless some peaks are lost in noise, or currently de-magnified, which is certainly possible. On the other hand the observation could be a smooth patch of masers, or a smooth distribution through the inner parts of the quasar jets, all of which also produce a single peak. A patch of spots will produce the most variability, in both velocity centroid and magnitude, the smooth patch will produce less -- it may drift in velocity but remain fairly constant in shape -- the jets will not change at all. A source with substructure like a patch of spots could produce changes in the spectrum within a couple of hundred years, lasting for 10-20 years. These results are just the beginning of the numerical analysis of this lensed maser, and continuing work would include: \\begin{itemize} \\item \\shl Developing physically consistent models based on future observations and new understanding of masers \\item \\shl Developing statistical tools for the extraction of histograms, structure functions, and other data from the high resolution magnification maps \\end{itemize} but it would be \\shl prudent to wait for observations that have a higher sensitivity and resolution of flux and velocity, and in particular that are able to detect the maser in images B and C. A long-term monitoring program \\ehl should be carried out on this source, either with a large single dish or a sensitive VLBI array (e.g. the European VLBI Network) with higher resolution than the original Effelsberg observations, with the following aims: \\begin{itemize} \\item monitoring should be conducted on monthly time scales to pick up short term fluctuations as well as building statistics for long term variability \\item the A1 and A2 image spectra should be separated, which allow a comparison of the spectral behaviour between the two images for all the models, adding information that may constrain the choice of model \\item \\sbl observations at higher spectral resolution may allow a more accurate determination of the spectral shape; differential fluctuations over time due to substructure will help to determine the morphology of a maser group. \\ehl \\end{itemize}" }, "1101/1101.4942_arXiv.txt": { "abstract": "We present new imaging at 12.81 and 11.7 $\\mu$m of the central $\\sim$40$''$$\\times$30$''$ ($\\sim$0.7$\\times$0.5 kpc) of the starburst galaxy M82. The observations were carried out with the COMICS mid-infrared (mid-IR) imager on the 8.2~m Subaru telescope, and are diffraction-limited at an angular resolution of $<$0$\\farcs$4. The images show extensive diffuse structures, including a 7\\arcsec--long linear chimney-like feature and another resembling the edges of a ruptured bubble. This is the clearest view to date of the base of the kpc-scale dusty wind known in this galaxy. These structures do not extrapolate to a single central point, implying multiple ejection sites for the dust. In general, the distribution of dust probed in the mid-IR anticorrelates with the locations of massive star clusters that appear in the near-infrared. The 10--21~\\micron\\ mid-IR emission, spatially-integrated over the field of view, may be represented by hot dust with temperature of $\\sim$160 K. Most discrete sources are found to have extended morphologies. Several radio \\hii\\ regions are identified for the first time in the mid-IR. The only potential radio supernova remnant to have a mid-IR counterpart is a source which has previously also been suggested to be a weak active galactic nucleus. This source has an X-ray counterpart in \\c\\ data which appears prominently above 3 keV and is best described as a hot ($\\sim$2.6 keV) absorbed thermal plasma with a 6.7 keV Fe K emission line, in addition to a weaker and cooler thermal component. The mid-IR detection is consistent with the presence of strong \\neii\\l 12.81\\micron\\ line emission. The broad-band source properties are complex, but the X-ray spectra do not support the active galactic nucleus hypothesis. We discuss possible interpretations regarding the nature of this source. ", "introduction": "The galaxy M82 (NGC~3034) hosts the nearest and best example of an ongoing massive starburst, making it an excellent target for detailed studies at all wavelengths. The galaxy is thought to have undergone an interaction event with its neighbor M81 about 10$^8$ yr ago \\citep{gottesman77}, triggering a massive nuclear starburst about 5$\\times$10$^7$~years ago \\citep{rieke80}. There is also evidence for several other star formation episodes, both older and younger \\citep{degrijs01, forsterschreiber03}. Around 40 supernova remnants (SNRs) have been identified in the core \\citep{fenech08}, and one new supernova (SN) is produced every $\\sim$3 years \\citep[e.g. ][]{rieke80, jones84}. The energy output of the super star clusters hosting the SNe is thermalised and drives a large scale superwind along the galactic minor axis \\citep{heckman90} which can be observed in detail because of the favorable edge-on inclination of the source. In the infrared, M82 has been extensively studied with all space missions. Its infrared luminosity is measured to be 5$\\times$10$^{10}$ $L_{\\odot}$, and it shows prodigious dusty outflows and polycyclic aromatic hydrocarbon (PAH) grains in the mid-infrared (mid-IR) extending on kpc-scales \\citep[e.g. ][]{iraspsc88, sturm00, forsterschreiber03b, engelbracht06, kaneda10, roussel10}. There are very few high-spatial-resolution mid-IR studies of the core itself on scales of order 100 pc, because of size limitations of space missions. With their large primary mirrors, ground-based telescopes can achieve the best spatial resolution currently possible. In the $N$ (8--13 \\micron) band atmospheric window, the highest resolution studies thus far are the works of \\citet[][hereafter TG92]{telesco92} and \\citet[][ hereafter AL95]{achtermannlacy95}, with nominal resolutions of 1$\\farcs$1 and 2\\arcsec, respectively. In this work, we present the first sub-arcsec mid-IR $N$-band images of the core of M82. The galaxy was observed as an extension of our recent work of Seyfert galaxies \\citep{g09_mirxray}, as part of a study to understand the mid-IR emission of galaxies at high angular resolution. The observations were carried out at the 8.2~m Subaru telescope. At wavelengths of 11.7 and 12.81 \\micron, our imaging is diffraction-limited at $<0\\farcs 4$. These new images provide the sharpest mid-IR view of structures at the base of the superwind, and allow an extensive multi-wavelength comparison of individual sources. Several \\hii\\ regions are newly identified in the mid-IR. We also discuss the nature of a putative active galactic nucleus candidate. Using a mid-IR detection and new \\c\\ data, we rule out the AGN hypothesis and discuss other possibilities, include supernova remnant ionization and emission from a starburst. Distance estimates to M82 have ranged over 3.2--5.2 Mpc \\citep[e.g. ][]{burbidge64,tully88,sakaimadore99}. Some of the latest measurements suggest a distance at the lower end of this range based upon accurate determination of the tip of the red giant branch magnitude \\citep{karachentsev04,dalcanton09}. We adopt a value of 3.53 Mpc herein, resulting in a physical scale of 17.1 pc per arcsec. Our imaging resolution limit corresponds to $\\approx$6.1 and 6.7 pc at 11.7~\\micron\\ and 12.8~\\micron, respectively. ", "conclusions": "We have presented the highest resolution imaging of the nuclear regions of M82 to date in the mid-IR. The Subaru diffraction limit at wavelengths of 11.7 and 12.81~\\micron\\ is 0\\farcs 36 and 0\\farcs 39, respectively. Compared to the previous best resolution observations (the 12.4~\\micron\\ imaging with PSF of 1$\\farcs$1 presented by TG92), our images have PSFs improved by factors of 3.1 and 2.8 at 11.7 and 12.81~\\micron, respectively, and we cover a core region larger by a factor of at least two. The referee also made us aware of a conference proceeding by \\citet{ashby94}, where an image at 11.7~\\micron\\ with an angular resolution of 0\\farcs 6 is presented. The absolute astrometry of their image agrees to within $\\sim$1\\arcsec\\ with our work, as well as with the works of AL95 and TG92. To our knowledge, no further details have been published. Multi-wavelength image comparisons show an anti-correlation between the observed stellar distribution (probed in the near-IR) with the distribution of warm dust that we probe. This means that obscuration in these dusty regions is heavy enough to scatter even near-IR radiation. Our observations provide the best view of the base of the dusty superwind that is known to exist in this galaxy, and reveal several elongated features on projected physical scales of up to $\\ell$=120 pc at least. If the mid-IR emitting dust is mixed in with and entrained in outflowing gas, then the travel time ($t$) is \\begin{equation} t=5.9\\left( \\frac{\\ell}{120\\ {\\rm pc}}\\right) \\left(\\frac{200\\ {\\rm km\\ s^{-1}}}{\\rm v}\\right) \\times 10^5\\ {\\rm yr}. \\end{equation} \\noindent where a velocity of ${\\rm v}$=200 km s$^{-1}$ identical to that found by \\citet{nakai87} for the molecular gas is used as reference. This time period suggests recent energy input from young starbursts. The kinetic energy required to expel this dust is only a fraction of that channeled into the total gas mass, and so expulsion by SNe occurring at various locations around the ring of star-formation can easily account for this \\citep{nakai87}. Using the integrated fluxes over the COMICS field-of-view alongside archival data at longer and shorter wavelengths, and assuming only a two temperature phase for the dust, we are able to describe the broad-band mid-IR diffuse emission with a $T$=160 K modified black body with a hot dust mass of $\\sim$1000~\\Msun, in addition to the 45 K cool dust component responsible for the far-IR emission. Assuming a standard gas:dust ratio of 100, the mass of the gas associated with the mid-IR emitting dust is then $\\sim$10$^5$~\\Msun. This is lower than the ionized gas mass of 2$\\times$10$^6$~\\Msun\\ (TH80; \\citealt{willner77}) by a factor of about 20, meaning that the ionized gas may be the predominant environment for the hot dust that we are observing. The cooler dust, on the other hand, is likely to be tracing the distribution of molecular gas (TH80). More than 20 discrete sources are detected, and most are found to have extended profiles. Matching against radio catalogs suggests that we have resolved at least four (and tentatively, five) \\hii\\ regions in the mid-IR for the first time. These \\hii\\ regions have monochromatic continuum dust luminosities ranging from $\\nu L_{\\nu}^{\\rm 11.7\\ \\mu m}$=2$\\times$10$^6$\\lsun\\ (radio ID 39.29+54.2) to 8$\\times$10$^6$\\lsun\\ (41.17+56.2), consistent with being powered by embedded super star clusters similar to those seen in some other nearby galaxies \\citep[e.g. ][]{galliano08}. As seen in Fig.~\\ref{fig:rgb}, no source is detected at the position of the radio transient identified by \\citet{brunthaler09} as SN2008iz, the brightest radio SNR in M82 over the past 20 years. The start of bright radio flaring activity is limited to the time interval of 2007 Oct 29 to 2008 Mar 24. Our observations were carried out one month after the final set of follow-up radio observations on 2009 Apr 04 reported by \\citeauthor{brunthaler09}, when they found the source to have an integrated 22.2 GHz flux of 9.2\\p 0.2 mJy. Our flux limits are $F_{\\rm NeII}$$<$38 mJy and $F_{\\rm N11.7}$$<$18 mJy, assuming a point source of angular size equal to the diffraction limits in each filter. In general, there is little overlap between confirmed SNRs and mid-IR detections. Additionally, no mid-IR counterpart is detected at the location of the unusual radio transient found by \\citet{muxlow10}. The source appeared on radio images obtained within the period of 1--5 May 2009 -- contemporaneous with our mid-IR observations -- and was not present one week before. It is located $\\approx$1~\\arcsec\\ from the position of our source \\# 18 (see also \\citealt{kong09_atel} for an identification of the X-ray counterpart). We estimate point source upper limits of 34 mJy and 15 mJy in the NeII and N11.7 filters. No space observatory is foreseen to have a better resolving power than Subaru at $\\sim$10~\\micron, though the Mid InfraRed Instrument (MIRI) onboard \\jwst\\ should provide excellent sensitivity for only a modest loss in angular resolution at the same wavelengths \\citep[e.g. ][]{jwstmiri}. MIRI is also expected to have a field of view larger by a factor of $\\gtsim$2 on a side, making source identification much more secure. On a longer timescale, the European/JAXA mission \\spica\\ \\citep{spica} will be crucial for deep searches of far-IR signatures of AGN activity at this source location. On the ground, the Gran Telecopio Canarias, with a primary mirror diameter of 10.4~m, can improve the resolution slightly to 0\\farcs 3 under good observing conditions. Even better resolution must await larger ground-based observatories such as the Extremely Large Telescope. This will not be diffraction-limited under natural seeing conditions and will require additional adaptive optics capabilities in the mid-IR to improve upon the results presented herein. \\subsection{Nature of AGN candidate source} The puzzling radio source 44.01+59.6 has a significant \\neiifilter\\ mid-IR detection (our source \\# 18 [I52.70+45.9]), and is also coincident with a source visible in high-resolution X-ray images (J095552.7), to well within our estimated absolute positional uncertainties. From \\c\\ data, it is immediately apparent that the X-ray spectrum is not consistent with that of AGN, which usually display broad-band power-laws with photon-indices $\\sim$1.9 characteristic of radiatively-efficient accretion \\citep[e.g. ][]{mateos05_wide}. Low luminosity AGN with radiatively-inefficient flows or jets, too, usually display hard X-ray continua \\citep[e.g. ][]{yu10}. Similarly, one can argue against direct association with other kinds of accreting sources such as X-ray binaries. The detection of a strong He-like Fe line instead of a neutral line centered on 6.4 keV implies the absence of cold reflecting matter, also atypical for an accreting source. Highly ionized lines have actually been attributed to heavily obscured AGN in some cases \\citep[e.g. ][ and references therein]{iwasawa05_arp220, nandraiwasawa07}, though on larger scales. In such a scenario, the AGN itself is not readily apparent; rather, the visible spectrum is dominated by surrounding gas which is irradiated by the AGN along directions out of our line of sight. If this is true for the case of source J095552.7, the underlying accreting object (be it an AGN or an X-ray binary) could power the radio jet found by \\citet{wills99} and also provide a ready source where accreting gas is in rotation (as inferred from the maser lines found by \\citealt{seaquist97}). But this would also result in X-ray photoionization spectral features or even an ionized reflection continuum. An extra layer of optically-thick obscuration is then also required (in addition to the two that we detect) with an extreme covering factor very close to unity so as to completely hide the reflection continuum and any neutral Fe K line. This is not consistent with the detection of a jet which would be expected to decrease the covering factor by clearing away some surrounding matter. Finally, the 2--10 keV luminosity of 5$\\times$10$^{38}$ erg s$^{-1}$ means that the source is unlikely to have a bolometric X-ray power which would place it in the ULX regime (e.g. \\citealt{makishima00}) because of the steep spectra of the fitted thermal models. Thus, our analysis allows us to reach some firm conclusions as to what the source is {\\em not}. The true source nature still remains uncertain, though, and in the following, we discuss plausible alternatives. In addition to our \\neii\\ image, the source appears most prominently as a compact object in hard X-rays. One possibility may then be that both the mid-IR and the hard X-ray components are associated with the SNR visible in the radio. To test this, we may compare the source power to that of Cas A, a powerful and young Galactic SNR. From \\suzaku\\ observations summed over the entire remnant, the 4--10 keV X-ray luminosity of Cas A is determined to be $L_{4-10\\ {\\rm keV}}^{\\rm Cas A}$$\\approx$4$\\times$10$^{35}$ erg s$^{-1}$ \\citep{maeda09}. This is about 150 times smaller than the deabsorbed X-ray luminosity $L_{4-10\\ {\\rm keV}}$$\\approx$6$\\times$10$^{37}$ erg s$^{-1}$ that we find for our source by using the high temperature {\\sc apec} component alone (\\S~\\ref{sec:xspec}). This requires that the electron density ($n_e$) in the X-ray emitting interstellar medium (ISM) around source \\# 18 be much larger than in the case of Cas A. $n_e$ may be determined by using the normalization values returned by the {\\sc apec$_2$} component and the emission volume which, in the scenario under consideration here, should correspond to the volume of the radio-emitting plasma. \\citet[][ see their Table 3]{fenech08} measure a diameter of 0.86~pc for the SNR, using which yields an electron density of $n_e$$\\sim$1.75(\\p 0.3)$\\times$10$^3$ cm$^{-3}$. A low plasma filling factor would only push this value up. This is, indeed, much larger than typical ISM densities of 1--10 cm$^{-3}$ relevant for Galactic SNRs. We also note that the present limit on any X-ray variability in the source ($\\ltsim$20\\%\\ over seven years) derived from archival data comparisons in \\S~\\ref{sec:variability} is consistent with the observed smaller flux changes in Cas A over similar timescales \\citep[e.g. ][]{patnaude10}. The detected collimated radio jet may then suggest some kind of axisymmetry in the progenitor (e.g. a binary merger event). On the other hand, the mid-IR power for source \\# 18 ($\\lambda L_{\\lambda(12.81\\ \\mu{\\rm m})}$=1.5$\\times$10$^{40}$ erg s$^{-1}$) is in excess of the integrated 12~\\micron\\ \\iras\\ power of Cas A ($\\lambda L_{\\lambda(12\\ \\mu{\\rm m})}^{\\rm Cas A}$=5$\\times$10$^{36}$ erg s$^{-1}$; \\citealt{saken92}) by a much larger factor of $\\approx$3000. But given the unknown dust temperature, the fraction of freshly formed vs. ambient heated dust, and the fact that our \\neiifilter\\ filter flux is likely to be dominated not by continuum but by the \\neii\\l12.81\\micron\\ line (which is known to be a strong cooling line for SNRs; e.g. \\citealt{rho08}), further comparison of the mid-IR fluxes is difficult. The other possibility to consider for the X-ray thermal plasma components is that of a hot ISM phase in a compact star cluster. In fact, collisionally ionized plasmas with a wide range of temperatures have been previously inferred to exist within starburst galaxies \\citep[e.g. ][]{iwasawa05_arp220, ranalli08, strickland09}, though on larger scales. Fig.~\\ref{fig:nucleussed} shows the compiled radio--to--X-ray fluxes for source \\# 18. A mid-IR flux dominating the SED is consistent with an origin in a starburst, though this is not a unique solution. The detection of two distinct layers of absorption (\\S~\\ref{sec:xspec}) may also suggest a scenario which combines the above possibilities as follows. The hot X-ray thermal plasma and the radio counterpart could both be associated with a SNR, which is embedded within a host star cluster. The high column density affecting the {\\sc apec}$_2$ component (Table~\\ref{tab:xspec}) may easily be explained by strong local absorption within molecular clouds, for instance. The point-like profile of the hard X-ray data is also consistent with this. The cluster, on the other hand, could appear prominently in soft X-rays and also in the mid-IR. The comparatively low aborption affecting this component is then due to gas along the plane of the galaxy on larger scales. Further insight into the nature of this source will be possible if it can be isolated at longer wavelengths characterizing the peak of typical star-formation SEDs (this may be within reach of \\jwst), or through detection of spectral features in the sub-mm with ALMA. Subsequent long \\c\\ exposures would be useful to place tighter constraints on X-ray variability. Meanwhile, the question of whether M82 hosts an AGN or not remains to be answered. \\begin{figure} \\begin{center} \\includegraphics[height=8.5cm,angle=90]{fig12.ps} \\end{center} \\caption{ Broad-band SED of source \\# 18. The radio data are the 1.3 cm to 74 cm fluxes (and one limit) from \\citet{allenkronberg98}. In the mid-IR, we plot our \\neiifilter\\ filter flux, and a 3$\\sigma$ \\n117filter\\ detection limit measured assuming Poisson statistics within an aperture equal in size to the diffraction limit. The X-ray regime shows the unfolded spectrum relevant for the model fitted to the observed data in Fig.~\\ref{fig:agn}. \\label{fig:nucleussed} } \\end{figure}" }, "1101/1101.1560_arXiv.txt": { "abstract": "{Ultraviolet emission from the first generation of stars in the Universe ionized the intergalactic medium in a process which was completed by $z \\sim 6$; the wavelength of these photons has been redshifted by $(1+z)$ into the near infrared today and can be measured using instruments situated above the Earth's atmosphere. First flying in February 2009, the Cosmic Infrared Background Experiment (CIBER) comprises four instruments housed in a single reusable sounding rocket borne payload. CIBER will measure spatial anisotropies in the extragalactic IR background caused by cosmological structure from the epoch of reionization using two broadband imaging instruments, make a detailed characterization of the spectral shape of the IR background using a low resolution spectrometer, and measure the absolute brightness of the Zodical light foreground with a high resolution spectrometer in each of our six science fields. This paper presents the scientific motivation for CIBER and details of its first two flights, including a review of the published scientific results from the first flight and an outlook for future reionization science with CIBER data.} \\FullConference{Cosmic Radiation Fields: Sources in the early Universe\\\\ November 9-12, 2010\\\\ DESY, Germany} \\begin{document} ", "introduction": "\\label{S:intro} The extragalactic background light (EBL) is the sum of all of the light emitted throughout the history of the Universe. At near infrared (IR) wavelengths, the EBL is predominantly due to stellar emission from nucleosynthesis (see \\citealt{Hauser2001} for a review). However, absolute measurements of the EBL at near IR wavelengths to date lack the systematic error control to constrain the radiative content of the cosmos to within an order of magnitude (\\citealt{Hauser1998}, \\citealt{Dwek1998}, \\citealt{Gorjian2000}, \\citealt{Wright2001}, \\citealt{Cambresy2001}, \\citealt{Matsumoto2005}, \\citealt{Levenson2007}). Further, the summed contribution of galaxies to the EBL does not reproduce the EBL measured by absolute photometric instruments. For example, at $\\lambda = 3.6 \\, \\mu$m the EBL measured by DIRBE from absolute photometry is $12.4 \\pm 3.2 \\,$nW m$^{-2}$ sr$^{-1}$ \\citep{Wright2000}, while the deepest pencil beam surveys with Spitzer give $6{-}9 \\,$nW m$^{-2}$ sr$^{-1}$ (\\citealt{Fazio2004}, \\citealt{Sullivan2007}). At shorter wavelengths this divergence is even more pronounced. The discrepancy between absolute photometric measurements and integrated number counts leaves open the possibility that there exists some truly diffuse emission comprising a fraction of the near IR background and that using suitable instruments it may be possible to measure it. Importantly, as the EBL traces star formation throughout the history of the Universe, it contains information about the earliest generation of stars which were responsible for ionizing it (see \\citealt{Fan2006} for a review). Though the spectrum of the near IR EBL does contain information about the epoch of reionization, it is expected to be faint compared to the contribution from galaxies. However, as the signature from the first stars is expected to be both diffuse and structured on large scales, the spatial fluctuations of its imprint on the near IR EBL yield a great deal of information about the formation of structure during the time of the first stars \\citep{Cooray2004}. The Cosmic Infrared Background Experiment (CIBER) is a series of sounding-rocket borne instrument payloads designed to measure both the spectrum of the EBL and the spatial fluctuations in the EBL imprinted during the epoch of reionization (REBL). The first CIBER payload configurations will probe the power spectrum of near IR EBL fluctuations for an REBL component, limit the strength of the Lyman cutoff signature of reionization in the near IR EBL between the optical and near-infrared EBL measurements, and measure the EBL from $0.7{-}2.1 \\, \\mu$m down to the zodiacal foreground subtraction limit. This paper reviews the scientific motivation for CIBER, gives a brief discussion of the instrumentation package, and finally discusses the outlook for measurement of the epoch of reionization using the near IR EBL over the coming decade. ", "conclusions": "" }, "1101/1101.5343_arXiv.txt": { "abstract": "{PSR J1713+0747 is a binary system comprising millisecond radio pulsar with a spin period of 4.57 ms, and a low-mass white dwarf (WD) companion orbiting the pulsar with a period of 67.8 days. Using the general relativistic Shapiro delay, the masses of the WD and pulsar components were previously found to be $0.28\\pm 0.03 M_{\\odot}$ and $1.3\\pm 0.2 M_{\\odot}$ (68\\% confidence), respectively.} {Standard binary evolution theory suggests that PSR J1713+0747 evolved from a low-mass X-ray binary (LMXB). Here, we test this hypothesis.} {We used a binary evolution code and a WD evolution code to calculate evolutionary sequences of LMXBs that could result in binary millisecond radio pulsars such as PSR J1713+0747. } {During the mass exchange, the mass transfer is nonconservative. Because of the thermal and viscous instabilities developing in the accretion disk, the neutron star accretes only a small part of the incoming material. We find that the progenitor of PSR J1713+0747 can be modelled as an LMXB including a donor star with mass $1.3-1.6 M_{\\odot}$ and an initial orbital period ranging from 2.40 to 4.15 days. If the cooling timescale of the WD is 8 Gyr, its present effective temperature is between 3870 and 4120 K, slightly higher than the observed value. We estimate a surface gravity of ${\\rm Log} (g) \\approx 7.38 - 7.40$.} {} ", "introduction": " ", "conclusions": "" }, "1101/1101.1491.txt": { "abstract": "We present a study of ionized gas, PAHs, and H$_{\\rm 2}$ emission in the halos of three edge-on galaxies, NGC 891, NGC 5775 and NGC 3044, based on $10-20$ $\\mu$m {\\it Spitzer Space Telescope} spectra. The [Ne$\\,$III]/[Ne$\\,$II] ratio, an excellent measure of radiation hardness, rises with $z$ in the halo of NGC 891. It is also higher in the halo of NGC 5775 than in the disk. NGC 3044 presents a more confusing situation. To explain the [Ne$\\,$III]/[Ne$\\,$II] as well as optical line ratio behavior in NGC 891, we carry out a simple exploration of parameter space with CLOUDY, which indicates a large increase in radiation temperature with height. Illustrative examples of physical models using a Monte Carlo radiative transfer code show that the rising neon ratio may be explained by adding a vertically extended, hot stellar source to a thin disk of massive stars. However, several other sources of hard spectra may be relevant. PAH features have scale heights of 430--530 pc in NGC 891 and 720--1080 pc in NGC 5775, suggesting they can be transported by disk-halo flows. Within NGC 891 and NGC 5775, scale heights are similar for all PAHs. For NGC 891, the scale heights exceed that of 8 $\\mu$m emission, indicating a transition from more ionized to more neutral PAHs with height. Most PAH equivalent widths are higher in the halos. H$_2$ 17.03 $\\mu$m emission with scale heights of $550-580$ pc in NGC 891 and 850 pc in NGC 5775 suggests a molecular component in a surprisingly thick layer. ", "introduction": "Gaseous halos of spiral galaxies have grown in importance in recent years as they are at the interface between galaxies and their environments, and as such are the sites of many physical processes that dictate galaxy evolution. Supernovae and stellar winds in disks can pump mass and energy into halos setting up a disk-halo cycle of gas (\\citealt{1989ApJ...345..372N}), depending on the level of star formation activity. Halos may also contain gas accreting onto galaxies, either primordial or from companions (\\citealt{2008A&ARv..15..189S}). Gas from these two sources may interact dynamically, thermally, and chemically. Such feedback from star formation is believed to be important for understanding the growth of galaxies and their current properties \\citep{1991ApJ...379...52W}. The discovery of thick layers of interstellar gas and dust, whether referred to as ``halos'' or ``extraplanar'' components, has opened up a window on these issues. In the past twenty years or so, such thick layers have been discovered in just about every component of the ISM, especially in X-rays (e.g., \\citealt{2006A&A...448...43T,2004ApJS..151..193S}), HI (e.g., \\citealt*{2007AJ....134.1019O}), radio continuum (e.g., \\citealt*{2006A&A...457..121D}), diffuse ionized gas (DIG; e.g., \\citealt{2003A&A...406..505R,1996ApJ...462..712R}), and dust (\\citealt{1999AJ....117.2077H,1998ApJ...507L.125A,2000A&AS..145...83A}). Particularly for the DIG, X-ray, radio continuum and dust components, their brightness and extent correlate with the level of disk star formation, both within and among galaxies, indicating an origin in a disk-halo flow (\\citealt{1996ApJ...462..712R,1998PASA...15..106R, 2003A&A...406..493R,2006A&A...457..779T,2004ApJS..151..193S,2006A&A...457..121D, 1999AJ....117.2077H}). Yet, some halo gas may be due to continued primordial infall onto galaxy disks, which some lines of evidence indicate is necessary to maintain star formation (e.g., \\citealt*{1997ApJ...477..765C}; \\citealt*{2010ApJ...717..323B}). Also, the metallicities of many Galactic High Velocity Clouds are as low as 0.1 Z$\\sun$ (\\citealt{2004Ap&SS.289..381W}), possibly indicating a mixing of primordial and processed gas. Finally, lagging DIG (e.g., \\citealt{2007ApJ...663..933H}) and HI (\\citealt{2007AJ....134.1019O}) halos have also been interpreted as a signature of such mixing (\\citealt{2008MNRAS.386..935F}), although other explanations are possible \\citep{2002ASPC..276..201B}. \\subsection{Spectroscopy of DIG} The DIG is a valuable tracer of energetic processes occurring in gaseous halos, largely through the substantial leverage provided by observations of emission lines. Studies of halo or \\lq\\lq extraplanar\\rq\\rq DIG in external edge-ons have been motivated by the detailed characterization of the DIG (or Reynolds Layer) of the Milky Way (\\citealt{1990ApJ...349L..17R,2009RvMP...81..969H}), and the questions that have arisen therefrom. One of the most important questions is what keeps the DIG ionized. Despite the large vertical extent of extraplanar DIG layers in many galaxies (e.g., \\citealt{1996ApJ...462..712R}), there is much evidence from various diagnostic line ratios that the primary source of ionization is radiation leaking from a thin disk of massive stars (e.g., \\citealt* {1997ApJ...474..129R,1998ApJ...501..137R,1997ApJ...483..666G,2001ApJ...551...57C, 2002ApJ...572..823O,1999ApJ...523..223H,2003ApJ...586..902H,2006ApJ...652..401M}). In the Milky Way, such radiation is also the only source that can meet the energetic requirement of keeping the DIG layer ionized (\\citealt{1990ApJ...349L..17R}). Most commonly observed is that the ratios of [S$\\,$II]$\\lambda\\lambda 6716,6731$ and [N$\\,$II]$\\lambda\\lambda 6548,6583$ to H$\\alpha$ generally increase with distance from the suspected ionizing source, in the Milky Way and in external galaxy disks and halos (all emission lines, and ionization energies necessary to create the responsible ions, relevant for this paper are listed in Table 1). This is a signature of a falling ionization parameter, $U$, (e.g., \\citealt{1994ApJ...428..647D,1998ApJ...501..137R}) but other factors have been argued to affect line ratios significantly -- namely elevated gas temperatures in the diffuse gas (\\citealt{2001ApJ...548L.221R}; \\citealt{1999ApJ...523..223H}) and hardening of the radiation field during its propagation (\\citealt{2003ApJ...586..902H,2004MNRAS.353.1126W}, hereafter WM). The situation is further complicated by the spatial behavior of other line ratios, such as [O$\\,$III]$\\lambda 5007$/H$\\beta$, [O$\\,$I]$\\lambda 6300$/H$\\alpha$, and He$\\,$I$\\,\\lambda 5876$/H$\\alpha$ (see \\citealt{2009RvMP...81..969H} for a summary). The first of these ratios, in particular, shows behavior in many galaxies suggesting that some DIG ionization occurs through processes other than photo-ionization by a thin disk of massive stars. In several external spiral galaxies, both edge-on and more face-on, this ratio increases or stays relatively constant with distance from the ionizing stars (e.g., \\citealt*{1998ApJ...501..137R, 2000ApJ...537L..13R,1999AJ....118.2775G}; \\citealt{2002ApJ...572..823O, 2003ApJ...586..902H,2003ApJ...592...79M}). An increase with height has also been found in the Reynolds Layer in the inner Galaxy (\\citealt{2005ApJ...630..925M}). This may be partially explained by elevated gas temperatures, although \\citet{2001ApJ...551...57C} find this explanation insufficient for NGC 5775 and UGC 10288. In cases such as these and the well-studied NGC 891, other sources of ionization are indicated, with shocks receiving the most attention (e.g., \\citealt{2001ApJ...551...57C}). Not all extraplanar DIG layers may require such a source, however. Some show a drop of [O$\\,$III]/H$\\beta$ with height, $z$ (\\citealt{2000A&A...362..119T,2003ApJ...592...79M}). One of these is NGC 3044, especially clear in the more westerly of the two long slit spectra presented by \\citet{2000A&A...362..119T}. A further complication is the assumed abundances. In general, solar or ``ISM'' abundances are assumed in photo-ionization models (e.g., \\citealt{1994ApJ...428..647D,2000ApJ...528..310S}; WM), but one can expect some sensitivity of line ratios to abundances, either directly, or through variation in the cooling efficiency. \\citet{2000ApJ...528..310S} find that most line ratios vary inversely with abundance in their photo-ionization models, although \\citet{1994ApJ...428..647D} and WM find a more complicated relationship. Given the possibility that halo gas contains a mix of disk-halo cycled and primordial gas, abundance variations may well be important. Finally, one should mention the complication caused by extinction, which has been found to be significant even at $z=1-2$ kpc from the disk in several edge-on spirals, including NGC 891 (\\citealt[][and references therein]{1999AJ....117.2077H,2000AJ....119..644H}). We crudely estimate in this paper that there should be many magnitudes of visual extinction in the midplanes of the galaxies we report on. Therefore, apart from the effect of extinction on line ratios, the depth along the line of sight to which we are probing optically likely varies significantly with $z$. With these various difficulties in interpreting optical line ratios, it is desirable to have a diagnostic of ionization with a more straightforward interpretation. Such an opportunity is provided by the Infrared Spectrograph (IRS; \\citealt{2004ApJS..154...18H}) on board the {\\it Spitzer Space Telescope} (\\citealt{2004ApJS..154....1W}), which allows measurement of the ratio of the 15.56 $\\mu$m [Ne$\\,$III] (ionization potential 41.0 eV) and 12.81 $\\mu$m [Ne$\\,$II] (ionization potential 21.6 eV) lines, as well as others. This ratio provides a diagnostic of the hardness of ionizing radiation that is relatively insensitive to extinction, gas-phase abundances, and gas temperature (being low excitation lines in warm gas) -- three of the biggest sources of confusion for the optical lines [the ratio is enhanced in low-metallicity galaxies, but this is attributed to the associated harder stellar radiation fields (e.g. \\citealt{2010ApJ...712..164H}, \\citealt{2009ApJ...704.1159H})]. Pioneering work with the {\\it Infrared Space Observatory} (ISO) first demonstrated the power of this diagnostic (e.g., \\citealt{2000ApJ...539..641T}; \\citealt{2002ApJ...566..880G}; and \\citealt{2003A&A...403..829V}). In \\citet*[hereafter RWB]{2008ApJ...680..263R} we presented our first measurements of this ratio in the DIG halo of NGC 891 (assumed distance 9.5 Mpc), at $z=1$ kpc on opposite sides of the disk. The ratio was found to be elevated in both extraplanar pointings relative to the disk. If the DIG is ionized by radiation leaking out of the disk, then the same stars, with the same metallicities, are responsible for disk and halo ionization, and our result for NGC 891 would therefore be due to propagation effects. To understand whether such a scenario is viable, we combined the neon ratio with optical line ratios along a long slit running through the IRS pointings, and modeled line emission from the DIG layer using the 2D Monte-Carlo radiative transfer code of \\citet*{2004MNRAS.348.1337W}. The code is also capable of 3D simulations. The models include the effects of radiation field hardening and explore various temperatures and luminosities of the radiation field, and were tailored to the diffuse ISM density distribution of NGC 891. As applied to optical line ratios for the Reynolds layer and the halo of NGC 891, WM had found that spectral hardening reduced the need for a non-ionizing heat source, yet it was still not possible to match all the line ratio data, so that additional heating and/or a secondary ionization source still seems to be required. With the neon ratio included, RWB found that no model could reproduce the observed low values of this ratio and its rise with $z$; neither could they predict a rising [O$\\,$III]/H$\\beta$, or the typical values of [O$\\,$I]/H$\\alpha$, and He$\\,$I/H$\\alpha$. They are more successful in reproducing the rise of [S$\\,$II]/H$\\alpha$ and [N$\\,$II]/H$\\alpha$, at least semi-quantitatively. We emphasize that including a non-ionizing heat source or a metallicity gradient should affect the optical line ratios significantly but not the neon ratio. These results are obviously problematic for such photo-ionization models. Clearly, one wishes to know whether the values of the neon ratio measured by RWB, and their enhancement off the plane, are typical of DIG layers, whether this infrared ratio reflects trends in optical ratios, especially [O$\\,$III]/H$\\beta$ (the second ionization potential of oxygen is 35.1 eV, nearly as high as that of neon), and whether the neon ratio in general presents such difficulties for photo-ionization models. We have therefore extended our observations to $z=2$ kpc on both sides of the plane in NGC 891 (where [O$\\,$III]/H$\\beta$ rises more dramatically), and observed disk and halo positions in two other edge-ons, NGC 5775 and NGC 3044 [we use the same distances, 24.8 Mpc and 16.1 Mpc, respectively, as in \\citet{2000ApJ...536..645C} and \\citet{1997ApJ...490..247L})], where coincident optical spectra along slits perpendicular to the disk are also available. The former galaxy is more actively star forming than NGC 891 (\\citealt{1996ApJ...462..712R}), while the long-slit spectroscopy indicates a strongly rising [O$\\,$III]/H$\\alpha$ with $z$ (\\citealt{2000ApJ...537L..13R}). By contrast, as mentioned above, NGC 3044 shows a falling [O$\\,$III]/H$\\beta$, especially for one long slit whose location is the focus of our IRS observations. It also has a relatively bright extraplanar DIG layer (\\citealt{2003A&A...406..493R}). We reconsider photo-ionization models in light of our results, focusing on NGC 891. We begin with a general exploration of parameter space using the CLOUDY code \\citep{1998PASP..110..761F}, asking the simple question of whether there is any combination of ionization parameter, radiation temperature and gas temperature that can reproduce the infrared and optical line ratios in NGC 891 independently at each height. That exploration opens up a wide range of parameter space with a new set of constraints to consider in physical models, and suggests a tractable way to explore multiple ionization sources in future work. Consistent with these findings, we present an initial exploration of this space with new Monte Carlo radiative transfer models that include a second, vertically extended stellar ionizing source. \\subsection{Dusty Halos} The dust content of gaseous halos can provide clues to their origin. There is plenty of evidence for dusty halos through both emission (e.g., \\citealt{2009MNRAS.395...97W}) and absorption (e.g., \\citealt{1999AJ....117.2077H}). That dusty halos are powered by star formation driven disk-halo cycling of gas is suggested by the correlation among galaxies between the presence of extraplanar H$\\alpha$ extinction and emission (e.g., \\citealt{1999AJ....117.2077H}). It seems most likely that dust is transported upwards with the gas. However, other explanations for halo dust are possible (see \\citealt{1997AJ....114.2463H} for a general discussion), including radiative acceleration of grains (\\citealt{1991ApJ...366..443F}), which may lift clouds to heights of a few hundred pc through gas-dust coupling. It has also been argued that grain lifetimes are much shorter than the typical residence time in the ISM (\\citealt{2003ARA&A..41..241D}), implying much dust formation and modification in the ISM. It seems more likely that such dust formation would take place in the denser disk environment than in halos, although the energetic processes involved in disk-halo flows may modify dust properties. This latter point may be important for understanding dust evolution and provides motivation for studies such as ours. The IRS provides the ability to study emission from Polycyclic Aromatic Hydrocarbons (PAHs; \\citealt*{1984A&A...137L...5L,1985ApJ...290L..25A,2001ApJ...556..501B}; \\citealt{2001ApJ...560..261B}; \\citealt*{2007ApJ...657..810D}) - molecules up to about 20$\\AA$ in size containing up to a few thousand C atoms which in SINGS galaxies are found to contribute a few percent of the dust mass (\\citealt{2007ApJ...663..866D}). Emission features in the $10-20$ $\\mu$m range are thought to arise from more neutral PAHs relative to features in the $6-9$ $\\mu$m range (\\citealt{2007ApJ...657..810D}; \\citealt{2008ApJ...679..310G}). It is uncertain how PAHs relate to the larger dust grains, but it has become clear in recent years that PAH emission is suppressed relative to the emission from larger grains in the immediate vicinity of recently formed stars (e.g., \\citealt{2007ApJ...665..390L}), a result attributed to destruction of PAH molecules by UV radiation (e.g., \\citealt{2008ApJ...682..336G}). A key piece of evidence for this picture is the decrease of PAH Equivalent Widths (EW) with increasing radiation field hardness (at least in the $20-40$ eV range), as measured by [Ne$\\,$III]/[Ne$\\,$II], in star forming regions (\\citealt{2006A&A...446..877M}; \\citealt{2006ApJ...639..157W}; \\citealt{2008ApJ...678..804E}), although there seems to be no correlation below [Ne$\\,$III]/[Ne$\\,$II]\\ $\\sim$ 1 (\\citealt{2006ApJ...653.1129B}; \\citealt{2008ApJ...682..336G}), the regime that will be relevant here. In halos, radiation fields are much weaker, whatever their hardness, and it is less likely that processing of grains by radiation will be important. Furthermore, it should be pointed out that in the general ISM as well as in halos, radiation field intensities are low enough so that even the continuum in the $10-20$ $\\mu$m range should be dominated by grains heated by single photons rather than ones in thermal equilibrium, according to \\citet{2007ApJ...657..810D}. We should therefore expect that the PAH spectrum should not vary significantly with radiation hardness or intensity for the kinds of environments relevant here. Rather, relative variations in PAH emission are likely to reflect real changes in the PAH population, and would point to other energetic processes likely to be relevant for halos, such as the aforementioned photo-levitation, or shocks, where PAHs may be produced (by grain collisions or sputtering; \\citealt*{1996ApJ...469..740J}), as well as destroyed \\citep*{2010A&A...510A..36M}. Imaging is beginning to reveal extended halos of PAHs in edge-on galaxies (e.g., \\citealt{2009MNRAS.395...97W} for NGC 891, \\citealt{2009MNRAS.396.1875H} for NGC 5775). Spectroscopy reveals more detailed information but fewer studies exist. In RWB our limited spectra provided a first look at scale heights and EWs of PAH features in NGC 891. Elsewhere, the halo of the starburst M82 has also been studied spectroscopically, first by \\citet{2006ApJ...642L.127E}. More recently, \\citet{2008ApJ...679..310G} find that the intensity ratios of both the 6.2 and 7.7 $\\mu$m features relative to the 11.2 $\\mu$m feature decrease in this halo, which they interpret as a drop in the contribution from ionized PAHs. Otherwise, they find the mix of PAHs in M82 is quite invariant. In the same galaxy, \\citet{2008ApJ...676..304B} find higher EWs of neutral PAHs, but not ionized PAHs, in the halo than in the disk, but they attribute this to a drop in the continuum longward of 10 $\\mu$m relative to the PAH strength. With our larger set of IRS spectra we can start to examine these issues in the halos of more normal, albeit still quite actively star forming, edge-on galaxies, and do so in more depth than was possible in RWB. Scale heights of PAH emission can be compared with those of other vertically extended components to relate PAH and gaseous halos. Scale heights of various features and disk-halo contrasts of their EWs can be compared to search for modification of the PAH population with height. ", "conclusions": "We have presented $10-20$ $\\mu$m spectroscopy of the disks and halos of the edge-on galaxies NGC 891, NGC 5775 and NGC 3044. Regarding the gas-phase lines, our main result is that [Ne$\\,$III]/[Ne$\\,$II] is higher in the halos of NGC 891 and NGC 5775 than in the disks. Scatter in NGC 3044 prevents any trend from being seen. These results exacerbate the problem of explaining DIG line ratio behavior with simple photo-ionization models featuring a thin disk of massive stars. Focusing on NGC 891, where we have the most observational constraints, we have explored parameter space in CLOUDY photo-ionization models to try to determine, as a function of height, what combination of radiation temperature (assuming blackbodies for simplicity), ionization parameter and gas temperature may reproduce the observed IR and optical line ratios. We find that a dramatic rise in radiation temperature with $z$ is required, as well as a more modest rise in gas temperature and a fall in ionization parameter. Not all line ratios can be reproduced without considering more complex models, however. We then considered representative physical photo-ionization models incorporating a thin disk of massive stars in combination with a thick disk of hot (50 kK), presumably evolved stars. We found that the neon ratio behavior in particular could be approximately reproduced for such a thick disk with scale height of order 1 kpc and a contribution of $5-7.5\\%$ of the ionizing radiation. Whether such a component exists in NGC 891 (or even the Milky Way) remains very unclear, but it cannot yet be ruled out based on space density, temperature and scale height arguments. In our models, such a component is also reasonably successful in reproducing the rising [O$\\,$III]/H$\\beta$ and approximate He I/H$\\alpha$ values in this galaxy. However its inclusion does introduce problems with other line ratios. A secondary stellar source is not the only way of increasing radiation hardness with height, and our exploration of parameter space with CLOUDY suggests a feasible way in which other such sources, including hard EUV/X-ray radiation, may be included in future work. PAHs in NGC 891 and NGC 5775 form a vertically extended distribution, with emission scale heights larger in NGC 5775, as is the case for the HI gas. The PAH emission scale heights are comparable to, or somewhat less than, the HI values. These results suggest that PAHs participate in the active disk-halo flows in these two galaxies. In NGC 891, the scale heights are larger than the 8 $\\mu m$ scale height, suggesting a transition to more neutral, and possibly larger, grains with $z$. %Although far from conclusive, shock speeds for which these %processes are plausible candidates roughly agree with speeds invoked %to explain the behavior of the [O$\\,$III]/H$\\beta$ ratios in these %two galaxies. Most PAH equivalent widths are higher in the halos than in the disks, again indicating some slight modification in the PAH population in the halo environment. Shocks and radiative acceleration are processes that may affect PAH populations in halos. The 17 $\\mu$m H$_2$ line is detected in all pointings, indicating a warm molecular gas component with a large vertical extent, reaching as high as $z=2$ kpc. The greater scale height of this emission in NGC 5775 {\\it vs.} NGC 891 again suggests a connection with disk-halo flows. Column densities and temperatures cannot be derived without measurement of additional transitions, however. Future observations will no doubt reveal the fraction of halo gas mass in this phase, and its importance for understanding the ISM of halos. Finally, we wish to stress here the importance of the neon ratio in understanding the energetics of the diffuse ISM, and the great advantage it holds over optical diagnostics. Future instruments capable of sensitive, high spatial and spectral resolution observations of these and other mid-IR lines should be a strong priority as these lines reveal a wealth of information. %% Included in this acknowledgments section are examples of the %% AASTeX hypertext markup commands. Use \\url without the optional [HREF] %% argument when you want to print the url directly in the text. Otherwise, %% use either \\url or \\anchor, with the HREF as the first argument and the %% text to be printed in the second." }, "1101/1101.4050_arXiv.txt": { "abstract": "{ Cluster media are dynamical, not static; observational evidence suggests they are turbulent. High-resolution simulations of the intracluster media (ICMs) and of idealized, similar media help us understand the complex physics and astrophysics involved. We present a brief overview of the physics behind ICM turbulence and outline the processes that control its development. High-resolution, compressible, isothermal MHD simulations are used to illustrate important dynamical properties of turbulence that develops in media with initially very weak magnetic fields. The simulations follow the growth of magnetic fields and reproduce the characteristics of turbulence. These results are also compared with full cluster simulations that have examined the properties of ICM turbulence. ", "introduction": "Observation and theory have revealed intracluster media (ICMs) to be very dynamic environments with active ``weather'' driven by a host of activities such as mergers, accretion, AGNs, galactic winds and instabilities. These drivers are common and cause the ICMs to be criss-crossed by large-scaled, complex flows that generate shocks, contact discontinuities (aka ``cold fronts'') and bulk shear. Inevitably such flows should lead to turbulence in the ICMs, an outcome supported by growing observational evidence. These include, for example, substantial ICM random velocities in Perseus reducing resonance scattering in the 6.7 keV iron line \\cite{chur04}, evidence for thermal ICM pressure fluctuations in the Coma cluster \\cite{schu04}, patchy Faraday rotation measure distributions in several clusters \\cite{bona10} and the absence of large scale polarization in cluster radio halos \\cite{kim90}, suggesting disordered magnetic fields. Turbulence in clusters is important to understand for many reasons. Turbulent pressure helps support the ICM, so relevant cluster mass measures. Turbulence transports entropy, metals and cosmic rays, all important cluster diagnostics. It transports and amplifies magnetic fields, which in turn control ICM viscosity, resistivity and thermal conductivity, as well as the propagation and acceleration of cosmic rays. The literature on turbulence is extensive including excellent reviews on MHD turbulence, which is most relevant to the ICM (e.g., \\cite{brand05}). Here we make a few observations pertinent to this meeting. ", "conclusions": "Processes such as shocks and outflows are likely to drive turbulence in ICMs. The detailed physics is difficult to model analytically, but simulations allow us to explore it in some detail. Magnetic fields are integral ingredients of both the microphysics of ICM transport properties and essential players in the large scale dynamics. Simulations are revealing important insights into the character of the turbulence, including properties of magnetic fields." }, "1101/1101.1991_arXiv.txt": { "abstract": "Muon neutrino astronomy is drown within a polluted atmospheric neutrino noise: indeed recent ICECUBE neutrino records at (TeVs) couldn't find any muon neutrino point source \\cite{01} being blurred by such a noisy sky. However at $24$ GeV energy atmospheric muon neutrinos, while rising vertically along the terrestrial diameter, should disappear (or be severely depleted) while converting into tau flavor: any rarest vertical $E_{\\mu}\\simeq 12$ GeV muon track at South Pole Deep Core volume, pointing back to North Pole, might be tracing mostly a noise-free astrophysical signal. The corresponding Deep Core $6-7-8-9$ channels trigger maybe point in those directions and inside that energy range without much background. Analogous $\\nu_{\\mu}$ suppression do not occur so efficiently elsewhere (as SuperKamiokande) because of a much smaller volume, an un-ability to test the muon birth place, its length, its expected energy. Also the smearing of the terrestrial rotation makes Deep Core ideal: along the South-North Pole the solid angle is almost steady, the flavor $\\nu_{\\mu}\\mapsto \\nu_{\\tau} $ conversion persist while the Earth is spinning around the stable poles-axis. Therefore Deep Core detector at South Pole, may scan at $E_{\\nu_{\\mu}} \\simeq 18-27$ GeV energy windows, into a narrow vertical cone $\\Delta \\theta \\simeq 30^{o}$ for a novel $\\nu_{\\mu}$, $\\bar{\\nu}_{\\mu}$ astronomy almost noise-free, pointing back toward the North Pole. Unfortunately muon (at $E_{{\\mu}}\\simeq 12$ GeV) trace their arrival direction mostly spread around an unique string in a zenith-cone solid angle. To achieve also an azimuth angular resolution a two string detection at once is needed. Therefore the doubling of the Deep Core string number, (two new arrays of six string each, achieving an average detection distance of $36.5$ m), is desirable, leading to a larger Deep Core detection mass (more than double) and a sharper zenith and azimuth angular resolution by two-string vertical axis detection. Such an improvement may show a noise free (at least factor ten) muon neutrino astronomy. This enhancement may also be a crucial probe of a peculiar anisotropy foreseen for atmospheric anti-muon, in CPT violated physics versus conserved one, following a hint by recent Minos results. \\vspace{1pc} ", "introduction": "Neutrino Astronomy is a hard and novel view of the Universe mostly ruled, at lowest energy, by solar MeV electron neutrino signal. At tens MeV a neutrino astronomy occur by rarest (nearly one a century) galactic Supernova events. At higher energies (GeVs, TeVs) the neutrino flux, detectable at best as muons, is drawn and smeared by an overabundant homogeneous atmospheric $\\nu$ background. They exist with high rate because their parent charged Cosmic Rays, C.R., while reaching the Earth, are bent and spread by stellar and galactic magnetic fields. Moreover for the same argument CR, while propagating randomly and twisted in space, are surviving much longer than direct photons or neutrino tracks. The CR flux (except maybe ZeV ones) is thus several order of magnitude more abundant than neutral gamma or neutrinos one (even if they were born at nearly the same rate). This is manifest in recent ICECUBE featureless records for TeVs neutrinos have (unfortunately) shown \\cite{01}. Other astrophysical sources, commonly offering a weak neutrino signal, are hard to be disentangled from such a noisy atmospheric (Cosmic Ray secondary) $\\nu$ background. If the primary source neutrino spectra is hard (for instance as Fermi suggested by $\\Phi_{\\nu} \\simeq E^{-2}$) than the atmospheric $\\nu$ background, $\\Phi_{\\nu} \\simeq E^{-2.7}\\rightarrow E^{-3.7} $, at energies $E \\geq 10^{14}$eV or $E \\geq 10^{15}$eV , atmospheric neutrino noise may be finally overcome by astrophysical signal. However their flux at those high energies are depleted and too low to be easy observed. At even highest energies, EeV, the tau \\cite{17} neutrino astronomy may also rise via up-going tau airshowers, possibly soon in Auger or T.A. Fluorescence Telescopes \\cite{03}\\cite{Auger08} . Consequently for the moment it maybe also important to reveal any muon neutrino signal at low energies in cleaned or filtered (from the atmospheric $\\nu$ background) sky: around $E_{\\nu} \\simeq 24$ GeV energy up-going muon neutrinos inside a $\\theta \\simeq 20-30^{o}$ cone pointing to North Pole are offering such a tuned noise-free $\\nu$ view. Any upgoing muon clustering in Deep Core at those $6-9$ channels \\cite{10},\\cite{22},\\cite{15} maybe much better revealed in next a few years. \\subsection{Gamma, Neutrino and Cosmic Rays} The role of radiations and particles in the Universe maybe summarized by a wide spectra , see Fig. \\ref{01}, see also \\cite{03b}. Most of us are waiting for an astrophysical signal at highest energies, PeVs up to EeVs as parasite secondaries of UHECR (GZK cut off, respectively for cosmogenic neutrinos by UHECR nuclei or nucleon \\cite{03a}), see also \\cite{Auger08}, \\cite{Auger10}; in Fig. \\ref{01} one see a narrow shadow window where $\\nu_{\\mu}\\mapsto\\nu_{\\tau} $ (the oscillating colored curve below an average atmospheric neutrino muon flux). In that window the absence of atmospheric muons $\\nu_{\\mu}$ favors a better noise free astrophysical view of the Universe. \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{AllSpectra.eps} \\caption{ The wide view flux number of radiation and cosmic rays. The integral flux number is shown in usual unity. The parasite atmospheric neutrinos and their oscillation \\cite{12},\\cite{19},\\cite{02} into tau are shown, in logarithmic scale. The vertical muon disappearance at $E_{\\nu_{\\mu}} \\simeq 20-28$ GeV is shown by a gray band. The neutrino oscillation role for atmospheric tau neutrino is drawn (while the corresponding muon one is not, to avoid confusion). There are two $\\nu_{\\tau}$ curves; the fast decreasing one related to horizontal $\\nu_{\\tau}$, and the vertical up-going curve reaching a maxima in the shaded area. } \\label{01} \\end{figure} \\subsection{Neutrino Rate}The expected number of muons produced by up-going $\\nu_{\\mu}$,$\\bar{\\nu}_{\\mu}$, fully contained and partially contained are derived extrapolating by size ratio SuperKamiokande \\cite{02a} events versus Deep Core effective mass, respectively at $15-25$ GeV energy band where most of the up-going atmospheric $\\nu_{\\mu}$, $\\bar{\\nu}_{\\mu}$ conversion into $\\nu_{\\tau}$, $\\bar{\\nu}_{\\mu}$ takes place \\cite{19},\\cite{12},\\cite{02}. The Fully Contained events in SK cannot account for most of these events because the $ \\mu$ tracks are too long to be totally contained inside the SK $\\simeq 40$ m height (out of very rare inclined upward trajectories). Therefore most of the events are based on Partially contained (PC) and Upward (UP) and Through going $\\mu$ tracks \\cite{02a}. The corresponding event rate a year are (for a nominal 4.8 Mton Deep Core effective mass in that energy range $25\\geq E \\geq 16$ GeV) within a vertical cone of $33^{o}$ opening angle as they have been recently reported \\cite{02}.The tracks by nearly horizontal muons will excite the vertical string with a characteristic arrival time similar to the vertical shower event or up-going vertical muon about five GeV. Indeed the time difference in arrival for spherical shower along a string (each DOM at $7$ m separation) is nearly $\\Delta t_{0}\\simeq t_{0}= h/c = 23 ns$; by triangulation any horizontal muon tracks and its Cherenkov cone will record a similar delay $\\Delta t_{0}\\simeq t_{0}\\cdot\\cot(\\theta_{C}) (1-\\frac{n_{ice}}{\\cos(\\theta_{C})}) \\simeq 1.03 t_{0} = 24 ns$ between two nearby phototube (DOM). This delay is due to the superior region of Cherenkov cone illuminating the phototube from below. This delay should not be confused with the other one discussed in next section. Therefore the $3-4-5$ channel might be polluted by horizontal muon and by shower originated by NC and by electron charged events with a very similar signature. These crowded low energy edge cannot be useful is neutrino astronomy. For a summary of the neutrino muon suppression along different channel group (see Fig. \\ref{08}). \\subsection{Zenith angle via timing scale} To test the arrival muon direction by an unique string at twenties GeV range one may exploit the Cherenkov signal timing train of events along the string, event due to the different geometry of Cherenkov light arrival along the muon track. This time delay by an arrival muon angle $\\theta$ (constrained within ($\\theta_{max} \\simeq 48.75^{\\circ}$)), complemental to Cherenkov angle ($\\theta_{Ch}\\simeq 41.25^{\\circ}$), is due to different path of the light flight toward the phototube. Its value is: \\begin{equation} \\delta t = \\frac{h}{c} \\left( \\frac{Sin(\\theta_C+\\theta)-n \\cdot Sin(\\theta)}{Cos(\\theta)\\cdot Sin(\\theta_C+\\theta)} \\right) \\end{equation} where h is the phototube distance ($h= 7$ m), n is the refractive index in ice, $\\theta_C$ is the Cherenkov angle in ice. \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{timing.eps} \\caption{ The delay time between two nearby consecutive DOM due to an inclined arrival muon at zenith angle $\\theta$ (angle between the vertical axis and the muon axis direction assumed coplanar with the string line). The continuous curve is the exact function eq.1, the dashed line is the linear approximation. The nearly linear correlation allow to estimate the zenith angle by such delay scale among the phototube detection, as in eq.2. } \\label{02} \\end{figure} The linear behavior shown in graph \\ref{02} can be approximately expressed by the following equation: $$\\delta t \\simeq 2.2\\cdot 10^{-8} \\left(1- \\left(\\frac{\\theta}{48.75 ^{\\circ}}\\right)\\right) s$$ From here we may express the arrival zenith angle as: \\begin{center} \\begin{equation} \\theta \\simeq 48.75 ^{\\circ} - \\frac{ \\delta t}{2.2 \\ 10^{-8}} \\end{equation} \\end{center} The characteristic channel exited by such twenties GeV neutrinos are 6-7-8-9; these 5-8 pairs offer a clear timing measure whose average value may strongly constrain the zenith muon angle, as shown by previous formula and graph. The validity of last approximation is within $ \\theta \\leq 30^{\\circ}$, also because time resolution of Deep Core array. \\subsection{Muon survival probability} Following our recent articles \\cite{02} the oscillating neutrino flavor offer different reading chart: the $\\nu_{\\mu}$ survival probability as a function of the arrival angle at given energy (mainly the most suppressed one at $20.5$ GeV),(see Fig. \\ref{03}); the additional view of the $\\nu_{\\mu}$ survival probability as a function of the distances (see Fig. \\ref{04}); the $\\nu_{\\mu}$ survival probability as well as the complemental $\\nu_{\\tau}$ appearance probability as a function of the energy crossing the Earth diameter (see Fig. \\ref{05}). In that figure one may observe the CPT violated scenario whose oscillation may be opposite to common CPT conserved one. Read more details in \\cite{02}. The lower energy band where the $\\nu_{\\mu}$ survival probability may be suppressed (at inclined-horizontal directions) as a function of the zenith angle is shown in (see Fig. \\ref{06}); different argument make unrealistic the use of such a clean sky, mostly polluted by horizontal muons and additional noises. A final $\\nu_{\\mu}$ survival probability is described for the higher energy (above $30$ GeV) where the conversion and suppression became smaller and smaller, making the filter of atmospheric noise almost useless.(see Fig. \\ref{07}) \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{20GeV_6-6ch.eps} \\caption{ The probability of $\\nu_{\\mu}$ survival as a function of the angular arrival direction, crossing the Earth, for an average $\\nu_{\\mu}$ energy $E_{\\nu_{\\mu}}\\simeq 20.5$ GeV. The role of the matter density (respect the vacuum) inside the Earth has a negligible role.} \\label{03} \\end{figure} \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{polosud-24_6.eps} \\caption{ As above the same probability of $\\nu_{\\mu}$ survival as a function of the distance across the Earth at $E_{\\nu_{\\mu}}\\simeq 24.6$ GeV, in vacuum. The dashed areas label the region where the suppression is more than one order of magnitude, i.e. where the sky is more clean from any atmospheric neutrino noise.} \\label{04} \\end{figure} \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{MU.eps} \\caption{The survival probability for muon neutrinos and the complementary tau conversion in CPT conserved model and in the new Minos CPT violated scenario. The probability is described as a function of the energy both for the mixing in vacuum and in Earth. The dashed area shows the energy windows where the neutrino astronomy maybe enhanced. } \\label{05} \\end{figure} \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{13GeV_4ch.eps} \\caption{As above at different energy windows, and at different solid angle region, where atmospheric muon (in CPT conserved scenario) are almost suppressed. This energy range corresponds to nearly $6.5$ GeV muon whose track, almost of $30$ m length, possibly contained and measured in SK. These signals maybe searched also in SK, but they are too rare because of the small size of SK (a few or ten event a year) and nearly horizontal, polluted by direct horizontal downward muons. Moreover in Deep Core these nearly horizontal muons will excite the vertical string with a characteristic arrival time similar to the vertical shower event or upgoing vertical muon about five-six GeV, made by $12$ GeV vertical $\\nu_{\\mu}$. In Deep Core these $13$ GeV signals (silent but horizontal) are very difficult to disentangle within the extremely abundant and polluted shower events (tens of thousands of event a year or more) by muon and tau neutral current interactions and also because of the up-going $5-6$ GeV (atmospheric muon) arrival made by $12$ GeV vertical $\\nu_{\\mu}$. Therefore these $3-4-5$ channel of events in Deep Core, might be extremely polluted and are useless to astronomical study.} \\label{06} \\end{figure} \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{32GeV_10-4ch.eps} \\caption{The survival probability for muon neutrinos as different energy windows, and at different solid angle region, where atmospheric muon (in CPT conserved scenario) are only partially suppressed ($20\\%$). In Deep Core these $32$ GeV astrophysical neutrino events maybe already sink in dominant polluting atmospheric signals, making difficult to disentangle any clear astronomy. At larger and larger energies the probability suppression fade away as well as the possibility to filter and cancel the atmospheric neutrino noise.} \\label{07} \\end{figure} \\begin{figure*}[htb] \\vspace{9pt} \\includegraphics[width=140mm]{Rate.eps} \\caption{The rate of upgoing muons based on SK rate and extrapolated to Deep Core, assuming a vertical cone view within $\\sim 33^{o}$. The rate is mostly based on PC,Upward stopping and Upward through-going signal in SK. The expected event rate in the narrow red area is strongly modulated in an anisotropy due to the nearly total flavor conversion. The muon suppression may reach at least a factor $10$ for an accuracy spread in the muon energy (and its length) : $\\frac{\\Delta E_{\\nu_{\\mu}}}{E_{\\nu_{\\mu}}}\\simeq 0.1$ ; see the suppression factor in channels $6-8$ that is reducing to a few hundred ($100-200$) event a year of the atmospheric muon noise. Any astrophysical source may better rise and sharply cluster around source in this energy-angular silent cone of view.} \\label{08} \\end{figure*} \\clearpage ", "conclusions": "$ astronomy at $20$ GeV} The muon neutrino almost complete conversion at Deep Core along vertical axis into tau, offer a rare opportunity to use this energy range and that sky view to search for astrophysical neutrino sources. The possibility to test the exact arrival direction by an unique string is poor: only the zenith angle may be found following eq.1,2. To obtain at twenty GeV neutrino direction (and a larger detector effective mass) we suggest the doubling of the Deep Core string: two contemporaneous string detection will mark zenith and azimuth muon (and neutrino) vector, opening the road to a sharp neutrino astronomy. \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=75mm]{Constellations.eps} \\caption{The stellar constellation sky, in galactic coordinates, pointing to the terrestrial North, where the muon disappearance at $\\simeq 25$ GeV occurs, as it maybe observed by Deep Core. The spread spinning sky may simulate the Deep Core ability to somehow disentangle the zenith angle (inner to outer rings respectively corresponding to channel 9-8-7-6) , but un-ability to fix the exact azimuth muon arrival direction, being the signal projected along the string axis in a unique conic solid angle. } \\label{09} \\end{figure} \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{northPole.eps} \\caption{As above the Very High Energy gamma sources and sky with the marked North sky area. Also the recent 69 UHECR events by AUGER have been shown, mostly in the South sky, where Argentina sky is, \\cite{Auger08}, \\cite{Auger10}. } \\label{10} \\end{figure} The angular resolution, the muon track detection and the energy estimate may offer an additional road to test muon suppression, tau appearance as well as eventual CPT violated mass terms. \\cite{14},\\cite{02}. The North sky may show the persistence of known VHE gamma sources also in neutrino form: the flaring of gamma sources observed by Magic, Hess, Veritas or Fermi satellite (as sources 0502+675, 0716+714,0710+591, 1959 + 650, as well as M82) may shine in this exciting and silent muon neutrino Northern sky in a very few years . \\begin{figure}[htb] \\vspace{9pt} \\includegraphics[width=70mm]{vhe5.eps} \\caption{As above the Very High Energy gamma sources and sky with the marked North sky area. Different extragalactic sources are labeled, following recent (Nov. 2010) record by Cherenkov Telescopes and Fermi satellite. The North sky may show the persistence of VHE gamma sources also in neutrino form: Magic,Hess,Veritas sources as 0502+675, 0716+714,0710+591 as well as M82, 1959 + 650, may shine in this (not just cold, but cool) Northern sky } \\label{11} \\end{figure}" }, "1101/1101.6055_arXiv.txt": { "abstract": " ", "introduction": "The process of accretion that drives most X-ray astrophysical phenomena can often be dramatic and short lived, with increases in accretion rates causing X-ray flux increases of up to 6 orders of magnitude from quiescent levels. In many cases these events lead to the discovery of previously unknown systems, or systems that were previously considered uninteresting. The flare-up of a Galactic X-ray transient typically heralds a rapid increase in the rate of accretion onto a compact object (white dwarf, neutron star or black hole), and provides an ideal laboratory for studying astrophysics in a relativistic regime. Even though transients have been studied for many years, our understanding of the processes behind extreme accretion events remains relatively poor. X-ray transients also cover a wide range of different system phenomenology, including Black Hole binaries systems (e.g.\\ Remillard \\& McClintock 2006), Low Mass X-ray Binaries, High Mass X-ray Binaries (HMXB), millisecond pulsars (e.g.\\ Campana et al.\\ 2008), Cataclysmic Variables, Novae and Supergiant Fast X-ray Transients (e.g.\\ Romano et al.\\ 2010). These transient events are rare, and often short lived, making detection and detailed study difficult. To obtain a good rate of detection of transient outbursts, X-ray instruments that cover very large areas of the sky are required. However, these wide field instruments typically lack the spatial resolution required to provide accurate localizations necessary for further optical and IR observations, and often do not have enough spectral sensitivity for a detailed analysis of the characteristics of the outburst. Therefore follow-up observations with complementary observatories, such as {\\em Swift}, are required to provide more accurate positions and simultaneously observe the broadband (Optical to Gamma-ray) spectral behavior. ``Monitor of All-sky X-ray Image'' (MAXI, e.g.\\ Ueno et al.\\ 2009), part of the Japanese Experiment Module on the International Space Station, provides a powerful tool for discovery of new X-ray transients. MAXI's ability to perform a near all-sky X-ray image of the sky in the 0.5--20 keV energy band, with sensitivities as low as 60 mCrab (5 sigma) in a single orbit and 15 mCrab in a day, makes it more capable of finding transients than other instruments like {\\em Swift}'s Burst Alert Telescope (BAT; Barthelmy et al.\\ 2004) which has too hard an energy band, or RXTE PCA Galactic Bulge scans (limited spatial and temporal coverage).9 MAXI's capability of finding transient X-ray phenomena has been proven in the first year of operations by the detection of outbursts of known sources, GRBs, and most recently the discovery of a new Galactic transients: MAXI J1659$-$152 and MAXI J1409$-$619. We report here on the results of a program to localize MAXI discovered X-ray Galactic X--ray transients utilizing NASA's {\\em Swift} Gamma-Ray Burst Explorer Mission (e.g.\\ Gehrels et al.\\ 2004). {\\em Swift} has proven capabilities in the follow-up of transient sources: fast slewing, flexible scheduling, and the ability to perform rapid TOO observations of targets through spacecraft commanding as quickly as within 1 hour of the announcement of a target's location. Recent developments of the {\\em Swift} planning infrastructure have been driven towards making these observations both quicker and less burdensome on the small {\\em Swift} mission operations team. {\\em Swift}'s X-ray Telescope (XRT; Burrows et al.\\ 2004) has a field of view of approximately $23.6'$ diameter, well matched to the typical $0.2$ degree error circle from MAXI. The {\\em Swift} UV/Optical Telescope (UVOT; Roming et al.\\ 2005) here also provides two valuable services in support of these observations: broad band optical/UV observations of the optical counterpart of the new transient (if visible); and astrometric correction of XRT data that allows X-ray localization errors to be reduced to as little as 1.5 arc-seconds radius (90\\% confidence). \\begin{figure*}[t] \\centering \\psbox[xsize=12cm] {HD347929.ps} \\caption{\\label{fi:HD347929local}The XRT field of view of HD 347929, with MAXI transient error circle marked. This observation is typical of what is seen in this program for other follow-ups.} \\end{figure*} ", "conclusions": "We have presented the results of a {\\em Swift} program to rapidly follow-up and localize MAXI discovered X-ray transients. Typically, the MAXI localization of a source ($\\sim 0.2$ degrees) is well matched to the XRT field of view, making XRT the ideal instrument to attempt to more accurately localize the transient. {\\em Swift}'s unique observing flexibility also allows us to rapidly follow-up and report on the detection of these transients, usually within 24 hours of being notified of the transient by the MAXI team. {\\em Swift} is capable of localizing transients to an accuracy of up to 1.5 arc-seconds radius (90\\% confidence), or even more accurate if a UVOT counterpart is found. We have reported on 6 triggers of the {\\em Swift}/MAXI transient follow-up program. Two of those triggers localized and confirmed the present of a previously unknown transient source, one (MAXI J1659$-$152) is a new blackhole binary system. In both cases {\\em Swift} follow-up and monitoring observations of these targets was triggered, giving valuable insight into the nature of these transients. For the other 3 triggers {\\em Swift} was able to positively identify the source of the MAXI trigger. In only 1 case was the result of MAXI follow-up inconclusive (SAX J1542.8$-$5949). The unique complementary capabilities of {\\em Swift} and MAXI has proven to be well matched in the search for and localization of new Galactic transients sources. The all-sky monitoring of MAXI, combined with {\\em Swift}'s rapid response TOO capability and the convenient similarity of the XRT's field of view with the typical MAXI error circle, has proved to be very successful in accurately localizing new transient sources, which provides an important service to the community." }, "1101/1101.3670_arXiv.txt": { "abstract": "A search for a diffuse flux of astrophysical muon neutrinos, using data collected by the ANTARES neutrino telescope from December 2007 to December 2009 is presented. A $(0.83\\times 2\\pi)$ sr sky was monitored for a total of 334 days of equivalent live time. The searched signal corresponds to an excess of events, produced by astrophysical sources, over the expected atmospheric neutrino background without any particular assumption on the source direction. Since the number of detected events is compatible with the number of expected background events, a 90\\% c.l. upper limit on the diffuse $\\nu_\\mu$ flux with a $E^{-2}$ spectrum is set at $E^2\\Phi_{90\\%} = 5.3 \\times 10^{-8} \\ \\mathrm{GeV\\ cm^{-2}\\ s^{-1}\\ sr^{-1}}$ in the energy range 20 TeV -- 2.5 PeV. Other signal models with different energy shape were also tested and some rejected. ", "introduction": "\\label{sec:introduction} The ANTARES high-energy neutrino telescope is a three-dimensional array of photomultiplier tubes (PMT) distributed over 12 lines installed deep in the Mediterranean Sea, each line including 75 PMTs \\cite{Paschal}. A neutrino telescope in the Northern hemisphere includes the Galactic Centre in its field of view and is complementary to the IceCube Antarctic telescope \\cite{Teresa}. The main goal of the experiment is the search for high-energy neutrinos from astrophysical sources. If the sensitivity of point source search techniques is too small to detect neutrino fluxes from individual sources, it is possible that many sources could produce an excess of events over the expected atmospheric neutrino background. In this proceeding the search for very-high energy extraterrestrial muon neutrinos from unresolved sources is presented using data collected by the ANTARES telescope from December 2007 to December 2009. Atmospheric muons and neutrinos are the main sources of background in a neutrino telescope. The former can be suppressed by applying requirements on the direction of the events, % the latter is an irreducible background. As the spectrum of cosmic neutrinos is expected to be harder than that of atmospheric neutrinos, the signal we are looking for corresponds to an excess of high energy events in the measured energy spectrum without any particular assumption on the source direction. Electrons (in the so-called ``leptonic models'') or protons and nuclei (``hadronic models'') can be accelerated in astrophysical processes. Hadronic models \\cite{chiarusi} predict that the energy produced in the sources is carried away by cosmic rays, $\\gamma$-rays and neutrinos. A benchmark flux for the measurement of diffuse neutrinos is the Waxman-Bahcall (W\\&B) upper bound \\cite{wb}. Using the CR observations at $E_{CR}\\sim 10^{19}$ eV ($E_{CR}^2 \\Phi_{CR} \\sim 10^{-8}$ GeV cm$^{-2}$s$^{-1}$sr$^{-1}$) the diffuse flux of muon neutrinos is constrained at the value: \\begin{equation} E^2_\\nu \\Phi_\\nu < 4.5/2 \\times 10^{-8} \\ \\mathrm{GeV\\ cm^{-2}\\ sr^{-1}\\ s^{-1}} \\label{eq:wb} \\end{equation} (the factor 1/2 is added to take into account neutrino oscillations). ", "conclusions": "Using data from 334 days of equivalent live time collected with the ANTARES telescope, a search for a diffuse flux of high energy cosmic muon neutrinos was made. The 90\\% c.l. upper limit on the diffuse $\\nu_\\mu$ flux with a $E^{-2}$ spectrum is set at $E^2\\Phi_{90\\%} = 5.3 \\times 10^{-8} \\ \\mathrm{GeV\\ cm^{-2}\\ s^{-1}\\ sr^{-1}} $ in the energy range 20 TeV -- 2.5 PeV. Other signal models with different energy shape are also tested and some of them excluded at the 90\\% c.l.. \\begin{table}[tb] \\caption{{\\small Tested flux models.}} \\label{tab:models} \\small \\begin{tabular}{ccccc} \\hline Model & R$^*$ & N$_{mod}$ & $\\Delta E_{90\\%}$ & $\\mu_{90\\%}/N_{mod}$ \\\\ & & &(PeV) & \\\\% {\\scriptsize $\\mu_{90\\%}/N_{mod}$} \\\\ \\hline MPR\t\t\t\t& 1.43 & 3.0 & 0.1$\\div$10 & 0.4 \\\\ P96$p\\gamma$\t& 1.43 & 6.0 & 0.2$\\div$10 & 0.2 \\\\ S05\t\t\t\t& 1.45 & 1.3 & 0.3$\\div$ 5 & 1.2 \\\\ SeSi\t\t\t\t& 1.48 & 2.7 & 0.3$\\div$20 & 0.6 \\\\ M$pp+p\\gamma$\t& 1.48 & 0.24 & 0.8$\\div$50 & 6.8 \\\\ \\hline \\end{tabular} {\\small Astrophysical flux models, the value of the R$^*$ which minimizes the MRF, the expected number of events $N_{mod}$, the energy range $\\Delta E_{90\\%}$ in which the 90\\% of events are expected, and the ratio $\\mu_{90\\%}/N_{mod}$. See \\cite{papero_diffusi} and references therein. } \\end{table}" }, "1101/1101.0997_arXiv.txt": { "abstract": "{In order to accurately model giant planets, a whole set of observational constraints is needed. As the conventional constraints for extrasolar planets like mass, radius, and temperature allow for a large number of acceptable models, a new planetary parameter is desirable in order to further constrain planetary models. Such a parameter may be the tidal Love number $k_2$.} {In this paper we aim to study the capability of $k_2$ to reveal further information about the interior structure of a planet.} {With theoretical planetary models we investigate how the tidal Love number $k_2$ responds to the internal density distribution of a planet. In particular, we demonstrate the effect of the degeneracy of $k_2$ due to a density discontinuity in the envelope of a three-layer planetary model.} {The effect of a possible outer density discontinuity masks the effect of the core mass on the Love number $k_2$. Hence, there is no unique relationship between the Love number $k_2$ and the core mass of a planet. We show that the degeneracy of $k_2$ with respect to a layer boundary in the envelope also occurs in existing planets, e.g. \\object{Saturn} and the Hot Neptune \\object{GJ\\,436b}. As a result of the degeneracy, the planetary parameter $k_2$ cannot be used to further constrain Saturnian models and for GJ\\,436b only a maximum possible core mass can be derived from a given $k_2$. To significantly narrow the uncertainty about the core mass of GJ\\,436b the combined knowledge of $k_2$ and atmospheric metallicity and temperature profile is necessary.} {} ", "introduction": "Ever since the first extrasolar planet around a solar-type star was detected \\citep{MayorQueloz95}, questions about the composition and origin of extrasolar planetary objects (exo-planets) have been of major interest. Models of exoplanets are often little constrained based on the observable parameters mass, radius, and effective temperature, in particular metal-rich planets \\citep{Adamsetal08}. For solar planets additional constraints are provided by the gravitational moments which have been measured by spacecraft or Earth-based observations of the motion of satellites and hence are not accessible for extrasolar planets. However, a similar quantity does exist: the tidal Love number $k_2$. To first order in the dimensionless number that describes the effect of rigid rotation or degree 2 tidal distortion $k_2$ is equivalent to $J_2$ \\citep[see e.g.][]{Hubbard84}. The tidal Love number $k_2$ is a potentially observable parameter. \\citet{RagozzineWolf09} showed that the dominant source of apsidal precession of Hot Jupiters is the tidal interaction between the planet and its star. This tidally induced apsidal precession creates a unique variation in the transit light curve which is detectable by space-based missions like \\emph{Kepler}. Another possibility of determining $k_2$ is the measurement of the orbital parameters of a two-planet system in apsidal alignment \\citep{Batyginetal09}. Due to tidal dissipation a coplanar two-planet system can evolve into a tidal fixed point which is characterized by the alignment of the apsidal lines \\citep{Mardling07} and both orbits precess with the same rate. \\citet{Batyginetal09} showed that in this state the Love number $k_2$ is a function of the inner planet's eccentricity. Like $J_2$ for the solar system planets, $k_2$, if known, can be used to further constrain the models of extrasolar planets as it is sensitive to the internal density distribution of the planet. Understanding the planetary interior is important for determining not only physical processes but also the formation history. Hence, it is crucial to analyze what information can be extracted from a measured $k_2$ and its implications on the planetary interior. First, we will describe the definition and calculation of the Love numbers in Sect.~\\ref{sec:DefCalc2L}. We also confirm the correlation between the central condensation of a planet and its Love number $k_2$ within a simple two-layer model. In Sect.~\\ref{sec:3L-model} we introduce a more sophisticated three-layer planetary model and demonstrate the degeneracy of $k_2$ with respect to the density discontinuity in the envelope. We apply these results to Saturn and to the Hot Neptune GJ\\,436b in Sect.~\\ref{sec:planets}. The main results of this paper are summarized in Sect.~\\ref{sec:sum}. ", "conclusions": "\\label{sec:sum} In this paper we investigated the effect of the density distribution of a planet on its tidal Love number $k_2$ in order to find out what conclusions can be drawn from a measured $k_2$ on the internal structure of a planet. We confirmed that the Love number $k_2$ is a measure of the level of central condensation of a planet. However, in a three layer model approach $k_2$ is \\emph{not} a unique function of the core mass. There is a degeneracy of $k_2$ with respect to a density discontinuity in the envelope of the planet. It is possible to have several acceptable models for a given $k_2$ value, which can differ significantly in core mass. The effect of the outer density discontinuity on $k_2$ is compensating the effect of the core. Furthermore, we showed that the Radau-Darwin relation is a too crude approxomation to describe the moment of inertia of gas giant planets. We verified our results on $k_2$ with models of existing planets. For Saturn the freedom to place the layer boundary in the envelope leads to a high uncertainty in the core mass. Regardless of the core mass all Saturn models have the same $k_2$, demonstrating the degeneracy caused by the outer layer boundary. For extrasolar planets the Love number $k_2$ can be an equivalent constraint to $J_2$ for the solar system planets. However, one has to be careful with estimates about the core mass derived from $k_2$ as degeneracy may also occur in extrasolar planets. For GJ\\,436b we find a highly degenerate area of $k_2<0.24$ where a measurement of $k_2$ would barely help to further constrain the interior models. Only a \\emph{maximum possible} core mass and for $k_2>0.24$ a large metallicity can be inferred. With additional knowledge about the atmospheric metal abundance the uncertainty about the core mass could be significantly narrowed. Tabulating $k_2$ values of various planetary models can prove to be very useful once $k_2$ is actually measured for extrasolar transiting planets. For instance, for the Super-Earth \\object{GJ\\,1214b} \\citet{Nettelmannetal10b} demonstrated that H/He or water envelopes result in significantly different values of $k_2$. Furthermore, we have shown in this paper that even though the Love number $k_2$ is a degenerate quantity it can help constraining the core mass of a planet. Knowledge about the core masses of planets is highly desired because it is thought to help to distinguish between the possible planet formation scenarios of core accretion \\citep[see e.g.][]{Pollacketal96} and gravitational instability \\citep[see e.g.][]{Boss97}. However, one has to keep in mind that core accretion models can also result in very small cores of 1.7\\,\\ME or 0.25\\,\\ME in the case of grain-free or even metal-free envelopes, respectively \\citep{HoriIkoma10}. On the other hand, gravitational instability models allow the formation of a massive core as well if the protoplanet is cold enough for grain settling to take place \\citep{HelledSchubert08}. A clear distinction between the two formation models can only be made for massive extra-solar giant planets $\\geq$ 5\\,\\MJ. \\citet{HelledSchubert08} showed that such massive protoplanets formed by disk instability cannot build up a core at all due to their high internal temperatures and evaporation of the grains." }, "1101/1101.4670_arXiv.txt": { "abstract": "{We present an optical/infrared study of the dense molecular cloud, L935, dubbed ``The Gulf of Mexico'', which separates the North America and the Pelican nebulae, and we demonstrate that this area is a very active star forming region. A wide-field imaging study with interference filters has revealed 35 new Herbig-Haro objects in the Gulf of Mexico. A grism survey has identified 41 H$\\alpha$ emission-line stars, 30 of them new. A small cluster of partly embedded pre-main sequence stars is located around the known LkH$\\alpha$ 185-189 group of stars, which includes the recently erupting FUor HBC~722.} ", "introduction": " ", "conclusions": "" }, "1101/1101.3320_arXiv.txt": { "abstract": " ", "introduction": "A key feature of the primordial seed fluctuations observed in the cosmic microwave background (CMB) and the large-scale structure (LSS) is their {\\it scale-invariance}~\\cite{WMAP,SDSS}. Furthermore, the observed near-Gaussianity of the data requires the fluctuations to be {\\it weakly coupled}\\,\\footnote{`Weak coupling' implies that the action for the fluctuations has a well-defined expansion in which the leading terms of order $n+1$ in the fluctuations are smaller than the terms of order $n$. We shall be more precise about this in \\S\\ref{sec:strong}.} over a large range of scales. In this paper we show that if the universe was dominated by a single degree of freedom at the time when the fluctuations were created, the above two facts together strongly constrain the background spacetime at that time. \\vskip 4pt Our basic point is very simple: % the scale-invariance of two-point correlations does not guarantee scale-invariance of interactions. Observations require scale-invariance of the power spectrum of curvature perturbations over at least $\\Delta N \\sim 10$ $e$-folds, from CMB scales ($\\sim 10^4$\\,Mpc) to galactic scales ($\\sim 1$\\,Mpc). These fluctuations are created while the conformal time $\\tau$ evolves by a factor of $e^{-\\Delta N}$. In quasi-de Sitter backgrounds, this exponential change in time appears only in the time-evolution of the scale factor $a(\\tau) = -(H \\tau)^{-1}$, which is not a physical observable. All observable couplings during inflation are nearly time-independent. This is a consequence of the time-translation invariance of the background. The scale-invariance of the fluctuations produced by inflation therefore applies to all $n$-point functions. In contrast, in non-de Sitter backgrounds with a single fluctuating degree of freedom\\footnote{All of our constraints can be evaded in {\\it multi-field} non-de Sitter cosmologies such as~\\cite{Creminelli:2007aq, Buchbinder:2007ad, Lehners:2007wc}.} (e.g.~\\cite{PaulJustin, Justin2}), it is hard to achieve both scale-invariant two-point correlations and time-independent interactions simultaneously. This is the case because the $n$-point functions of adiabatic (or `single-clock') fluctuations around Friedmann-Robertson-Walker (FRW) backgrounds aren't independent, but are related by the symmetries of the background~\\cite{Creminelli:2006xe, Cheung}: time translations are spontaneously broken in FRW backgrounds and the symmetry is {\\it non-linearly} realized by a Goldstone boson. % This non-linear realization of time translations forces unavoidable relations between the quadratic terms and the higher-order terms in the Lagrangian. The interactions are particularly constrained when the Goldstone boson of time translations is the only relevant degree of freedom. It is precisely these constraints that will allow us to make statements about the interaction Lagrangian in any cosmology that produces a scale-invariant two-point function. We find that any time-dependence that is not in the scale factor typically leads to an exponential growth of interactions, so that one has to be worried that the fluctuations become strongly coupled at some point during the evolution. At this point perturbative control of the theory would be lost. \\vskip 4pt In Section~\\ref{sec:two} we review the conditions for scale-invariant two-point correlations in the effective theory of adiabatic fluctuations. At lowest order in derivatives, we identify two classes of backgrounds that allow for scale-invariant two-point functions for the primordial curvature perturbation $\\zeta$. These two possibilities are distinguished by whether $\\zeta$ is constant on superhorizon scales ({\\sf Case I}) or grows in time by an exponentially large amount ({\\sf Case II}). These two possibilities correspond to the background being a dynamical attractor or not. Within these two classes of solutions there are important subclasses depending on whether the background is expanding or contracting, has a strongly time-dependent equation of state or a strongly time-dependent speed of sound. In Section~\\ref{sec:strong} we explain that higher-order interactions of $\\zeta$ and the associated higher-point correlations severely limit the possibilities. We present the leading cubic interactions, identify the time-dependent `coupling constants', and discuss when they lead to a strong coupling problem. In Section~\\ref{sec:attractor} we study the attractor solutions in detail. For the case of a constant speed of sound we find {\\it exact} solutions which allow us to study FRW backgrounds in full generality. We show that only near-de Sitter backgrounds, with intrinsically small time-variations of all quantities, avoid the strong coupling problem. We explain why our conclusions are not affected significantly when the speed of sound is allowed to vary in time. We present our conclusions in Section~\\ref{sec:conclusions}. \\vskip 4pt Four appendices contain further technical details: In Appendix~\\ref{sec:A2} we derive the quadratic action for the curvature perturbation $\\zeta$ in the framework of the effective theory of adiabatic fluctuations~\\cite{Creminelli:2006xe, Cheung}. We show that, at leading order in a derivative expansion, the dynamics is characterized by two free functions: the Hubble rate $H(t)$ and the sound speed $c_s(t)$. In the limit in which $M_{\\rm pl}^2 \\dot H$ becomes smaller than other parameters in the Lagrangian, higher-derivative terms can be the leading effects. We discuss this limit in Appendix~\\ref{sec:new}. In Appendix~\\ref{sec:nonA} we turn our attention to the non-attractor cases. These theories are significantly less predictive and require a variety of additional assumptions about the evolution both before and after the scale-invariant modes are produced. We will show that even granting the most optimistic such assumptions the parameter space is significantly constrained by the weak-coupling requirement. Finally, in Appendix~\\ref{sec:bispectra} we compute the full bispectra for the non-attractor cases of theories with non-trivial speed of sound, following earlier work by Khoury and Piazza~\\cite{Justin}. ", "conclusions": "\\label{sec:conclusions} Cosmic microwave background (CMB) and large-scale structure (LSS) observations constrain the primordial seed fluctuations to be {\\it i}) nearly scale-invariant and {\\it ii}) approximately Gaussian (or weakly coupled). In this paper we asked what these two observational facts together teach us about the cosmological background at the time when these fluctuations exited the horizon. We considered the most general effective theory of adiabatic fluctuations around completely arbitrary FRW backgrounds allowing both for expansion and contraction. We assumed that the observed cosmological perturbations arose from quantum mechanical fluctuations of a single degree of freedom. For theories with constant sound speed and requiring the background to be a dynamical attractor, we reduced the requirement of scale-invariant two-point correlations to a non-linear differential equation for the scale factor $a(\\tau)$, \\beq q^2 = a^2 \\epsilon = a^2 \\left[ 2 - \\frac{a'' a}{(a')^2}\\right] \\propto \\frac{1}{\\tau^2}\\ . \\eeq This equation has an exact solution---a special case of solutions to the generalized Emden-Fowler equation~\\cite{RussianBook}. Using this solution we reproduced two known limits: slow-roll inflation~\\cite{TASI} ($a \\propto \\tau^{-1}$) and ekpyrotic contraction (or expansion) with rapidly changing equation of state~\\cite{PaulJustin,Justin2} ($\\epsilon \\propto \\tau^{-2}$) . \\begin{figure}[h!] \\centering \\includegraphics[width=.8\\textwidth]{ziScanMinus6.pdf} \\caption{\\sl The value of the $\\eta$ parameter 10 $e$-folds after the beginning of the scaling phase illustrates the strong coupling problem arising in backgrounds that deviate too much from quasi-de Sitter ($z_i = - 1 - \\epsilon_i$).} \\label{fig:ziScan2} \\end{figure} The non-linear realization of time diffeomorphisms in the effective theory of the fluctuations forces specific relationships between the coefficients of the quadratic Lagrangian and the interaction Lagrangian. This drastically limits the number of viable models that can produce a scale-invariant two-point function while staying weakly coupled for a sufficiently long time. In fact, we showed that only inflation leads to nearly time-independent couplings of all higher-order interactions such as $\\epsilon$, $\\eta$ and $c_s^{-2}$. This implies that only inflationary spacetimes allow the fluctuations to stay weakly coupled over the range of scale relevant to cosmological observations. For all non-de Sitter backgrounds we identified a strong coupling problem that limits the predictivity of these solutions. Within 10 $e$-folds after the beginning of the scaling phase the perturbative expansion of the theory breaks down. This makes it challenging to produce the modes observed in the CMB and the LSS in a consistent non-de Sitter background. Dropping the requirement that the background should be an attractor, we found non-de Sitter solutions with rapidly time-varying speed of sound (following previous work by Khoury and Piazza~\\cite{Justin}). In Appendix~\\ref{sec:nonA} we identified the regime of parameter space in which those theories evade both quantum mechanical and classical strong coupling problems. Since these backgrounds are not attractors, the curvature perturbations $\\zeta$ evolve after horizon exit and predictions depend on assumptions about the physics both before and after the regime during which the scale-invariant modes exit the horizon. We described the basic model-building requirements for non-attractor non-de Sitter backgrounds with weakly-coupled scale-invariant fluctuations. Whether these physical elements can be realized in a coherent theoretical framework remains an open question. In contrast, it is quite remarkable how the time-translation invariance of quasi-de Sitter backgrounds without any additional physical ingredients solves the horizon and flatness problems, while at the same time allowing for scale-invariant $n$-point functions. \\vskip 8pt {\\it Note added.} While this paper was being completed, Ref.~\\cite{Justin2} appeared which has some overlap with our Section~\\ref{sec:attractor}. \\subsubsection*{Acknowledgements} D.B.~thanks Justin Khoury, Federico Piazza and Paul Steinhardt for helpful discussions and correspondence. The research of D.B.~is supported by the National Science Foundation under PHY-0855425 and a William D.~Loughlin Fellowship at the Institute for Advanced Study. L.S.~is supported in part by the National Science Foundation under PHY-0503584. M.Z.~is supported by the National Science Foundation under PHY-0855425 and AST-0907969, as well as by the David and Lucile Packard Foundation and the John D.~and Catherine T.~MacArthur Foundation. \\newpage \\appendix" }, "1101/1101.0924_arXiv.txt": { "abstract": "The magnetic field in many astrophysical plasmas -- such as the Solar corona and Earth's magnetosphere -- has been shown to have a highly complex, three-dimensional structure. Recent advances in theory and computational simulations have shown that reconnection in these fields also has a three-dimensional nature, in contrast to the widely used two-dimensional (or 2.5-dimensional) models. Here we discuss the underlying theory of three-dimensional magnetic reconnection. We also review a selection of new models that illustrate the current state of the art, as well as highlighting the complexity of energy release processes mediated by reconnection in complicated three-dimensional magnetic fields. ", "introduction": "Magnetic reconnection is a fundamental process that is ubiquitous in astrophysical plasmas. It facilitates the release of energy stored in the magnetic field by permitting a change in the magnetic topology in an almost ideal plasma. As such, reconnection is universally accepted to be a key ingredient in the behaviour of many astrophysical plasmas, including the interiors and atmospheres of stars such as the Sun, planetary magnetospheres, accretion disks, and pulsar magnetospheres. Much of the literature on reconnection focusses on the two-dimensional problem, due to the theoretical and computational simplifications that this allows. However, it is now becoming clear that magnetic reconnection in an even weakly three-dimensional (3D) setting is crucially different from the planar 2D case. In this article, we review recent advances in three-dimensional reconnection theory, and the complex picture that is emerging of the possible regimes of reconnection in 3D. This review is by necessity limited and misses a number of important facets of reconnection research. Complementary reviews include those by \\cite{priest2000,biskamp2000,zweibel2009,yamada2010}. In Section \\ref{topsec} we introduce some key measures and features of magnetic field structure in 3D that are crucial to understanding where and how reconnection operates in 3D, while in Section \\ref{recpropsec} we discuss some fundamental differences between 2D and 3D reconnection. In Sections \\ref{nonnullsec}-\\ref{sepsec} we review the current picture of the various different 3D reconnection regimes, and in Section \\ref{obssec} we touch briefly on recent results from large-scale numerical simulations and observations. We finish with a summary in Section \\ref{sumsec}. ", "conclusions": "\\label{sumsec} Magnetic reconnection is a universal process in astrophysical plasmas. It facilitates the release of stored magnetic energy by permitting changes of magnetic topology (or just changes of field line connectivity for reconnection in the absence of topological structures) and as such is a key ingredient of many energetic processes in these environments. It is only in recent years that the rich geometrical and topological structure of these astrophysical plasmas has begun to be appreciated, following great advances in observations. In 3D, the qualitative properties of reconnection are much richer than in 2D, where reconnection occurs only at magnetic X-type null points. There are many topological and geometrical features of a magnetic field that may be favourable sites for current growth. 3D reconnection processes in different magnetic field structures have different characteristic properties, and as such can be classified into separate regimes. Reconnection may occur in the absence of any null points, with the presence of counter-rotational flows on either side of the diffusion region being a characteristic feature. Reconnection may also occur at isolated null points, with `torsional spine', `torsional fan' and `spine-fan' reconnection modes thus far identified. The most common of these -- spine-fan reconnection -- involves a local collapse of the null to form a current sheet (focussed at the null) that locally spans both the spine and fan, with flux transfer through spine line and fan (separatrix) surface. Furthermore, reconnection may occur at separator lines connecting pairs of nulls, in which the properties of the reconnection process may share some properties with null or non-null reconnection modes. While each of the reconnection modes has its own characteristics, they all share some fundamental properties that make them distinct to 2D ($\\EE\\cdot\\BB=0$) reconnection. In particular, the non-existence of a unique flux-conserving velocity anywhere within the diffusion region (i.e.~for any field lines threading the diffusion region), which implies that reconnection occurs throughout the diffusion region, not at a single point as in 2D. As a result, there is no one-to-one cut-and-paste rejoining of field lines. While we are now beginning to understand some of the properties of 3D reconnection, there is much left to discover. Future advances are likely to be led by large-scale MHD simulations being developed alongside fundamental theory. % Many important open questions remain, such as: \\begin{itemize} \\item What are the quantitative properties of the different 3D reconnection regimes? Thus far the vast majority of our knowledge is of a qualitative nature, and quantitative studies are required to probe, for example, the diffusion region dimensions. What is the range of possible reconnection rates in each regime, and what determines the reconnection rate? How do these numbers scale with different plasma parameters? While 3D reconnection studies have so far (by necessity) used the MHD approximation, it will be important in the future to investigate the implications of including additional physics such as Hall and electron pressure tensor electric fields in Ohm's law, to better model the physics within the diffusion region. \\item What is the relationship between null, non-null, and separator reconnection? \\item Which of these modes of reconnection is most important in more realistic (complex) magnetic fields -- when the local models discussed above are embedded within global magnetic field configurations? One step towards answering this question would be to determine the observational signatures of the different reconnection regimes. \\item What are the most common mechanisms of current sheet formation in complex 3D magnetic fields? How does the reconnection influence the global evolution of the magnetic field? (This may be more complex than previously appreciated -- recent MHD simulations have shown that the distribution of reconnection sites may be highly complex and flux may be reconnected multiple times in generic 3D MHD evolutions \\cite{parnell2008,pontin2010}) \\item How is the magnetic energy transferred to other forms, and ultimately dissipated, in 3D reconnection? This includes the important question of the pattern and efficiency of particle acceleration in each 3D reconnection regime. \\end{itemize}" }, "1101/1101.2921_arXiv.txt": { "abstract": "We study two large classes of alternative theories, modifying the action through algebraic, quadratic curvature invariants coupled to scalar fields. We find one class that admits solutions that solve the vacuum Einstein equations and another that does not. In the latter, we find a deformation to the Schwarzschild metric that solves the modified field equations in the small coupling approximation. We calculate the event horizon shift, the innermost stable circular orbit shift, and corrections to gravitational waves, mapping them to the parametrized post-Einsteinian framework. ", "introduction": "Although black holes (BHs) are one of the most striking predictions of General Relativity (GR), they remain one of its least tested concepts. Electromagnetic observations have allowed us to infer their existence, but direct evidence of their non-linear gravitational structure remains elusive. In the next decade, data from very long-baseline interferometry~\\cite{2009ApJ...695...59D,Fish:2009ak} and gravitational wave (GW) detectors~\\cite{Eardleyprl,Will:2004xi,Berti:2004bd,2005ApJ...618L.115J,Alexander:2007kv,Arun:2009pq,Sopuerta:2009iy,Schutz:2009tz,Yunes:2009bv,Yunes:2009ke,2009PhLB..679..401A,2010GWN.....4....3S,Mishra:2010tp,Yunes:2010yf,Stein:2010pn,2010PhRvD..82l2003T,DelPozzo:2011pg,Molina:2010fb} should allow us to image and study BHs in detail. Such observations will test GR in the dynamical, non-linear or strong-field regime, precisely where tests are currently lacking. Testing strong-field gravity features of GR is of utmost importance to physics and astrophysics as a whole. This is because the particular form of BH solutions, such as the Schwarzschild and Kerr metrics, enter many calculations, including accretion disk structure, gravitational lensing, cosmology and GW theory. The discovery that these metric solutions do not accurately represent real BHs could indicate a strong-field departure from GR with deep implications to fundamental theory. Such tests require parametrizing deviations from Schwarzschild or Kerr. One such parameterization at the level of the metric is that of {\\emph{bumpy BHs}}~\\cite{Collins:2004ex,Vigeland:2009pr,Vigeland:2010xe}, while another at the level of the GW observable is the {\\emph{parameterized post-Einsteinian}} (ppE) framework~\\cite{2009PhRvD..80l2003Y,2010PhRvD..82h2002Y}. In both cases, such parameterizations are greatly benefited from knowledge of specific non-GR solutions, but few, 4D, analytic ones are known that represent regular BHs (except perhaps in dynamical Chern-Simons (CS) gravity~\\cite{2009PhRvD..79h4043Y,Alexander:2009tp} and Einstein-Dilaton-Gauss-Bonnet (EDGB) gravity~\\cite{Kanti:1995vq,Torii:1996yi,Kanti:1997br,Pomazanov:2003wq,Pani:2009wy}). Most non-GR BH solutions are known through numerical studies. In this approach, one chooses a particular alternative theory, constructs the modified field equations and then postulates a metric ansatz with arbitrary functions. One then derives differential equations for such arbitrary functions that are then solved and studied numerically. Such an approach was used, for example, to study BHs in EDGB gravity~\\cite{Kanti:1995vq,Torii:1996yi,Kanti:1997br,Pomazanov:2003wq,Pani:2009wy}. Another approach is to find non-GR BH solutions analytically through approximation methods. In this scheme, one follows the same route as in the numerical approach, except that the differential equations for the arbitrary functions are solved analytically through the aid of approximation methods, for example by expanding in (a dimensionless function of) the coupling constants of the theory. Such a {\\emph{small-coupling approximation}}~\\cite{Campanelli:1994sj,Cooney:2008wk,2009PhRvD..79h4043Y} treats the alternative theory as an {\\emph{effective and approximate}} model that allows for small GR deformations. This approach has been used to find an analytic, slowly-rotating BH solution in dynamical CS modified gravity~\\cite{2009PhRvD..79h4043Y,Alexander:2009tp}. But not all BH solutions outside of GR must necessarily be different from standard GR ones. In fact, there exists many modified gravity theories where the Kerr metric remains a solution. This was the topic studied in~\\cite{2008PhRvL.100i1101P}, where it was explicitly shown that the Kerr metric is also a solution of certain $f(R)$ theories, non-dynamical quadratic gravity theories, and certain vector-tensor gravity theories. Based on these fairly generic examples, it was then inferred that the astrophysical observational verification of the Kerr metric could not distinguish between GR and alternative theories of gravity. Such an inference, however, is not valid, as it was later explicitly shown in~\\cite{2009PhRvD..79h4043Y}. Indeed, there are alternative gravity theories, such as dynamical CS modified gravity, where the Kerr metric is not a solution. This prompted us to study what class of modified gravity theories admit Kerr and which do not. We begin by considering the most general quadratic gravity theory with dynamical couplings, as this is strongly motivated by low-energy effective string actions~\\cite{1992PhLB..285..199C,Boulware:1985wk,Green:1987mn,Green:1987sp,lrr-2004-5}. When the couplings are static, we recover the results of~\\cite{2008PhRvL.100i1101P}, while when they are dynamic we find that the Kerr metric is not a solution. In the latter case, we find how the Schwarzschild metric must be modified to satisfy the corrected field equations. We explicitly compute the shift in the location of the event horizon and innermost stable circular orbit. Such modifications to the BH nature of the spacetime induce corrections to the waveforms generated by binary inspirals. We compute such modifications and show that they are of so-called second post-Newtonian (PN) order, i.e.~they correct the GR result at ${\\cal{O}}(v^{4})$ relative to the leading-order Newtonian term, where $v$ is the orbital velocity. We further show that one can map such corrections to the parameterized post-Einsteinian (ppE) framework~\\cite{Yunes:2009ke}, which proposes a model-independent, waveform family that interpolates between GR and non-GR waveform predictions. This result supports the suggestion that the ppE scheme can handle a large class of modified gravity models. The remainder of this paper is organized as follows. Sec~\\ref{sec:quad-grav} defines the set of theories we will investigate and computes the modified field equations. Sec.~\\ref{sec:non-spin-sol} solves for BH solutions in this class of theories. Sec.~\\ref{sec:prop-of-sol} discusses properties of the solution and Sec.~\\ref{sec:impact} studies the impact that such BH modifications will have on the GW observable. Sec.~\\ref{sec:future-work} concludes by pointing to future possible research directions. For the remainder of this paper, we use the following conventions: latin letters in index lists stand for spacetime indices; parentheses and brackets in index lists stand for symmetrization and antisymmetrization respectively, i.e.~$A_{(ab)} = (A_{ab} + A_{ba})/2$ and $A_{[ab]} = (A_{ab} - A_{ba})/2$; we use geometric units with $G = c = 1$. ", "conclusions": "" }, "1101/1101.2260_arXiv.txt": { "abstract": "We report on oxygen abundances determined from medium-resolution near-IR spectroscopy for a sample of 57 carbon-enhanced metal-poor (\\cemp{}) stars selected from the Hamburg/ESO survey. The majority of our program stars exhibit oxygen-to-iron ratios in the range $+0.5<$ \\ofe{}$ <+2.0$. The \\ofe{} values for this sample are statistically compared to available high-resolution estimates for known CEMP stars, as well as to high-resolution estimates for a set of carbon-normal metal-poor stars. Carbon, nitrogen, and oxygen abundance patterns for a sub-sample of these stars are compared to yield predictions for very metal-poor asymptotic giant-branch abundances in the recent literature. We find that the majority of our sample exhibit patterns that are consistent with previously studied CEMP stars having s-process-element enhancements, and thus have very likely been polluted by carbon- and oxygen-enhanced material transferred from a metal-poor asymptotic giant-branch companion. ", "introduction": "Carbon-enhanced metal-poor (\\cemp{}) stars are quite common in the halo populations of the Milky Way, and are of particular interest, as they preserve important astrophysical information concerning the early chemical evolution of the Galaxy (Beers \\& Christlieb 2005). Previous work has indicated that at least 20\\% of stars with metallicities [Fe/H] $< -2.0$ exhibit large over-abundances of carbon ([C/Fe] > +1.0; Lucatello et al. 2006; Marsteller et al. 2009), although recent studies (e.g. Cohen et al. 2005; Frebel et al. 2006), have claimed that this fraction is somewhat lower (9\\% and 14\\%, respectively, for [Fe/H] $< -2.0$). In any case, the fraction of CEMP stars rises to 30\\% for [Fe/H] $< -3.0$, 40\\% for [Fe/H] $< -3.5$, and 100\\% for [Fe/H] $< -4.0$ (Beers \\& Christlieb 2005; Frebel et al. 2005; Norris et al. 2007). There exist a number of classes of \\cemp{} stars, some of which have been associated with proposed progenitor objects. CEMP-s stars (those with s-process-element enhancement), for example, are the most commonly observed type to date. High-resolution spectroscopic studies have revealed that around 80\\% of \\cemp{} stars exhibit s-process-element enhancement (Aoki et al. 2007). The favored mechanism invoked to account for these stars is mass transfer of carbon-enhanced material from the envelope of an asymptotic giant-branch (AGB) star to its binary companion; it is this surviving binary companion that is now observed as a CEMP-s star. The class of CEMP-no stars (which exhibit no strong neutron-capture-element enhancements) are particularly prevalent among the most metal-poor stars. Possible progenitors for this class include massive, rapidly-rotating, mega metal-poor ([Fe/H] $< -6.0$) stars, which models suggest have greatly enhanced abundances of CNO due to distinctive internal burning and mixing episodes, followed by strong mass loss (Meynet et al. 2006; Hirschi et al. 2006; Meynet et al. 2010). Another suggested mechansim is pollution of the interstellar medium by so-called faint supernovae associated with the first generations of stars, which experience extensive mixing and fallback during their explosions (Umeda \\& Nomoto 2003, 2005; Tominaga et al. 2007); high [C/Fe] and [O/Fe] ratios are predicted in the ejected material. This model well reproduces the observed abundance pattern of the CEMP-no star BD+44:493, the ninth-magnitude [Fe/H] $= -3.7$ star (with [C/Fe]$ = +1.3$, [N/Fe] $= +0.3$, [O/Fe] $= +1.6$) discussed by Ito et al. (2009). The great majority of known CEMP stars were originally identified as metal-poor candidates from objective-prism surveys, such as the HK survey (Beers et al. 1985, 1992), and the Hamburg/ESO Survey (HES; Christlieb 2003; Christlieb et al. 2008), based on a weak (or absent) \\ion{Ca}{2} K line. Some candidate CEMP stars also come from a list of HES stars selected from the prism plates based on their strong molecular lines of carbon (Christlieb et al. 2001). Medium-resolution spectra for most of these objects have been obtained over the past few years (Goswami et al. 2006; Marsteller 2007; Goswami et al. 2010, Sivarani et al., in preparation). Inspection of these data indicate that at least 50\\% of these targets are consistent with identification as CEMP stars, while the others are roughly solar-metallicity carbon-rich stars. Dedicated surveys for CEMP stars covering a wide range of carbon abundance and metallicities are just now getting underway, based on the observed strength of the CH G band measured from the HES prism plates (e.g., Placco et al. 2010). In order to more fully test the association of CEMP-no stars with massive primordial stars and/or faint supernovae, and to better explore the nature of the s-process in low-metallicity AGB stars (which is still rather poorly understood; Herwig 2005), we require measurements of the important elements C, N, and O for as large a sample of CEMP stars as possible. While estimates of carbon and nitrogen abundances can be determined from medium-resolution optical or near-UV spectra of CEMP stars (e.g., Rossi et al. 2005; Beers et al. 2007b; Johnson et al. 2007; Marsteller et al. 2009), high-resolution spectroscopy is usually required in order to obtain estimates of oxygen abundances from the forbidden [\\ion{O}{1}] $\\lambda$6300 \\AA{} line, the $\\lambda$7700 \\AA{} triplet (e.g., Schuler et al. 2006; Sivarani et al. 2006; Fabbian et al. 2009, and references therein), or the OH lines at 1.5-1.7 $\\micron$ (Mel\u00b4endez \\& Barbuy 2002). Masseron et al. (2010) provides a useful compilation of known elemental abundances for CEMP stars. In addition to abundance measurements for metal-poor halo stars, oxygen abundances have also been measured directly in the gas phase in damped Lyman $\\alpha$ systems (Pettini et al. 2002, 2008). If a star has a measured carbon abundance (and, assuming C/O $> 1$, which applies for most CEMP stars), essentially all of the O is locked up in CO molecules, and medium-resolution spectroscopy of the \\co{} ro-vibrational bands in the near-IR can be used for estimation of \\ofe{} (e.g., Beers et al. 2007b, and references therein). Although one sacrifices measurement accuracy, relative to high-resolution studies, this approach has the great advantage that medium-resolution spectroscopy can be gathered far faster than high-resolution spectroscopy, ensuring that much larger samples of stars can be investigated. In addition, the large separation of the $^{13}$CO lines from the $^{12}$CO lines at 2.3$\\mu$m provides a straightforward means to measure the important mixing diagnostic $^{12}$C/$^{13}$C, as long as the S/N of the spectra are sufficient. This paper is outlined as follows. In Section 2 we discuss details of the observations and data reduction procedures used in the present study. Section 3 describes the previously determined atmospheric parameter estimates and their origins, as well as details about the synthetic spectra. Methods used for determination of \\ofe{} for our sample of stars are described in Section 4. Our results, and a statistical comparison to high-resolution estimates of [C/Fe] and [O/Fe] for a subset of our program stars can be found in Section 5. Section 6 is a short discussion of our results; conclusions follow in Section 7. ", "conclusions": "We have used near-IR medium-resolution spectroscopy in order to estimate \\ofe{} for a sample of candidate carbon-enhanced stars selected from the Hamburg/ESO Survey. This method of abundance analysis allows us to obtain oxygen abundances accurate to about 0.4 dex. The use of four separate CO features to estimate oxygen abundances from the near-IR spectra allows for more precise estimates, based on a robust average of the independently determined fits. A large spread of derived \\ofe{} values are obtained for this sample, ranging from near the solar value to as much as one hundred times greater. A comparison of our abundance determinations with high-resolution estimates was carried out. The values of \\ofe{} for our full set of 57 CEMP stars largely fall within regions of parameter space occupied by the high-resolution estimates of oxygen for other CEMP stars. We also found that the majority of our stars have oxygen abundances that are consistent with known CEMP-s and CEMP-r/s stars. Only a few stars could be considered CEMP-no stars, based on the data compiled in Masseron et al. (2010). This is likely due to the fact that CEMP-no stars commonly have lower metallicities than most of the stars in this sample. Oxygen enhancements (on the order of \\ofe{} $= +0.7$) have also been observed in very metal-poor stars without significant carbon enhancement, indicating that there were early oxygen-producing nucleosynthetic sites in the Galaxy independent of any enhancement by AGB evolution. However, we find that the \\cfe{}, \\ofe{}, and \\nfe{} (when available) estimates follow the patterns from Herwig (2004) closely enough that mass-transfer from a AGB companion is a likely scenario for many of the stars in our sample, especially when the effects of dilution are considered. Our measured carbon abundances always exceed the available high-resolution \\nfe{} abundances. If the origin of CNO abundance patterns comes from hot-bottom-burning (HBB) in an intermediate mass (AGB) star, one would expect to see elevated \\nfe{} relative to \\cfe{} and \\ofe{}. This signature is not found in our sample, but it has been suggested that other mechanisms, such as cool-bottom-processing (Wasserburg et al. 1995; Denissenkov \\& VandenBerg 2003) or the occurrence of HIF, can alter the levels of nitrogen enhancement. It is likely that the majority of CEMP stars in this sample will turn out to be enhanced in neutron-capture elements. Consistency of most of our program stars with the CEMP-s class, based both on comparison to AGB models and existing high-resolution data, is expected since that CEMP stars with s-process-element enhancement are the most commonly observed type to date. However, recent chemical evolution models (Cescutti \\& Chiappini 2010) have revealed that the winds from massive, rapidly-rotating metal-poor stars can result in a large scatter in the predicted abundances of C, N, and O, presumably without the production of neutron-capture elements. Therefore, we are currently unable to assign classification to this sample of CEMP stars. High-resolution spectra of the stars in our sample will help clarify these questions." }, "1101/1101.2867_arXiv.txt": { "abstract": "We use a sample of 8298 galaxies observed as part of the HST $H_{160}$-band GOODS NICMOS Survey (GNS) to construct the galaxy stellar mass function both as a function of redshift and stellar mass up to $z = 3.5$. Our mass functions are constructed within the redshift range $z=1-3.5$ and consist of galaxies with stellar masses of $M_{*}=10^{12} M_{\\odot}$ down to nearly dwarf galaxy masses of $M_{*}=10^{8.5} M_{\\odot}$ in the lowest redshift bin. We discover that a significant fraction of all massive $M_{*}>10^{11} M_{\\odot}$ galaxies are in place up to the highest redshifts we probe, with a decreasing fraction of lower mass galaxies present at all redshifts. This is an example of `galaxy mass downsizing', and is the result of massive galaxies forming before lower mass ones, and not just simply ending their star formation earlier as in traditional downsizing scenarios, whose effect is seen at $z<1.5$. By fitting Schechter functions to our mass functions we find that the faint end slope ranges from $\\alpha=-1.36$ to $-1.73$, which is significantly steeper than what is found in previous investigations of the mass function at high redshift. We demonstrate that this steeper mass function better matches the stellar mass added due to star formation, thereby alleviating some of the mismatch between these two measures of the evolution of galaxy mass. We furthermore examine the stellar mass function divided into blue/red systems, as well as for star forming and non-star forming galaxies. We find a similar mass downsizing present for both blue/red and star-forming/non-star forming galaxies, and further find that red galaxies dominate at the high mass end of the mass function, but that the low mass galaxies are mostly all blue, and therefore blue galaxies are creating the steep mass functions observed at $z > 2$. We furthermore show that, although there is a downsizing such that high mass galaxies are nearer their $z=0$ values at high redshift, this turns over at masses $M_{*}\\sim10^{10} M_{\\odot}$, such that the lowest mass galaxies are more common than galaxies at slight higher masses, creating a `dip' in the observed galaxy mass function. We argue that the galaxy assembly process may be driven by different mechanisms at low and high masses, and that the efficiency of the galaxy formation process is lowest at masses $M_{*}\\sim10^{10} M_{\\odot}$ at $1 < z < 3$. Finally, we calculate the integrated stellar mass density for the total, blue and red populations. We find the integrated stellar mass density of the total and blue galaxy population is consistent with being constant over $z=1-2$, while the red population shows an increase in integrated stellar mass density over the same redshift range. ", "introduction": "A deep understanding of the high redshift universe is vital in order to complete our knowledge of galaxy formation, and hence uncover the history of the universe as a whole. With the continued development of instrumentation and of new analysis techniques our knowledge of the high redshift universe is increasing rapidly. Thanks to new deep imaging and multiobject spectroscopy we now have the power to routinely look back at the universe over cosmic time to witness step by step evolution. The result of this has been a wealth of observations of large samples of galaxies over a large redshift range in various surveys, such as within GOODS (\\citealt{Giav04}), COSMOS and z-COSMOS (\\citealt{Scov07} and \\citealt{Lill07}) and AEGIS (\\citealt{Davi07}). These large samples give us the power to achieve statistically meaningful results concerning the evolution of galaxy properties. As a result a detailed picture is beginning to form regarding when galaxy stellar mass is built up over cosmic time. Various studies (\\citealt{Dick03}, \\citealt{Dror05}, \\citealt{Cons07}, \\citealt{Elsn08}, \\citealt{Pere08}) have focused on the change of stellar mass density with time, and have seen generally consistent results. These studies find that the integrated stellar mass density decreases at higher redshift, as expected since the ongoing process of star formation increases the amount of stellar mass in the universe over time. These studies also show that roughly 50\\% of the mass density of the universe is in place by $z \\sim 1$. This implies prior to this redshift the comoving stellar mass density has a more rapid evolution. This ties in with studies that show that the star formation rate peak is at $z\\ge1-2$ (\\citealt{Mada96}, \\citealt{Hopk04}, \\citealt{Hopk06}). While studies of the star formation history has been the traditional way to examine and probe galaxy evolution, a great deal of research has been carried out investigating the evolution of galaxies using the galaxy stellar mass function. Early examples of the measurement of the local galaxy mass function at $z \\sim 0$ have been carried out by e.g. \\citet{Cole01} and \\citet{Bell03}, using surveys such as 2MASS, 2dF, and the Sloan Digital Sky Survey. data. These and other investigations construct the stellar mass function, and integrated stellar mass density within the local universe, and hence provide a vital benchmark for comparison with higher redshifts. It is important to understand the stellar mass function of the local universe in itself as this traces the integrated star formation and mass assembly history over the entire universe. However, by extending similar studies to higher redshift we can investigate not only galaxy growth as a function of stellar mass, but also the growth of integrated galaxy stellar mass with time. That is, we can trace the evolution of stellar mass for galaxies of different stellar masses, i.e., very massive galaxies vs. lower mass galaxies, over time. This gives us insights into the growth of galaxies of different stellar masses due to star formation, mergers and other assembly processes. Many studies have investigated the stellar mass function up to $z=2$ using relatively large area surveys (\\citealt{Font04}, \\citealt{Bund06}, \\citealt{Borc06}, \\citealt{Fran06} \\citealt{Bell07}, \\citealt{Bolz09}). This has been further extended to even higher redshifts using deeper surveys, generally within a much smaller area (\\citealt{Cons05}, \\citealt{Font06}, \\citealt{Pere08}, \\citealt{Kaji09}). These studies of the stellar mass function at high redshifts allows us to form a picture of the high mass end of the stellar mass function as a function of redshift. It is widely agreed that the dominance of star formation within massive galaxies ends much earlier than within low stellar mass galaxies by $z \\sim 1$ (\\citealt{Bund06}). This is one form of galaxy downsizing, whereby the higher stellar mass galaxies have their star formation truncated, or gas depleted, earlier than lower mass galaxies. This formation scenario was first observed by \\citet{Cow96} and subsequently observed in various studies including e.g. \\citet{Baue05}, \\citet{Feul05}, \\citet{Bund06}, and \\citet{Verg08}. Downsizing is now accepted as part of the formation scenario of galaxies, but is not yet fully understood. What is also not fully understood is whether the related, but different process of mass downsizing, is occurring in tandem, such that the high mass galaxies form their stellar mass earlier than lower mass galaxies, and when this stellar mass differentiated galaxy formation process first reveals itself. Despite the considerable work done investigating the high mass end of the stellar mass function, there are still many issues that are not yet fully understood. Firstly, the generally shape of the high redshift stellar mass function is not well described. Nearby stellar mass functions are well fit by a form of the Schechter function, and there are various investigations regarding how the parameters of this fit change with redshift. Secondly, difficulties with obtaining deep data at high redshifts mean that the low mass end of the stellar mass function has not been explored as fully, or to as high a redshift. More recent work has started to uncover a possible steepening with redshift of the low stellar mass end, and a ``dip\" in the intermediate stellar mass range (e.g. \\citealt{Kaji09}). It has been suggested that this is a result of evolution of different galaxy populations driven by their mass (\\citealt{Bolz09}, \\citealt{Dror09}, \\citealt{Ilbe09}, \\citealt{Pozz09}). The exact nature and reasons behind such features of the stellar mass function are not well understood, and this can only be improved upon with deeper, more robust data. This will then lead to a better understanding of how the populations of galaxies, as defined by stellar mass, change and evolve over time. In this paper we use data from the GOODS NICMOS survey (GNS) to investigate how the stellar mass function evolves from $z= 1$ to $z \\sim 3.5$. By examining the stellar mass functions of galaxies ranging in stellar mass from $M_{*}=10^{12} M_{\\odot}$ to as low as $M_{*}=10^{8.5} M_{\\odot}$ we investigate when, and which galaxies are forming at various epochs in the universe. The depth of the GNS data is such that we are able to probe over a factor of 10$^{3}$ the stellar mass evolution up to $z \\sim 3.5$ and importantly trace how galaxies of different masses are evolving with time. We find throughout this paper a differential in the stellar mass function and how it evolves, revealing a strong stellar mass dependence in the galaxy formation process. We describe this and give some general explanations for how this differential evolution can occur due to different physical processes. The paper is set out as follows: Section \\ref{sec:GNS} discusses the GOODS-NICMOS Survey, the galaxy sample and how the data used in this paper was obtained. Section \\ref{sec:totalMF} examines the galaxy stellar mass functions of all the galaxies in various redshift bins. In Sections \\ref{sec:MFforBR} and \\ref{sec:MFforSF} we split the galaxies into blue and red and star forming and non-star forming respectively. Section \\ref{sec:massden} describes the calculation of the stellar mass densities for the total sample and for the red and blue populations. Sections \\ref{sec:diss} and \\ref{sec:summ} contain the discussion and summary of our findings respectively. Throughout this paper we assume $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$ and $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$. AB magnitudes and a Salpeter IMF are used throughout. ", "conclusions": "\\label{sec:diss} \\subsection{Comparison With Other Surveys and Data} There is a large amount of research exploring the stellar mass function over a large range in redshifts, from $z=0-6$. What is unique about the GNS, and this work in particular, is the depth and redshift range provided by the data. A large number of surveys (e.g. COSMOS, zCOSMOS and COMBO-17) previously examine the stellar mass function at a lower redshift range than the GNS. For $z<1$ these surveys probe a large stellar mass range e.g. \\citet{Pozz09}, with mass limits of $M_{*}\\sim 10^{8.5} M_{\\odot}$ at $z=0$ to $10^{10} M_{\\odot}$ at $z=0.75$ and \\citet{Dror09}, with mass limits of $M_{*}\\sim 10^{8.5} M_{\\odot}$ at $z=0.2$ to $10^{9} M_{\\odot}$ at $z=1$. For $z=1-1.5$ our stellar mass functions are complete down to $M_{*}\\sim 10^{8.5} M_{\\odot}$, and for $z=3-3.5$ we find we are complete to $M_{*}\\sim 10^{9.5} M_{\\odot}$. Other work, which probes a similar redshift range, does not probe the faint end of the mass function as low as we do here. For example \\citet{Pere08} has a mass limit of $M_{*}\\sim 10^{9.5} M_{\\odot}$ at $z=1-1.3$ and the mass limits of \\citet{Kaji09} range from $M_{*}\\sim 10^{9} M_{\\odot}$ to $10^{10} M_{\\odot}$ over $z=1-3.5$. We are thus able to examine stellar mass evolution at $z>1$ as a function of mass for the first time. \\subsection{The High Mass Galaxies and Downsizing} In all of the stellar mass functions presented in this work, we find that the high mass galaxies have formed earlier than the low mass galaxies. For the total galaxy stellar mass functions, as shown in Figure \\ref{totMF}, galaxies whose stellar mass is $M_{*}>10^{11.5} M_{\\odot}$ are close to the local value as early as $z=3-3.5$ (for $M_{*}\\sim 10^{11.5} M_{\\odot}$ the number densities are $31^{+13}_{-14}$ \\% of the local value at this redshift). These galaxies have reached the local number density by $z=1-1.5$ ($94^{+6}_{-39}$\\% the local value). Downsizing conventionally says that the high mass galaxies have stopped star forming before the low mass galaxies. However what we see here is galaxy ``mass downsizing\", where by the mass of the high mass galaxy is formed before that in the low mass galaxies. This is occurring at $z>3$ unlike star formation downsizing which starts at $z<1.5$ (Bauer et al. 2010). This means the majority of the stellar mass of a galaxy is already in place before the star formation largely stops. This suggests that it is not the star formation making the galaxy massive, instead it is some other process such as mergers, which we do see happening often at high redshift (\\citealt{Cons03}, \\citealt{Cons06} and \\citealt{Cons08}). This mass downsizing effect is also seen for the halos of galaxies (\\citealt {Fouc10}). Figure \\ref{totresid} also shows the relative ratio between the local galaxy stellar mass function of \\citet{Cole01} at $z=0$ and the number densities calculated in this work. The green horizonal line is how the ratio would appear if all the number densities were at the local value. The black line is a linear fit to the points from the lowest to highest ratio, hence its steepness correlates to how dominant downsizing is in that redshift range. We find that although at $z=2-2.5$ the gradient of the line is very steep, there is a generally decreasing trend with lower redshift. We find similar results for the red and blue mass functions in Figure \\ref{BRMF}, with the red population of $M_{*}\\sim10^{11.5} M_{\\odot}$ galaxies being very close to their local number density in the range $z=2.5-3$ ($15^{+18}_{-3}$\\% for galaxies in this redshift range). The evolution of the high mass blue galaxies towards the local value is slower than that of the red. We find that the blue galaxies with stellar mass $M_{*}>10^{11.5} M_{\\odot}$ are mostly present by $z=1-1.5$ ($56^{+27}_{-26}$\\% for $M_{*}>10^{11.5} M_{\\odot}$ for galaxies with in this redshift range). This is because these blue galaxies are likely to evolve into red galaxies at some stage. We also find downsizing in the star forming and non-star forming mass function, where the star forming/ non-star forming populations behave very similarly to the blue/red populations. \\subsection{The Intermediate Mass Dip} Figure \\ref{totresid} shows various lines indicating the ratio between the local stellar mass function and the local stellar mass function with slightly varying parameters (the coloured and dashed/dotted lines). For each line only one parameter is changed at one time. We see that between $z= 1.0$ to $1.5$ galaxies in the mass range $M_{*}\\sim10^{10.5} M_{\\odot} - 10^{9.5} M_{\\odot}$, have generally higher ratios, hence they have generally lower number densities relative to galaxies with $M_{*}>10^{11} M_{\\odot}$ and $M_{*}<10^{9.5} M_{\\odot}$ . This dip feature also seems to be present for $z>2$, but is much weaker and shifted to the stellar mass range $M_{*}\\sim10^{11.5} M_{\\odot} - 10^{10.5} M_{\\odot}$. The various coloured lines show the features of the mass function is not just an effect of the changing of the shape of the Schechter function. The dip in the total stellar mass function has been seen before and can be explained by the differential evolution of the blue and red galaxy populations. This is consistent with studies such as \\citet{Pozz09} and \\citet{Ilbe09} who find a dip in the total stellar mass function that can be fit by the combination of these two population. We also find this to be the case here, as shown in Figure \\ref{totcomp}. We find that a double Schechter function (the black dash dotted line), computed using the parameters of the blue and red fits, matches the form of the total Schechter function for all galaxies. The only large discrepancy occurs at the high mass end in the range $2 2 \\times 10^{20}$ cm$^{-2}$) and sub-DLAs ($1 \\times 10^{19} \\lesssim N$(\\ion{H}{1}) $\\leq 2 \\times 10^{20}$] provide sensitive probes of the chemical enrichment history of the Universe from $z = 0$ to $z > 4$ (e.g., Prochaska et al. 2003; Wolfe et al. 2005, Meiring et al. 2009,2011), but the context of the DLAs (i.e., the environment and nature of the absorbers) is generally hard to study, and currently only limited information is available regarding the origins of DLAs (e.g., Chen \\& Lanzetta 2003; Rao et al. 2003; Battisti et al. 2012). Considering the cross section of a disk or halo of gas, it is likely that many of the DLAs arise in the outer regions of galaxies. The Milky Way is a damped Ly$\\alpha$ absorber. Since its absorption context can be scrutinized in great detail, the Milky Way provides a valuable laboratory for understanding the nature of DLAs/sub-DLAs, but measurements of abundances patterns in the more distant {\\it outer} Galaxy are still relatively limited. Third, while some abundances in the outer galaxy and abundance gradients have been measured using \\ion{H}{2} region emission lines, there has long been concern about whether \\ion{H}{2} region abundances are biased by ``self pollution'', i.e. whether the abundances are boosted by freshly formed metals from massive stars embedded within the \\ion{H}{2} region. Abundances from absorption lines in distant background objects probe random locations with respect to foreground \\ion{H}{2} regions and can test the self-pollution hypothesis, and some absorption-line studies have indeed suggested that self pollution does occur (Cannon et al. 2005 and references therein). The discovery that many gas-rich galaxies have ``extended ultraviolet disks'' (Thilker et al. 2007) underscores the importance of investigating this issue -- \\ion{H}{2} regions may be biased in favor of the UV-bright star clusters that comprise the extended UV disks and might not accurately represent abundance gradients for testing theoretical work as discussed above. Fourth, outer galaxy abundances can be used to investigate whether some gaseous structures of the Milky Way could be due to interactions with satellite galaxies. This is related to the questions raised above: one means to bring gas into galaxies is to strip the interstellar media of dwarf satellite galaxies as they plunge into, and merge with, the central galaxy. Satellite interactions can also stimulate the growth of galactic structures such as warps (e.g., Weinberg \\& Blitz 2006; Quillen et al. 2009) and may drive continuing (low-level) star formation in elliptical galaxies (Kaviraj et al. 2011). Since satellite galaxies can have significantly different abundances compared to each other and the Galactic disk, the outer galaxy abundances provide clues about the origins of galactic structures and the importance of this mechanism for bringing additional gas into galaxies. For these reasons we have conducted a study of the abundances in the outer galaxy using absorption lines recorded in high-resolution spectra of two QSOs observed with the {\\it Hubble Space Telescope (HST)} and the {\\it Far Ultraviolet Spectroscopic Explorer (FUSE)}. We selected two QSOs for this study, HS0624+6907 and H1821+643. These sight lines are unique among the QSOs and AGNs that have been in observed in the ultraviolet at high spectral resolution and with good signal-to-noise (S/N) ratios because the QSOs are at relatively low Galactic latitudes and thus provide an opportunity to study outer-galaxy gas near the plane. Moreover, these QSOs lie behind a high-velocity gaseous structure with a large angular extent in the outer galaxy known as the ``Outer Arm'' (OA). This structure is also unique because its distance has recently been constrained, which is valuable for understanding its nature and implications. In addition, there are several high-velocity cloud complexes near the OA with similar kinematics, including Complexes C, G, and H. The {\\it HST} archive includes high-quality spectra of many stars in the directions of these gas clouds, and these stellar spectra provide an opportunity to obtain new constraints on the distance of these objects. In this paper, we present a study of these outer-Galaxy gaseous structures. We focus on the abundances and physical conditions in the Outer Arm and whether the OA is related to the HVCs in its proximity. In \\S~\\ref{outerarmnotes}, we provide some comments on the Outer Arm and the QSO sight lines that probe this part of the Galaxy. We present the QSO and star observations and absorption-line measurements in \\S~\\ref{obssec} and \\S~\\ref{absmeas}, respectively, and we discuss new constraints on the distance of the Outer Arm, and the nearby high-velocity cloud Complex G, in \\S~\\ref{distance_section}. In \\S~\\ref{ionabun}, we examine the physical conditions of the OA, and we use models to evaluate the impact of ionization corrections on the metallicity measurements. We also make some remarks on the nature of the highly ionized gas in the Outer Arm. We discuss our results in \\S~\\ref{disc} with an emphasis on the possible origin and implications of the Outer Arm. ", "conclusions": "\\label{disc} As we have summarized above, the extended gas cloud known as the Outer Arm is usually considered to be part of the warp in the outer Galaxy and possibly the most distant spiral arm. However, the recent observations of Lehner \\& Howk (2010) have revealed aspects of the OA that are not expected in this scenario -- the OA has a high-velocity component that is inconsistent with Galactic rotation and instead indicates that the OA kinematics are similar to those of Complex C, which is close to the OA in velocity and on the sky. We have presented additional observations that can be used to further probe the nature of the OA. Briefly, we find: \\begin{enumerate} \\item Based on ultraviolet absorption lines, we have measured OA abundances in the directions of two QSOs, H1821+643 and HS0624+6907. The OA oxygen abundance in the direction of H1821+643 is [O/H] = $-0.30^{+0.12}_{-0.27}$. The metallicity derived from the HS0624+6907 sight line suggests that the OA could have a range of metallicities in different locations with $Z_{\\rm OA} = 0.2 - 0.5 Z_{\\odot}$, but the HS0624+6907 metallicity is more uncertain and is consistent with the H1821+643 metallicity when uncertainties are taken into account. The OA metallicity is only marginally higher than the abundances usually measured in Complex C, $Z_{\\rm Comp. C} = 0.1 - 0.3 Z_{\\odot}$ (e.g., Tripp et al. 2003; Collins et al. 2007; Shull et al. 2011). \\item Both the OA and Complex C are underabundant in nitrogen. This is often interpreted to be an indication that the gas is ``chemically young'' since nitrogen is synthesized in intermediate-mass stars, and thus more time is required to build up the nitrogen abundance than is required for species such as oxygen that are rapidly produced in Type II supernovae (e.g., Vila Costas \\& Edmunds 1993; Pettini et al. 1995). \\item High-resolution spectroscopy of several stars in the direction of the OA indicates that the object is at a Galactocentric radius of $9 - 18$ kpc. Based on currently available constraints, it is possible that the OA and Complex C are at similar distances. In addition, we have detected the HVC Complex G, which is close to the Outer Arm, in absorption toward two stars. This places Complex G relatively close to the solar Galactocentric radius at $R_{G}$ = $8 - 10$ kpc. \\item The OA absorption profiles toward H1821+643 show that the OA is a complex, multiphase entity with several narrow components, including narrow features in the profiles of highly ionized species. This suggests that the OA is, at least in part, interacting with the ambient gas of the Milky Way. \\end{enumerate} This ensemble of information suggests that the Outer Arm might have a more complicated origin than the usual attribution to the outer warp. This concept has been proposed before: Davies (1972) suggested a connection between the OA and Complex C, which he proposed to be generated by the Large Magellanic Cloud. More recently, Kawata et al. (2003) attempted to model the production of both of these structures by the interaction of a satellite galaxy with the Milky Way. While they found that structures similar to the OA + Complex C could be generated, the absence of the interacting satellite (which they require to have a mass comparable to the LMC) in the expected part of the sky poses a problem for this model. At approximately the same time, the Monoceros Ring structure was discovered (Newberg et al. 2002), and given the similarity of the kinematics and location of the Monoceros Ring to the Outer Arm, it is interesting to ask if these high-velocity gas clouds could be related to the merging satellite that produced the stellar Monoceros Ring. We consider this hypothesis in a future paper." }, "1101/1101.4664_arXiv.txt": { "abstract": "\\noindent The requirement that SUSY should solve the hierarchy problem without undue fine-tuning imposes severe constraints on the new supersymmetric states. With the MSSM spectrum and soft SUSY breaking originating from universal scalar and gaugino masses at the Grand Unification scale, we show that the low-fine-tuned regions fall into two classes that will require complementary collider and dark matter searches to explore in the near future. The first class has relatively light gluinos or squarks which should be found by the LHC in its first run. We identify the multijet plus $E_{T}^{miss}$ signal as the optimal channel and determine the discovery potential in the first run. The second class has heavier gluinos and squarks but the LSP has a significant Higgsino component and should be seen by the next generation of direct dark matter detection experiments. The combined information from the 7~TeV LHC run and the next generation of direct detection experiments can test almost all of the CMSSM parameter space consistent with dark matter and EW constraints, corresponding to a fine-tuning not worse than 1:100. To cover the complete low-fine-tuned region by SUSY searches at the LHC will require running at the full 14~TeV CM energy; in addition it may be tested indirectly by Higgs searches covering the mass range below 120~GeV. ", "introduction": "Weak scale supersymmetry (SUSY) has been proposed as a solution to the hierarchy problem, {\\it i.e.}\\ it ensures that the electroweak breaking scale is consistent with radiative corrections without undue fine-tuning. However, to achieve this, the new SUSY states must be relatively light. To quantify how light the SUSY states should be and and how much stress experimental limits already put on SUSY, one can apply a measure of fine-tuning. In~\\cite{Cassel} an analysis using SOFTSUSY~\\cite{Allanach:2001kg} was made of the status of SUSY searches in the constrained minimal supersymmetric standard model (CMSSM) using the electroweak (EW) fine-tuning measure, $\\Delta$, introduced in~\\cite{Ellis:1986yg} computed to two-loop order. A scan of the CMSSM parameter space was performed, requiring acceptable radiative electroweak breaking, non-tachyonic SUSY particle masses (avoiding colour and charge breaking vacua), consistency with the experimental bounds on superpartner masses and with constraints from BR$(b\\rightarrow s\\gamma)$, BR$(B_{s}\\rightarrow\\mu^{+}\\mu^{-})$ and the muon $(g-2)$ as detailed in~\\cite{Cassel}. In Figure \\ref{higgsp} the envelope of the shaded region shows the EW fine-tuning \\newrevision{$\\Delta$}. One may see that, imposing all the constraints listed above except for the LEPII bound on the Higgs mass, there is a minimum of \\newrevision{the EW} fine-tuning, $\\Delta\\approx 9$, for a Higgs mass of $m_h\\approx 114$~GeV.\\footnote{Note that the calculation of $m_h$ is subject to a theoretical uncertainty of about 2--3~GeV~\\cite{Degrassi:2002fi}, see also \\cite{Cassel}.} \\strike{Although we consider here only the low-fine-tuned regions of the CMSSM, this covers a large part of the low-fine-tuned MSSM parameter space because many models beyond the CMSSM are more fine-tuned. For example gauge mediation is more fine-tuned because the former relies on the scalar focus point~\\cite{Feng:2000bp} whereas the latter does not have the focus point as it does not have degenerate scalars at a high scale.} \\revision{Although the analysis presented here is concerned with the CMSSM, this class of model permits the presence of a scalar focus point, which favours small fine tuning of the electroweak scale~\\cite{Feng:2000bp} relative to more generic MSSM models. % } Thus if one excludes the low-fine-tuned regions of the CMSSM one can say that much larger ranges of the MSSM parameter space are also disfavoured. % To date the only class of non-CMSSM models with MSSM spectrum that have lower fine-tuning are those with non-universal gaugino masses with a `natural' relation between them that reduces the gluino mass relative to the CMSSM case \\revision{ \\cite{Gogoladze:2009bd,Horton:2009ed} }. In addition to EW fine-tuning, another important constraint is that the thermal relic abundance of the lightest supersymmetric particle (LSP), which contributes to dark matter, should be consistent with cosmological observations, \\revision{under the assumption of R-parity conservation}. Here we use MicrOMEGAs~\\cite{Belanger:2006is} as computational tool. Imposing the constraints of \\cite{Dunkley:2008ie} limits the allowed region to be in the coloured regions of Figure~\\ref{higgsp}. In regions~1 (red) and 5 (black points superposed on the green region) $h^{0}$ and $A^{0}/H^{0}$ resonant annihilation respectively are responsible for reducing the dark matter abundance within current bounds. Region~2 (purple) has significant bino-higgsino mixing in the LSP, and annihilation proceeds via higgsino t-channel exchange to EW gauge bosons. Finally, for regions~3 (green) and 4 (blue), the dominant processes are stau and stop co-annihilation, respectively. Overall, requiring that the SUSY dark matter relic density should be within present bounds raises the minimal fine-tuning to $\\Delta=15$, still quite reasonable, corresponding to $m_{h}\\approx 115$~GeV. % \\begin{figure}[t] \\centering \\includegraphics[width=9cm]{EWFThiggs_noCDMS.eps} \\def\\baselinestretch{1.} \\caption{{\\protect\\small Two-loop fine-tuning versus Higgs mass for the scan over CMSSM parameters with no constraint on the Higgs mass. The solid line is the minimum fine-tuning with $\\left( \\protect\\alpha_s^{}, M_t^{} \\right) =$ (0.1176, 173.1~GeV). The dark green, purple, crimson and black coloured regions have a dark matter density within $\\Omega h^{2} = 0.1099 \\pm 3 \\times 0.0062$~\\cite{Dunkley:2008ie} (i.e. 3$\\sigma$ saturation) while the lighter coloured versions of these regions lie below this bound. The colours and their associated numbers refer to different LSP structures as described in the text. Regions 1, 3, 4 and 5 have an LSP which is mostly bino-like. In region~2, the LSP has a significant higgsino component. \\newrevision{Representative phase space points for regions 1,2...5, denoted SUG1, SUG2,...,SUG5, respectively, will be analysed in detail later on, together with the point of minimal EW $\\Delta\\approx 9$ denoted SUG0.} }} \\label{higgsp} \\end{figure} An immediate question is what is the best way to test the low-fine-tuned region of SUSY parameter space and will it be tested soon? In this the search for the Higgs boson plays a very important role because, {\\it c.f.} Figure~\\ref{higgsp}, for large Higgs masses, quantum corrections make the fine-tuning exponentially sensitive to the Higgs mass. Thus, if the Higgs is not found below $\\approx$120~GeV, it will imply that the fine-tuning is uncomfortably % large, $\\Delta>100$. However, a LHC Higgs discovery at such low mass will be difficult and is likely to take several years. Given this, it is of interest to consider to what extent direct SUSY searches will probe the regions of low fine-tuning. In this paper we consider both collider searches and dark matter searches. For this it is important to analyse the nature of the SUSY spectrum in the region of low fine-tuning ($\\Delta<100$) in light of the Higgs mass and dark matter constraints. Of course, if the LHC is to detect new SUSY states in its first run with of ${\\cal{O}}(1)$fb$^{-1}$ luminosity at a CM energy of 7~TeV, some states with coupling to the gluon, {\\it i.e.}\\ squarks or gluinos, must be light enough to have a sizeable cross section. In \\cite{Baer:2010tk}, the discovery reach of the LHC % with 1~(2)~fb$^{-1}$ luminosity at a CM energy of 7~TeV was determined as $m_{\\tg}\\sim 1100~(1200)$~GeV for $m_{\\tq}\\sim m_{\\tg}$, and $m_{\\tg}\\sim 620~(700)$~GeV for $m_{\\tq}\\gg m_{\\tg}$. The first results from CMS \\cite{cms} for 35~pb$^{-1}$ of data exclude gluino masses below 500~GeV for $m_0\\lesssim 350$~GeV, but no limit is obtained for $m_0\\gtrsim 500$~GeV. \\begin{figure}[t] \\centering\\def\\baselinestretch{1.} \\begin{tabular}{cc} \\includegraphics[width=0.48\\textwidth]{EWFTm0_noCDMS_Higgs114.eps}& \\includegraphics[width=0.48\\textwidth]{EWFTgluino_noCDMS_Higgs114.eps}\\\\ (a) & (b) \\end{tabular} \\caption{ {\\protect\\small (a) Fine-tuning versus scalar mass parameter, (b) fine-tuning versus gluino mass; in both cases the constraint on the Higgs mass, $m_{h}>114.4$~GeV is applied.} } \\label{higgsp1} \\end{figure} How does this compare to the expectation for the regions of low fine-tuning? From Figure~\\ref{higgsp1}(a) one sees that these regions % (which we take as $\\Delta<100$) can either have light squarks (region~3, green points) or be close to the scalar ``focus point'' corresponding to heavy squarks.\\footnote{The few points of region 4 that satisfiy the Higgs mass limit have very large fine-tuning, $\\Delta\\sim 10^3$, and are not shown in Figures~\\ref{higgsp1}--\\ref{dmcdmsbound}. \\revision{Relaxing the Higgs mass constraint by 2-3~GeV has a negligible effect on the plots, except for region 3 where reduced fine tuning ($\\Delta \\gtrsim 45$) becomes possible}} Figure~\\ref{higgsp1}(b) shows that the low fine-tuned points with heavy squarks (regions 1, 2 and 5) have two components. The first, region 1, has a small gaugino mass parameter and corresponds to gluinos with mass of about 400--500~GeV, potentially accessible to LHC discovery in the first run. The remaining regions (2 and 5) have gluino masses beyond the LHC reach in the 7~TeV run. However these regions may be accessible to dark matter searches. Such searches put limits on, {\\it e.g.}, the spin-independent scattering cross section for neutralino dark matter and this in turn is dominated by Higgs-boson exchange coupling to the higgsino and bino components in the neutralino. Typically the LSP has a large bino component but in restricted regions of parameter space it may also have a sizeable higgsino component. In Figure~\\ref{higgsp2} we plot the higgsino component versus the fine-tuning measure. One can see that only the purple points, region~2, have a significant higgsino component. \\begin{figure}[t] \\centering\\def\\baselinestretch{1.} \\begin{tabular}{cc} \\includegraphics[width=0.48\\textwidth]{EWFTh1ino_noCDMS_Higgs114.eps}& \\includegraphics[width=0.48\\textwidth]{EWFTh2ino_noCDMS_Higgs114.eps}\\\\ (a) LSP $\\tilde{h}_{1}^{0}$ component & (b) LSP $\\tilde{h}_{2}^{0}$ component \\end{tabular} \\caption{{\\protect\\small The higgsino components in the LSP, again requiring $m_{h}>114.4$~GeV.}} \\label{higgsp2} \\end{figure} Once one knows the composition of the LSP it is straightforward to compute the spin independent cross section relevant to the direct dark matter searches. This is shown in Figure~\\ref{dmcdmsbound}, plotted against the LSP mass. Here, the cross-section has been rescaled by $R= \\Omega h^{2}/0.1099$, to take into account of the dark matter abundance at each of the points. Also shown is the current best bound coming from the CDMS experiment~\\cite{Ahmed:2009zw}. For more details see \\cite{cassel-thesis}. One sees that this already provides a significant test of region 2. A factor of 10 improvement in the dark matter sensitivity, which should be achieved by SuperCDMS in 2013, will probe almost the full range of region 2, the exception being points that do not saturate the dark matter density. For these latter points a two orders of magnitude improvement will be needed. The points in region 1, 3 and 5 have very small higgsino component and dark matter searches do not test this region. However, as may be seen from Figure~\\ref{higgsp}, the bulk of this data is very fine-tuned with $\\Delta>100$. The situation for $\\Delta<100$ is depicted in Figure~\\ref{dmcdmsboundft100}. We see that the parameter space still needing to be explored has shrunk considerably: only a small part of regions 1 and 3 and an even smaller part of region 5 remains to be tested. These results are also illustrated in Figure~\\ref{m0m12} in the $(m_0,m_{1/2})$ plane, which shows the same points of low fine-tuning that pass the constraints mentioned. Notice that all these points are also consistent with the latest CMS observed exclusion area~\\cite{cms} (situated below the \\strike{red} \\revision{black} curve), \\revision{and all but a fraction of the region 3 points are consistent with the latest ATLAS exclusion limit \\cite{daCosta:2011qk} }. We will discuss below the LHC configuration needed to scan the SUSY spectrum for these residual regions. We conclude this Introduction with a side-remark: one often states that the very small area of points left in the moduli space $(m_0, m_{1/2})$, that respect all experimental constraints, renders supersymmetry an unlikely solution to the hierarchy problem. However, even if this area is reduced to few points due to further experimental constraints, recall that many of them have acceptable fine-tuning (in our case $\\Delta<100$). That is, the density and size of the area of points allowed in Figure~\\ref{m0m12} does not necessarily have a physical relevance, and cannot be used to conclude that only few points left would immediately invalidate supersymmetry as a solution to the hierarchy problem. It would rather indicate the most likely values of these moduli $(m_0,m_{1/2})$, that a fundamental theory beyond MSSM should fix dynamically, to avoid degenerate vacua. \\begin{figure}[t]\\centering \\def\\baselinestretch{1.} \\begin{tabular}{cc} \\includegraphics[width=0.48\\textwidth]{SIscaled_Higgs114.eps} & \\includegraphics[width=0.48\\textwidth]{SIscaled_Higgs114_tan50.eps}\\\\ (a) $\\tan \\beta \\leq 45$ & (b) $50 \\leq \\tan \\beta \\leq 55$ \\end{tabular} \\caption{ {\\protect\\small Scaled spin independent cross section for LSP-proton scattering, with $m_{h}>114.4$~GeV. The scaling factor $R=\\Omega h^{2}/0.1099$ has been applied. The solid line is the CDMS-II limit. All points satisfy $\\Omega h^{2} < 0.1285$, with those with darker shading lying within $3\\sigma$ of the WMAP bound, $\\Omega h^{2} = 0.1099 \\pm 3 \\times 0.0062$.} \\revision{The mSUGRA phase space scan was discrete only in the $\\tan \\beta$ dimension with 30 slices in the range $2 \\leq \\tan \\beta \\leq 45$, and two further slices at $\\tan \\beta = 50, 55$.} } \\label{dmcdmsbound} \\end{figure} \\begin{figure}[t]\\centering \\def\\baselinestretch{1.} \\begin{tabular}{cc} \\includegraphics[width=0.48\\textwidth]{SIscaled_Higgs114_Delta100.eps} & \\includegraphics[width=0.48\\textwidth]{SIscaled_Higgs114_tan50_Delta100.eps}\\\\ (a) $\\tan \\beta \\leq 45$ & (b) $50 \\leq \\tan \\beta \\leq 55$ \\end{tabular} \\def\\baselinestretch{1.} \\caption{ {\\protect\\small Same as Figure~\\ref{dmcdmsbound} but imposing in addition $\\Delta<100$. } } \\label{dmcdmsboundft100} \\end{figure} \\begin{figure}[t] \\centering\\def\\baselinestretch{1.} \\begin{tabular}{cc} \\includegraphics[width=0.48\\textwidth]{m0m12_CDMS_Higgs111_Delta100_ATLASline.eps} & \\includegraphics[width=0.48\\textwidth]{m0m12_CDMS_Higgs114_Delta100_ATLASline.eps}\\\\ (a) $m_h^{} > 111$\\,GeV & (b) $m_h^{} > 114.4$\\,GeV \\end{tabular} \\caption{{\\protect\\small Regions of low fine-tuning ($\\Delta <100$) in the $m_0$ versus $m_{1/2}$ plane, summed over $\\tan\\beta$ and $A_0$. All points satisfy the SUSY and Higgs mass limits, $\\Omega h^{2} < 0.1285$ (dark points having $0.0913<\\Omega h^{2} < 0.1285$), the B-physics and $\\delta a_\\mu$ constraints, and the CDMS-II bound on the dark matter detection cross section. The area below the \\strike{red} \\revision{black} line shows the CMSSM exclusion (for $\\tan\\beta=3$ and $A_0=0$) from the CMS dijet$+E_T^{miss}$ analysis \\cite{cms}, \\revision{and that below the red line the ATLAS exclusion area \\cite{daCosta:2011qk}}. \\newrevision{ The different lower bounds on $m_h$ of these plots are applied to show the impact of the 2-3 GeV theoretical uncertainty in the calculation of $m_h$ \\cite{Cassel,Degrassi:2002fi}.}}} \\label{m0m12} \\end{figure} ", "conclusions": "In summary, using the fine-tuning measure, we have made a detailed study of the possibility of testing in the near future the CMSSM as a solution to the hierarchy problem. Broadly, the regions of low fine-tuning split into two characteristic classes. The first class has light gluinos or light squarks and will likely be tested in the 7~TeV run at the LHC. The second class has a heavy gluino but the LSP has a significant higgsino component; this class is testable by direct dark matter searches in the near future. Together, these complementary experiments will be able to cover almost all of the parameter space with fine-tuning $\\Delta<100$. To cover all of this parameter space by SUSY searches at the LHC will require running at the full 14~TeV CM energy. In addition, the low-fine-tuned regions can be tested indirectly by Higgs searches covering the mass range $m_h\\le 120$~GeV." }, "1101/1101.4387_arXiv.txt": { "abstract": "In this paper we discuss the excess gamma rays from the Galactic center, the WMAP haze and the CoGeNT and DAMA results in WIMPless models. At the same time we also investigate the low energy constraints from the anomalous magnetic moment of leptons and from some lepton flavor violating decays. It is found that, for scalar or vector WIMPless dark matter, neither the WMAP haze nor the CoGeNT and DAMA observations could be explained simultaneously with the excess gamma rays from the Galactic center. As to fermion WIMPless dark matter, it is only marginally possible to accommodate the CoGeNT and DAMA results with the excess gamma rays from the Galactic center with vector connector fields. On the other hand, only scalar connector fields could interpret the WMAP haze concerning the constraints of anomalous magnetic moment of leptons. Furthermore, if there is only one connector field for all the charged leptons, some lepton flavor violating decays could happen with too large branching ratios severely violating the experimental bounds . ", "introduction": "There are lots of gravitational evidences suggesting the dominance of dark matter (DM) over baryonic visible matter in our universe, though little is known about its non-gravitational interaction. If DM can annihilate or decay into the standard model (SM) particles, gamma rays will be generated either directly or by the final state radiation, inverse Compton scattering and bremsstrahlung of the final charged particles. As these gamma rays depend on the DM density as $\\rho^{2}$ for annihilations and $\\rho$ for decays, they might be detectable at locations with very high density of dark matter, such as the Galactic center (GC). Recently Hooper and Goodenough \\cite{Hooper:2010mq} investigated the gamma rays from the GC using the first two years of data from the Fermi Gamma Ray Space Telescope. They modeled the backgrounds and found that the morphology and spectrum of the gamma rays between $1.25^\\circ$ and $10^\\circ$ from the GC can be well understood. However excess emission of gamma rays seems to be present within $1.25^\\circ$ degree of the GC. This may well indicate that we have to improve our understanding on the astrophysical background. Nevertheless, this additional component could also be interpreted by the DM annihilation into tau leptons, with a mass around $8$ GeV and the annihilation cross section $\\langle \\sigma v \\rangle \\sim 10^{-26}\\mbox{cm}^3/$s. There are also anomalous microwave emission from the inner Galaxy in the WMAP data \\cite{Finkbeiner:2003im,Dobler:2007wv} which is known as the \"WMAP haze\". It was further observed by Hooper and Linden \\cite{Hooper:2010im} that, if the DM annihilates equally into $e^+ e^-$, $\\mu^+ \\mu^-$ and $\\tau^+ \\tau^-$, the WMAP haze can be simultaneously accounted for by the synchrotron radiation emitted from the annihilation final state of electrons and positrons. Notice that WMAP haze could also be explained by astrophysical processes \\cite{Crocker:2010qn}. Coincidentally, the annual modulation effect measured by the DAMA/LIBRA collaboration \\cite{Bernabei:2010mq} together with the recent results about excesses of events (over the backgrounds) reported by CDMS-II \\cite{Ahmed:2009zw} and CoGeNT \\cite{Aalseth:2010vx}, if interpreted as elastically scattering DM, favor the DM mass about 5 - 10 GeV \\cite{Bottino:2009km,Kopp:2009qt,Fitzpatrick:2010em,Chang:2010yk,Foot:2010rj}. The null results from Xenon100 \\cite{Aprile:2010um} and CDMS-II \\cite{Ahmed:2010wy} are in tension with the previous data. Nevertheless, concerning the uncertainties of the DM velocity distribution \\cite{Lisanti:2010qx} and of the relative scintillation efficiency of the liquid xenon at low recoil energy \\cite{Savage:2010tg,Collar:2010nx,Collar:2010ht}, it is still possible to interpret the CoGeNT and DAMA/LIBRA observations as light DM with mass around $7$ GeV \\cite{Hooper:2010uy}, and at the same time, consistent with the null results of Xenon100 and CDMS-II. Such a light DM is beyond the conventional paradigm of weakly-interacting massive particle (WIMP) with the mass around $100$ GeV to several TeV, though it is certainly possible to realize this scenario in SUSY models, e.g. in an effective minimal supersymmetric extension of the Standard Model \\cite{Bottino:2003iu,Bottino:2003cz,Bottino:2009km,Fornengo:2010mk,Scopel:2011qt}. Interestingly, in the WIMPless models \\cite{Feng:2008ya} the DM particle naturally has the correct thermal relic density with a wide range of mass varying from TeV down to subGeV \\cite{Feng:2008mu}. The key observation in \\cite{Feng:2008ya} is that, the thermal relic density of a stable particle is proportional to $m^2/g^4$, with $m$ and $g$ the typical mass and coupling entering the annihilation cross section. The correct thermal relic density can be easily realized for the DM particle X in the hidden sector if \\begin{equation} \\frac{m_X}{g_X^2} \\sim \\frac{m_{weak}}{g_{weak}^2}~, \\end{equation} with the weak scale $m_{weak}\\simeq 100$ GeV and $g_{weak}\\simeq 0.65$. Such relation can be satisfied for example in the gauge-mediated supersymmetry breaking models if there is only one supersymmetry(SUSY) breaking sector, so that the SUSY-breaking soft scales in both the visible sector and the hidden sector are generated by the gauge interactions with the same SUSY-breaking sector \\footnote{There exists many light DM scenarios, of which dark atoms of DM is an exotic example \\cite{Khlopov:2010pq}}. In the following, we will then study the possible implications of the above experimental results for the WIMPless models. Notice that we shall not attempt to explore the full parameter space of the WIMPless models. Instead, as illustration only benchmark models are examined in this paper. For example, following Ref. \\cite{Hooper:2010im}, we only consider the WIMPless DM annihilating democratically into $e^+ e^-$, $\\mu^+ \\mu^-$ and $\\tau^+ \\tau^-$ to interpret the WMAP haze. But it is definitely possible to explain the WMAP haze with other choices of parameters provided the DM can significantly annihilate into $e^+ e^-$. Nevertheless, our main conclusions should remain qualitatively unchanged for general WIMPless models. Phenomenologically, the WIMPless DM has been used to explain the DAMA signal \\cite{Feng:2008dz}. The indirect detection of fermion WIMPless DM at the neutrino telescopes IceCube and DeepCore has been discussed in \\cite{Barger:2010ng}. The low energy constraints on WIMPless DM can be found in \\cite{McKeen:2009rm,McKeen:2009ny}. It is also worth noting that the particle physics implications for the gamma ray excess from the GC together with CoGeNT and DAMA/LIBRA has been studied in \\cite{Buckley:2010ve}, though in a model independent way and without the inclusion of the WMAP haze. This paper is organized as follows. In the next section, we adopt scalar and vector WIMPless models to discuss the excess gamma rays from the GC, the WMAP haze and the CoGeNT and DAMA results. We also investigate the low energy constraints from the anomalous magnetic moment of leptons and from some lepton flavor violating decays. We then extend the discussions to fermion WIMPless models in the third section. Finally we conclude with a summary in section IV. ", "conclusions": "In this paper we discuss in WIMPless models the possibility to interpret the excess gamma rays from the Galactic center, the WMAP haze and the CoGeNT and DAMA results. At the same time some low energy constraints must be satisfied, such as the anomalous magnetic moment of leptons and lepton flavor violating $\\ell' \\to \\ell \\gamma$ decays. Notice that a connector sector Y is introduced to couple the DM candidate X to the SM particles via Yukawa-like interactions. As shown in \\cite{Hooper:2010mq}, the excess gamma rays from the GC implies that the DM particles should annihilate dominantly into lepton pairs with ${\\bar b} b$ or ${\\bar c}c$ final states less than $20$\\% of the time. For scalar or vector WIMPless DM, this limits the coupling strength of $XYb$ to be too small to account for the CoGeNT and DAMA observations. To interpret the WMAP haze, the $XY\\ell$ coupling is required to be roughly universal for $e$, $\\mu$ and $\\tau$ leptons, which will lead to too large contribution to anomalous magnetic moment of electron and muon. As to fermion WIMPless DM with scalar or vector connector sector, $\\lambda/m_Y \\sim $ a few $\\times 10^{-3}~\\mbox{GeV}^{-1}$ could accommodate the excess gamma rays from the GC and the CoGeNT and DAMA results. This corresponds to a scalar connector particle $Y_H$ less than $300$ GeV or a vector $Y_H$ less than about $400$ GeV. As $Y_H$ couples to quarks, it behaves like an exotic fourth generation quark, which has been restricted by the Tevatron to be heavier than $330$ GeV \\cite{Alwall:2010jc}. Therefore the scalar Y case seems to be disfavored, while the vector $Y^\\nu$ case is only marginally consistent with the Tevatron limit. Instead, to interpret the WMAP haze in the same framework, the constraints of anomalous magnetic moment of leptons can only be satisfied for the scalar connector fields. Furthermore, if there is only one connector field for charged leptons, the lepton flavor violating $\\ell \\to \\ell' \\gamma$ decays could happen in both cases with too large branching ratios severely violating the experimental bounds . Therefore it is difficult, if not impossible, for the scalar or vector WIMPless DM to interpret the excess gamma rays from the GC together with either the WMAP haze or the direct detection experiments CoGeNT and DAMA. While for fermion WIMPless DM, it may be possible to accommodate the excess gamma rays from the GC and the CoGENT and DAMA results with vector connector fields, though it is just marginally consistent with the Tevatron limit. On the contrary, only scalar connector fields could explain the WMAP haze under the constraints of anomalous magnetic moment of the leptons. In addition, if one connector field can couple to any charged leptons, the lepton flavor violating $\\ell \\to \\ell' \\gamma$ decays could happen with the branching ratios surpassing the current experimental bounds by (at most) even several orders of magnitude. Quantitatively, the above conclusions are only valid in the WIMPless models. But going beyond this specific model, one may at least learn that the excess gamma rays from the Galactic center, if interpreted as DM annihilation, could impose strong constraint on the DM couplings to quarks. Then one should check, under this constraint, the interpretation of the CoGENT and DAMA results in terms of elastic DM-nucleon scattering. It is also interesting to investigate the interplay between the direct and indirect DM searches and low energy precision observables, such as anomalous magnetic moment of leptons and the lepton flavor violating decays." }, "1101/1101.3758_arXiv.txt": { "abstract": "We describe the CRASH (Center for Radiative Shock Hydrodynamics) code, a block adaptive mesh code for multi-material radiation hydrodynamics. The implementation solves the radiation diffusion model with the gray or multigroup method and uses a flux limited diffusion approximation to recover the free-streaming limit. The electrons and ions are allowed to have different temperatures and we include a flux limited electron heat conduction. The radiation hydrodynamic equations are solved in the Eulerian frame by means of a conservative finite volume discretization in either one, two, or three-dimensional slab geometry or in two-dimensional cylindrical symmetry. An operator split method is used to solve these equations in three substeps: (1) solve the hydrodynamic equations with shock-capturing schemes, (2) a linear advection of the radiation in frequency-logarithm space, and (3) an implicit solve of the stiff radiation diffusion, heat conduction, and energy exchange. We present a suite of verification test problems to demonstrate the accuracy and performance of the algorithms. The CRASH code is an extension of the Block-Adaptive Tree Solarwind Roe Upwind Scheme (BATS-R-US) code with this new radiation transfer and heat conduction library and equation-of-state and multigroup opacity solvers. Both CRASH and BATS-R-US are part of the publicly available Space Weather Modeling Framework (SWMF). ", "introduction": "As photons travel through matter, the radiation field experiences changes due to net total emission, absorption, and scattering, see for instance \\citet{mihalas1984,pomraming2005,drake2006}. At high enough energy density the radiation will heat and accelerate the plasma. This coupled system is radiation hydrodynamics. The radiation can at a fundamental level be described by the time evolution of the spectral radiation intensity $I_\\nu({\\bf r},t,{\\bf n},\\nu)$, which is the radiation energy per unit area, per unit solid angle in the direction of photon propagation ${\\bf n}$, per unit interval of photon frequency $\\nu$, and per unit time interval. Several method have been developed to solve the radiation field in various degree of physics fidelity. In Monte Carlo radiative transfer methods, the radiation is statistically evaluated. Small photon packets are created with their energy and propagation direction statistically selected. The packets are propagated through matter using the radiation transfer equation \\citep{nayakshin2009,maselli2009, baek2009}. Characteristic methods use integration along rays of various lengths to solve for the angular structure of the radiation transport. A recent conservative, causal ray-tracing method was developed and combined with a short characteristic ray-tracing for the transfer calculations of ionizing radiation \\citep{mellema2006}. Solar surface magneto-convection simulations are increasingly realistic and use a three-dimensional, non-gray, approximate local thermodynamic equilibrium (LTE), radiative transfer for the heating and cooling of plasma. These simulations are typically formulated for four frequency bins in the radiative transport equation \\citep{vogler2005,stein2007,martinez2009}. For some applications, simplifications to the radiation transfer can be made by calculating moments of the radiation intensity over the solid angle $\\Omega$. The spectral radiation energy and the spectral radiation energy flux are defined by the 0th and 1st moments as \\begin{equation} E_\\nu({\\bf r},t,\\nu) = \\frac{1}{c}\\int_{4\\pi} I_\\nu({\\bf r},t,{\\bf n},\\nu) d\\Omega, \\qquad {\\bf F}_\\nu({\\bf r},t,\\nu) = \\int_{4\\pi} {\\bf n} I_\\nu({\\bf r},t,{\\bf n},\\nu) d\\Omega. \\end{equation} In addition, the spectral radiation pressure tensor ${\\bf P}_\\nu$ is defined by the second moment \\begin{equation} {\\bf P}_\\nu({\\bf r},t,\\nu) = \\frac{1}{c} \\int_{4\\pi} {\\bf n}{\\bf n} I_\\nu({\\bf r},t,{\\bf n},\\nu)d\\Omega. \\label{eq:pressuretensor} \\end{equation} A whole class of radiation transfer models are based on solving the corresponding radiation moment equations, using a closure relation between the spectral pressure tensor and the spectral intensity \\citep{mihalas1984,pomraming2005,drake2006}. A radiation-hydrodynamics code based on variable Eddington tensor (VET) methods \\citep{stone1992} can still capture the angular structure of the radiation field by relating the spectral radiation pressure tensor to the spectral radiation intensity and the method is applicable for both the optically thin and thick regime. Optically thin versions of the VET method have been used in the context of cosmological reionization \\citep{petkova2009}. Further simplification assumes that the radiation pressure is isotropic and proportional to the radiation energy. This is the diffusion approximation. Several codes have been developed using this approximation. HYDRA \\citep{marinak2001} is an arbitrary Lagrange Eulerian code for 2D and 3D radiation hydrodynamics. The radiation transfer model is based upon either flux limited multigroup or implicit Monte Carlo radiation transport. The Eulerian code RAGE \\citep{gittings2008} uses a cell-based adaptive mesh refinement to achieve resolved radiative hydrodynamic flows. HYADES \\citep{larsen1994} solve the radiation hydrodynamic equations on a Lagrangian mesh, while CALE \\citep{barton1985} can use either a fixed Eulerian mesh, an embedded Lagrangian mesh, or a partially embedded, partially remapped mesh. Our newly developed radiation-hydrodynamic solver uses an Eulerian grid together with a block-based adaptive mesh refinement strategy. We limit the discussion of the radiation hydrodynamics implementation in CRASH to plasmas in the absence of magnetic field. Most of the description in this paper can, however, easily be extended to magnetohydrodynamic (MHD) plasmas as well. Indeed, since the CRASH code is essentially the magnetohydrodynamic BATS-R-US code \\citep{powell1999,toth2010} extended with libraries containing radiation transport, equation-of-state (EOS), and opacity solvers, the implementation for the coupling between the radiation field and MHD plasmas is readily available. The CRASH code uses the recently developed block adaptive tree library (BATL, \\citet{toth2010}). Here we will focus on the radiation implementation. Both the CRASH and BATS-R-US codes are publicly available as part of the Space Weather Modeling Framework (SWMF, \\citet{toth2005}) or can be used as stand-alone codes. In the following, Section \\ref{sec:radhydro} introduces the radiation hydrodynamic equations for multi-material plasmas, in a form general enough to apply at high energy density. Section \\ref{sec:method} describes the numerical algorithms to solve these equations. Next, in Section \\ref{sec:verification} verification tests, are presented for radiation and electron heat conduction on non-uniform meshes in 1D, 2D, and 3D slab geometry and in axially symmetric ($rz$) geometry. We also show a full system multi-material radiation hydrodynamic simulation on an adaptively refined mesh and demonstrate good scaling up to 1000 processors. The paper is summarized in Section \\ref{sec:summary}. ", "conclusions": "\\label{sec:summary} We have extended the BATS-R-US code \\citep{powell1999,toth2010} with a new radiation transfer and heat conduction library. This new combination together with the equation-of-state and multigroup opacity solver is called the CRASH code. This code uses the recently developed parallel Block Adaptive Tree Library (BATL, see \\citet{toth2010}) to enable highly resolved radiation hydrodynamic solutions. The implemented radiation hydrodynamic schemes solve for the gray or multigroup radiation diffusion models in the flux limited diffusion approximation. In high energy density plasmas, the electrons are most of the time strongly coupled to the ions by collisions. An important exception is at hydrodynamic shocks, where the ions are heated by the shock wave and the electrons and ions are out of temperature equilibrium. Since radiative shocks are the main application for CRASH, we have implemented a separate electron pressure equation with the electron thermal heat conduction. For the electron heat conduction, we have added the option of a flux limiter to limit the thermal flux with the free-streaming heat flux. The multi-material radiation hydrodynamic equations are solved with an operator split method that consists of three substeps: (1) solving the hydrodynamic equations with standard finite volume shock-capturing schemes, (2) the linear advection of the radiation in frequency-logarithm space, and (3) the implicit solution of the radiation, heat conduction, and energy exchanges. For the implicit solver, standard Krylov solvers are used together with a Block Incomplete Lower-Upper decomposition (BILU) preconditioner. This preconditioner scales well up to 500 or 1000 processors. For future work, we may explore for the implicit multigroup diffusion a multi-level preconditioner to better scale the radiation solver beyond 1000 processors. We have presented a suite of verification tests that benchmark the performance. These tests verify the correctness and accuracy of the implementation for the gray and multigroup radiation diffusion algorithm and the heat conduction in 1D, 2D, and 3D slab and 2D $rz$ geometry. To demonstrate the full capability of the implementation, we have presented a 3D multi-material simulation of a radiative shock wave propagating through an elliptic nozzle. This configuration will be used in future validation studies. Since this radiation transfer library is an extension of the BATS-R-US code, the implementation is readily available for MHD simulations as well. This allows for validation studies of the radiation MHD implementation using laboratory-astrophysics experiments or for the simulations of astrophysical plasmas." }, "1101/1101.3875_arXiv.txt": { "abstract": "Motivated by the recent low-threshold measurements of the solar $^8$B neutrino spectrum performed by Borexino, Super-Kamiokande and the Sudbury Neutrino Observatory -- all now monitoring the transition regime between low-energy (vacuum-like) and high-energy (matter-dominated) flavor conversions -- we consider the role of sub-dominant dynamical terms induced by new flavor-changing interactions. We find that the presence of such perturbations with strength $\\sim 10^{-1}G_{F}$ is now favored, offering a better description of the anomalous behavior suggested by the new results, whose spectrum shows no sign of the typical low-energy upturn predicted by the standard MSW mechanism. Our findings, if interpreted in a 2-flavor scheme, provide a hint of such new interactions at the $\\sim 2 \\sigma$ level, which is rather robust with respect to 3-flavor effects possibly induced by non-zero $\\theta_{13}$. ", "introduction": " ", "conclusions": "" }, "1101/1101.1088_arXiv.txt": { "abstract": "We present the observed offsets of short-duration gamma-ray bursts (SGRBs) from their putative host galaxies and compare them to the expected distributions of merging compact object binaries, given the observed properties of the hosts. We find that for all but one burst in our sample the offsets are consistent with this model. For the case of bursts with massive elliptical host galaxies, the circular velocities of the hosts' haloes exceed the natal velocities of almost all our compact object binaries. Hence the extents of the predicted offset distributions for elliptical galaxies are determined largely by their spatial extents. In contrast, for spiral hosts the galactic rotation velocities are smaller than typical binary natal velocities and the predicted burst offset distributions are more extended than the galaxies. One SGRB, 060502B, apparently has a large radial offset that is inconsistent with an origin in a merging galactic compact binary. Although it is plausible that the host of GRB~060502B is mis-identified, our results show that the large offset is compatible with a scenario where at least a few per cent of SGRBs are created by the merger of compact binaries that form dynamically in globular clusters. ", "introduction": "The merger of a compact binary consisting of neutron stars (NSs) or black holes (BHs) is a prime model for the origin of the short-duration gamma-ray bursts (SGRBs). Compact binaries may be formed by numerous channels including those which are formed in primordial binaries \\citep{bel2002b}, and those which are formed dynamically in dense cluster cores \\citep[e.g.][]{davies1995,grindlay2005}. Once formed the binary loses energy and angular momentum from its orbit in the form of gravitational radiation, and the compact objects spiral together and eventually merge. The merger time is highly sensitive to both separation and eccentricity and thus has an extremely broad distribution. Evidence for the origin of SGRBs in the final coalescence of such systems potentially comes both from the host galaxy types \\citep[e.g.][]{zheng2007} and the measured offsets of SGRBs from their hosts \\citep[e.g.][]{bloom99,bel2006,troja08}. The kicks imparted to double compact object (DCO) systems on formation, both in terms of dynamical kicks to the binary and natal kicks to neutron stars, endow DCO binaries with spatial velocities of up to several hundred $\\kms$. This means that they may merge far from their birth sites, even outside their parent galaxies. Equally, if the inspiral takes $>10^8\\,{\\rm years}$, then the host galaxy may no longer exhibit significant ongoing star formation at the time of merger. Observations of SGRBs obtained by {\\it Swift} so far offer broad support for this picture. Several of the host galaxies observed are old \\citep{gehrels2005,berger2005,barthelmy2005,bloom2006}, while only a small fraction are starbursts. However, the sample has been building up slowly, and we are only now approaching a stage where it is possible to compare the observed offsets with predictions for DCO binaries. We build on previous work that modelled the offset distributions of DCO mergers \\citep[e.g.][]{bloom99,fryer99} by using the observed properties of the host galaxies to predict offset distributions on a host-by-host basis. \\begin{table*} \\begin{center} \\begin{tabular}{llllllllllllll} \\hline GRB & Satellite & $t_{90}$ & Opt & Host Mag. & Host $M_B$ & z& Type & $R_{\\rm e}$ & $R_{\\rm proj}$& $R_{\\rm err}$& References\\\\ & & s & & & & & & kpc & kpc & kpc &\\\\ \\hline 050509B & {\\it Swift} & 0.04 & No & 16.75 $\\pm$ 0.05 & $-23.25$ & 0.225 & E & 20.98 & 63.7 & 12.1 & (a), (b)\\\\ 050709 & {\\it HETE-2} & 0.07 & Yes & 21.05 $\\pm$ 0.07 & $-18.19$ & 0.161 & S/I & 1.75 & 3.55 & 0.27 & (c), (d), (b)\\\\ 050724 & {\\it Swift} & 3.0 & Yes & 18.19 $\\pm$ 0.03 & $-22.11$ & 0.257 & E & 4.00 & 2.54 & 0.08 & (e), (b)\\\\ 051221A & {\\it Swift} & 1.4 & Yes & 21.81 $\\pm$ 0.09 & $-20.20$ & 0.546 & S & 2.17 & 1.53 & 0.31 & (f), (g)\\\\ 060502B & {\\it Swift} & 0.05 & No & 18.71 $\\pm$ 0.01 & $-21.84$ & 0.287 & E & 10.5 & 73 & 13 & (h)\\\\ 060801 & {\\it Swift} & 0.5 & No & 22.97 $\\pm$ 0.11 & $-20.64$ & 1.130 & S & - & 19.7 & 14.0 & (i), (j)\\\\ 061006 & {\\it Swift} & 0.4 & Yes & 22.65 $\\pm$ 0.09 & $-18.86$ & 0.438 & S & 3.67 & 1.44 & 0.29 & (b)\\\\ 061201 & {\\it Swift} & 0.8 & Yes & 19.65 $\\pm$ 0.10 & $-18.76$ & 0.111 & S & 1.8 & 33.9 & 0.4 & (i), (j)\\\\ 061210 & {\\it Swift} & 0.2 & No & 21.00 $\\pm$ 0.02 & $-20.36$ & 0.410 & S & - & 10.7 & 6.9 & (i), (j)\\\\ 061217 & {\\it Swift} & 0.3 & No & 23.33 $\\pm$ 0.07 & $-19.61$ & 0.827 & S & - & 55 & 20 & (i), (j)\\\\ 070429B & {\\it Swift} & 0.5 & Yes & 23.22 $\\pm$ 0.10 & -19.91 & 0.902 & S & - & 4.7 & 4.7 & (k)\\\\ 070714B & {\\it Swift} & 2.0 & Yes & 24.92 $\\pm$ 0.23 & -18.26 & 0.923 & S & - & 3.08 & 0.47 & (l)\\\\ 070724A & {\\it Swift} & 0.4 & Yes & 20.53 $\\pm$ 0.03 & -21.07 & 0.457 & E & - & 4.76 & 0.06 & (m)\\\\ 070809 & {\\it Swift} & 1.3 & Yes & 21.7 $\\pm$ 0.3 & -18.23 & 0.219 & S & - & 19.61 & 1.9 & (n)\\\\ 071227 & {\\it Swift} & 1.8 & Yes & 20.54 $\\pm$ 0.03 & -20.73 & 0.394 & S & - & 16.1 & 0.2 & (o)\\\\ 080905A & {\\it Swift} & 1.0 & Yes & 18.0 $\\pm$ 0.5 & -20.63 & 0.122 & S & - & 18.11 & 0.42 & (p)\\\\ \\hline \\end{tabular} \\end{center} \\caption{ Properties of SGRBs with known X-ray, optical or radio counterparts. Data is taken from the literature sources indicated other than for a small number of the offsets which we have re-measured in described in Section~2. Columns list in order: the burst identifier, the satellite which detected the burst, the duration over which 90\\% of the total fluence was seen ($t_{90}$), whether an optical counterpart was detected, the host apparent $R$-band magnitude and error, the host absolute $B$-band magnitude, the redshift, the host type (spiral/elliptical), the host effective radius $R_{\\rm e}$, the offset $R_{\\rm proj}$ and the error on the offset $R_{\\rm err}$. The duration, $t_{90}$, is measured in the 15-350 keV range for the {\\it Swift} bursts and in the 30-400 keV range for the {\\it HETE-2} burst. The absolute magnitudes have been calculated assuming a source with a flat spectrum in $F_{\\nu}$. All offset errors are one-sigma; 90\\% confidence limits in the literature have been converted assuming a Gaussian PSF. References: (a) \\citet{gehrels2005}; (b) \\citet{fong09}; (c) \\citet{fox05}; (d) \\citet{villasenor05}; (e) \\citet{berger2005}; (f) \\citet{soderberg06}; (g) this work; (h) \\citet{bloom2007}; (i) \\citet{bergerhiz}; (j) \\citet{troja08}; (k) \\citet{Cenko09}; (l) \\citet{graham09}; (m) \\citet{berger09}; (n) Perley et al. (GCN 7889); (o) \\citet{davanzo09}; (p) Rowlinson et al. in prep. } \\label{tab:bursts} \\end{table*}% In this paper we take the observed sample of SGRBs that have identified hosts and model the production and galactic trajectories of DCO binaries inside those hosts. In Section 2 we review the sample of bursts; we describe our models in Section 3, discussing the results and possible selection effects in Section 4. The alternative possibility of DCO binaries forming dynamically in the cores of globular clusters is discussed in Section 5. ", "conclusions": "" }, "1101/1101.5568_arXiv.txt": { "abstract": " ", "introduction": "% Polars are a subclass of the cataclysmic variables (CVs) in which the accreted material is channeled by the magnetic field of the white dwarf (WD) forming an accretion column. Near the WD surface the material forms a shock front increasing temperature and densities. The bulk of the emission comes from the base of the column also called post-shock region. The mean emission processes are cyclotron (optical and infrared) and bremsstrahlung (X-ray). The radiative cooling results in a structured region with temperature and density varying as a function of the height of the post-shock region. The X-ray emission of polars is commonly reproduced combining shock models and plasma emission models, as Meka model \\cite{mewe:1986} available in the {\\sc xspec} Spectral Fitting Package \\cite{arnoudi:1996}. First, the post-shock region is divided in layers and each layer has a temperature and density given by the post-shock structure. Then, the emissivity is calculated for each layer and the total emission is the sum of the plasma emissivity of all layers, assuming an optically thin region \\cite{cropper:2000}. The X-ray fittings are usually done without considering the optical data, in part because distinct codes are used to model each spectral range. An exception is the modeling of polar RX J2115-5840 \\cite{ramsay:2000}. The cyclotron 2D emitting region, obtained from the fit of optical data, is used to define the location of the X-ray emission. However, the density and temperature profiles adopted in fitting each energy band are not the same. % The absorption due the pre-shock flow located above the post-shock region is important in high energies. Customary, the absorption, which varies with orbital phase, is arbitrarily choosen in order to fit the observational data. The {\\sc cyclops} code solves the radiative transfer of the cyclotron process in the post-shock region adopting a 3D treatment and specific functions to describe the variation of temperature and density due the radiative cooling. The code provides the four Stokes parameters which can be used to reproduce optical light and polarimetric curves from polars \\cite{costa:2009}. Here we present the modification of {\\sc cyclops} to include bremsstrahlung emission and to prepare the code to simultaneously model the optical and X-ray continuum emission of post-shock regions of polars. ", "conclusions": "We present the {\\sc cyclops-x}, a code that calculates the optical and X-ray emission of post-shock regions of polars. This code includes the absorption caused by the pre-shock region in a manner consistent with the geometry defined by the magnetic lines. We used the {\\sc cyclops-x} to calculate the X-ray light curves of the two good optical models for CP Tuc. The optical emission of the two models are very similar, what is not true for the X-ray fluxes. This simple application shows the importance of consider as many data as possible to understand the post-shock region of a polar. We are planing to incorporate in {\\sc cyclops-x} the simultaneous fit of X-ray and optical data and, then, apply this tool to the CP Tuc multiwavelength data." }, "1101/1101.2512_arXiv.txt": { "abstract": "{} {Stellar activity is a potential important limitation to the detection of low mass extrasolar planets with indirect methods (RV, photometry, astrometry). In previous papers, using the Sun as a proxy, we investigated the impact of stellar activity (spots, plages, convection) on the detectability of an Earth-mass planet in the habitable zone (HZ) of solar-type stars with RV techniques. We extend here the detectability study to the case of astrometry. } {We used the sunspot and plages properties recorded over one solar cycle to infer the astrometric variations that a Sun-like star seen edge-on, 10 pc away, would exhibit, if covered by such spots/bright structures. We compare the signal to the one expected from the astrometric wobble (0.3 $\\mu$as) of such a star surrounded by a one Earth-mass planet in the HZ. We also briefly investigate higher levels of activity. } { The activity-induced astrometric signal along the equatorial plane has an amplitude of typ. less than 0.2 \\muas (rms=0.07 \\muas), smaller than the one expected from an Earth-mass planet at 1 AU. Hence, for this level of activity, the detectability is governed by the instrumental precision rather than the activity. We show that for instance a one Earth-mass planet at 1 AU would be detected with a monthly visit during less than 5 years and an instrumental precision of 0.8 \\muas. A level of activity 5 times higher would still allow such a detection with a precision of 0.35 \\muas. We conclude that astrometry is an attractive approach to search for such planets around solar type stars with most levels of stellar activity. } {} ", "introduction": "Stellar activity is now recognized as a potentially strong limitation for the indirect detection of planets. Indeed, spots and bright structures (plages, network) produce brightness inhomogeneities at the stellar surface that affect the photometric, astrometric and radial velocity (RV) signals. The RV signal is also affected by the inhibition of convection in the active area. The amplitude of the activity-related stellar noise depends on the activity pattern and intensity, which is related to the stellar properties (age, temperature, etc). In the cases of young, active stars, or late-type stars, the signal may mimic that of a giant planet with periods of the order of the star rotational period. In the case of less active, solar-type Main-Sequence stars, the noise is much lower, but can nevertheless affect the detection of terrestrial planets. In two recent papers, we have investigated the impact of spots (\\cite{lagrange2010}; herefater paper I), plages and convection (\\cite{meunier2010}; herefater paper II) on the detectability of an Earth mass planet located in the habitable zone (HZ) of the Sun, as seen edge-on and observed in RV or in photometry. To do so, we took into account all spots and plages recorded over one solar cycle. We showed that providing a very tight and long temporal sampling (typically twice a week over more than 2 orbital periods) and an RV precision in the 10 cm/s range, the photometric contribution of plages and spots should not prevent the RV detection of Earth-mass planets in the HZ. This is no longer true when convection is taken into account. Here, we complete the previous work by estimating the astrometric signal produced by the same structures over the same solar cycle, and we compare it to the astrometric signal induced by an Earth-mass planet located at 1 AU from the Sun. We also compare this signal to the RV and photometric signals. ", "conclusions": "\\subsection{Planet detection in astrometry} We see in Fig.~\\ref{shift_variations} that the activity-induced shifts are most of the time smaller than 0.2 $\\mu$as, ie smaller than the amplitude of the shifts induced by an Earth mass planet at 1 AU (0.33 \\muas along the equatioral plane and 0 \\muas perpendicular to the equator). These values are in agreement with the estimations of \\cite{makarov2010}. The rms of the equatorial shift over the entire cycle is $\\simeq$ 5 times smaller than the amplitude of the Earth signal, and 4 times smaller during high activity. This means that the Sun activity would not prevent to detect an Earth located in the HZ, provided the precision of the data allows such mesurements. This is illustrated in Fig.~\\ref{perio} where we show the periodogramme of the astrometric shifts along the equatorial plane once a one Earth-mass planet has been added (no noise). In this example, we use a limited amount of data instead of the whole set of data, in order to be closer to a real case. The chosen temporal sampling is of 1 month $^{+}_{-}$ 5 days over 50 months, typical of the strategy adopted for the NEAT instrument recently proposed to ESA (\\cite{leger2010}), dedicated to a systematic astrometric search for exo-Earths in the HZ of nearby stars. We then explored different instrumental noises. An example is given in the lower panel of Fig.~\\ref{perio} where we assume a noise of 0.8 \\muas \\footnote{this value corresponds to the precision that NEAT will provide in 1 hour observation. In practice, each visit will be longer (up to $\\simeq$ 5 hours) to ensure a better precision (precision is proportional to $\\sqrt{t_{\\rm obs}}$). }. The temporal sampling is the same as before. The planet is still detectable. However, in that case, the correlation seen in the previous section between the astrometric shift along the equatorial plane and the RV variations is not any longer present. This is due to the fact that the instrumental noise is larger than the activity-induced noise. The consequence is that if we take the instrumental noise into account, for this level of activity, it will not be possible to estimate the level of convection using the astrometric and RV data. \\begin{figure}[t] \\centering \\begin{tabular}{cc} \\includegraphics[angle=0,width=.45\\hsize]{deltax_terre_low_periodogram_obs}& \\includegraphics[angle=0,width=.45\\hsize]{deltax_terre_high_periodogram_obs}\\\\ \\includegraphics[angle=0,width=.45\\hsize]{deltax_terre_bruit_low_periodogram_obs}& \\includegraphics[angle=0,width=.45\\hsize]{deltax_terre_bruit_high_periodogram_obs}\\\\ \\end{tabular} \\caption{Periodogrammes of the shift along the equatorial plane of the Sun viewed edge-on and surrounded by a one Earth-mass planets at 1AU, and observed once a month approximately, during 50 months. From top to bottom: 1) low activity period, no instrumental noise; 2) high activity period, no instrumental noise; 3) low activity period, 0.8 \\muas rms instrumental noise; 4) high activity period, 0.8 \\muas rms instrumental noise.} \\label{perio} \\end{figure} We finally explore different levels of activity. We found that an instrumental precision similar to that expected for the NEAT instrument for a 5 hours-long visit of a G2 star located 10 pc away (0.35 \\muas) will easily allow the detection even for stars 5 times more active than the Sun (we assume conservatively that such a star will produce shifts five times higher than the Sun) \\footnote {in practice, a level of activity five times higher will probably induce shifts smaller than 5 times the shifts, as statistically, there will be more chances that the effects of structures located on a given side of the hemisphere be balanced by that those of structures on the opposite side}. This is illustrated in Fig.~\\ref{perio_5act} where we show the periodogrammes of the shift along the equatorial plane (same sampling as before) during the low and high activity periods. Given \\cite{lockwood07} relation between the relative photometric variations and the Ca H \\& K line strength index log(R'HK), and given the log(R'HK) for the Sun (between $\\simeq$ $-$4.85 and $-$5.0 respectively during the high and low activity phases), this corresponds to a log(R'HK) smaller than $-$4.25. % Most of active stars fall into this regime. Indeed, among the 385 Hipparcos stars closer than 20 pc with a known R'HK, only 8 have log(R'HK) $\\geq-$4.2 (see Fig.~\\ref{stats_hip}; \\cite{maldonado10})). \\begin{figure}[t] \\centering \\includegraphics[angle=0,width=.45\\hsize]{5act_noisep35_deltax_terre_bruit_low_periodogram_obs} \\includegraphics[angle=0,width=.45\\hsize]{5act_noisep35_deltax_terre_bruit_high_periodogram_obs}\\\\ \\caption{Periodogrammes of the shift along the equatorial plane assuming a Sun 5 times more active than observed, viewed edge-on and surrounded by a one Earth-mass planets at 1 AU, and observed once a month approximately, during 50 months. A 0.35 \\muas rms instrumental noise has been added to the data. Left: low activity period; right: high activity period.} \\label{perio_5act} \\end{figure} \\begin{figure}[t] \\centering \\includegraphics[angle=0,width=.8\\hsize]{paper_astrometry_hipparcos_activity}\\\\ \\caption{Histogram of log(R'HK) for all Hipparcos stars closer than 20 pc, with a known R'HK.} \\label{stats_hip} \\end{figure} So far we have assumed that the angle of the planet orbit (supposedly identical to the equatorial plane) is known. In the case of real observations though, this information will not be known a priori and the planet may orbit outside the star's equatorial plane. It appears that even in the presence of stellar activity and instrumental noise at a level of 0.35 \\muas per measurement, the astrometric data will allow to recover this information. We consider the example of a planet on an orbit inclined by -60 degrees with respect to the North. The star is supposed to have either the same activity level as the Sun, or 5 times this level. The astrometric shifts are measured in a referential which is allowed to rotate from 0 (equatorial plane) to 360 degrees. We measure in each case the shifts projected on this referential, and the associated rms along each axis of the referential. We consider either no or 0.35 \\muas level instrumental noise, and either a solar-like level of activity or 5 times this level. The temporal sampling is the same as before (50 data points). We show in Fig.~\\ref{fig_angle} the rms as a function of the rotation angle of the referential. We see a clear maximum at an angle equal to the planet inclination. Hence the position angle of the planet orbit can be retrieved. \\begin{figure}[t] \\centering \\includegraphics[angle=0,width=.8\\hsize]{fig_angle2}\\\\ \\caption{Rms of the shift along the x-axis of a rotating referential as a function of its rotational angle (see text). In red, case of solar activity. In blue, case of 5 times the solar activity. In all cases, we consider that the instrumental noise is 0 (plain lines) or 0.35 microas (dashed lines). The vertical dashed line shows the angle of the inclination of the planet orbital plane. } \\label{fig_angle} \\end{figure} We conclude that in most cases, stellar activity will not be the dominant factor for the planet detectability in astrometry. The dominant factor will rather be the instrumental precision. Note that if present, other larger planets will also induce noises, which are not considered in this paper, but the series of blind tests conducted to estimate the detection capabilities of Earth-like planets in multiple systems by space-borne astrometry (\\cite{traub2010}) have showed that this noise can be well circumvented. \\subsection{Comparison to RV data} The situation is then different from the RV data, as the rms of the RV due to spots and bright structures only (i.e. convection not taken into account) during resp. the entire cycle, the low activity period and the high activity periods are 3, 1 and 4 times larger than the amplitude of the Earth RV signal (see Table~\\ref{stats}). When taking convection into account, the corresponding ratio become 24, 4 and 14, respectively. Therefore, for RV techniques, the detection limit is set by the stellar activity (and mainly by convection effects) rather than the instrumental precision (provided precisions of 10 cm/s are available). To prepare and accompany a space astrometric misson dedicated to Earth-mass planet detection, a number of previous RV, or photometric monitoring would certainly be very precious. In particular, RV measurements are very powerful to derive the properties of the star's activity. They are in fact more powerfull than astrometry as the rms of the activity-induced RV variations is much larger than the instrumental precision (conversely to astrometry). Previous long-term RV monitoring can allow to identify the presence of larger bodies in the system and characterize them (note that lower precision astrometric monitoring can also do this). During the astrometric measurements, simultaneous RV observations with a precision of about 10 cm/s would provide detailed information on the actual activity of the star and will help the analysis of the astrometric data." }, "1101/1101.5036.txt": { "abstract": "The interest in X/$\\gamma$-ray Astronomy has grown enormously in the last decades thanks to the ability to send X-ray space missions above the Earth's atmosphere. There are more than half a million X-ray sources detected and over a hundred missions (past and currently operational) devoted to the study of cosmic X/$\\gamma$ rays. With the improved sensibilities of the currently active missions new detections occur almost on a daily basis. Among these, neutron-star X-ray binaries form an important group because they are among the brightest extra-solar objects in the sky and are characterized by dramatic variability in brightness on timescales ranging from milliseconds to months and years. Their main source of power is the gravitational energy released by matter accreted from a companion star and falling onto the neutron star in a relatively close binary system. Neutron-star X-ray binaries divide into high-mass and low-mass systems according to whether the mass of the donor star is above $\\sim$8 or below $\\sim$2 $\\msun$, respectively. Massive X-ray binaries divide further into supergiant X-ray binaries and Be/X-ray binaries depending on the evolutionary status of the optical companion. Virtually all Be/X-ray binaries show X-ray pulsations. Therefore, these systems can be used as unique natural laboratories to investigate the properties of matter under extreme conditions of gravity and magnetic field. The purpose of this work is to review the observational properties of Be/X-ray binaries. The open questions in Be/X-ray binaries include those related to the Be star companion, that is, the so-called \"Be phenomenon\", such as, timescales associated to the formation and dissipation of the equatorial disc, mass-ejection mechanisms, V/R variability, and rotation rates; those related to the neutron star, such as, mass determination, accretion physics, and spin period evolution; but also, those that result from the interaction of the two constituents, such as, disc truncation and mass transfer. Until recently, it was thought that the Be stars' disc was not significantly affected by the neutron star. In this review, I present the observational evidence accumulated in recent years on the interaction between the circumstellar disc and the compact companion. The most obvious effect is the tidal truncation of the disc. As a result, the equatorial discs in Be/X-ray binaries are smaller and denser than those around isolated Be stars. ", "introduction": "In very general terms, one can simply define X-ray binaries as systems that consist of a compact object orbiting an optical companion. They are \"close\" binary systems because there exists a transfer of mass from the optical component to the compact object. By \"optical companion\" it is understood that nuclear burning is still taking place in its interior. Figure~\\ref{xrb_class} shows a tree-diagram depicting all the different subsystems that comprise the generic group of X-ray binaries. In referring to the two components in an X-ray binary one should be careful and learn which is the subject of investigation as the same name can be used to mean different components. In massive X-ray binaries, the most massive star is normally termed {\\em primary} whereas the less massive one is called {\\em secondary}. In low-mass systems, the term {\\em primary} refers to the neutron star while the word {\\em secondary} is reserved for the late-type companion. Other names also used are \"optical\" or \"donor\" for the larger star and \"compact\", \"gainer\" or \"accreting\" for the denser companion. They admit several classification schemes depending upon whether the emphasis is put on the type of the compact companion or the physical properties of the optical star. X-ray binaries divide up into black-hole systems, neutron star X-ray binaries or cataclysmic variables (if the compact object is a white dwarf). Nevertheless, the term \"X-ray binaries\" is normally reserved to designate binaries with neutron stars. %----------------------FIG 1------------------------------------- \\begin{figure}[t] \\includegraphics[width=0.9\\linewidth]{./preig_fig01.eps} \\caption{Classification of X-ray binaries.} \\label{xrb_class} \\end{figure} %----------------------------------------------------------- \\subsection{High-mass X-ray binaries} \\label{intro_hmxb} Neutron-star X-ray binaries are divided up into high-mass (HMXBs) and low-mass (LMXBs) X-ray binaries depending on the spectral type of the mass donor, as this feature determines the mode of transferring mass to the compact object and the environment surrounding the X-ray source. HMXBs contain early-type (O or B) companions, while the spectral type of the optical star in LMXBs is later than A. HMXBs are strong emitters of X-ray radiation. Sometimes they appear as the brightest objects of the X-ray sky. The high-energy radiation is produced as the result of accretion of matter from the optical companion onto the neutron star. The term accretion refers to the gradual accumulation or deposition of matter onto the surface of an object under the influence of gravity. If the accreting object is a neutron star (or black hole), then matter falls down onto an enormous well of gravitational potential and is accelerated to extremely high velocities. When the matter reaches the surface of the neutron star, it is rapidly decelerated and the free-fall kinetic energy radiated away as heat which is available to power the X-ray source. The luminosity class serves to subdivide HMXBs into Be/X-ray binaries (\\bex), when the optical star is a dwarf, subgiant or giant OBe star (luminosity class III, IV or V) and supergiant X-ray binaries (SGXBs), if they contain a luminosity class I-II star. In SGXBs, the optical star emits a substantial stellar wind, removing between $10^{-6}-10^{-8}$ M$_{\\odot}$ yr$^{-1}$ with a terminal velocity up to 2000 km s$^{-1}$. A neutron star in a relatively close orbit will capture some fraction of this wind, sufficient to power a bright X-ray source. If mass transfer occurs via Roche lobe overflow, then the X-ray emission is highly enhanced and an accretion disc is formed around the neutron star. At present, there is known only one disc-fed SGXB in the Galaxy (Cen X-3) and three in total (SMC X-1 and LMC X-4), while there are about a few tens of wind-fed SGXBs. Because of their brightness and persistent X-ray emission, SGXBs were the first to be discovered. They were initially thought to represent the dominating population of HMXBs, whereas \\bex\\ were considered atypical cases. Hence, the name {\\em classical} or {\\em standard} was given to SGXBs. In \\bex, the optical companion is a Be star. Be stars are non-supergiant fast-rotating B-type and luminosity class III-V stars which at some point of their lives have shown spectral lines in emission, hence the qualifier \"$e$\" in their spectral types \\citep{port03, balo00, slet88}. The best studied lines are those of hydrogen (Balmer and Paschen series) but Be stars may also show He, Fe in emission \\citep[see e.g][]{hanu96}. They also show an amount of IR radiation than is larger than that expected from an absorption-line B star of the same spectral type. This extra long-wavelength emission is known as infrared excess. The origin of the emission lines and infrared excess in \\bex\\ is attributed to an equatorial disc, fed from material expelled from the rapidly rotating Be star in a manner that it is not yet understood \\citep{port03}. During periastron, the neutron star passes close to this disc, sometimes may even go through it causing major disruption. A large flow of matter is then accreted onto the compact object. The conversion of the kinetic energy of the in-falling matter into radiation powers the X rays. \\bex\\ have large orbital periods and by definition they are non-supergiant systems. Hence, the Be star is well within the Roche lobe. However, transient Roche lobe overflow may occur during periastron passage in systems with large eccentric orbits or during giant X-ray outbursts when a large fraction of the Be star's disc is believed to be accreted. Most \\bex\\ are transient systems and present moderately eccentric orbits ($e\\simmore0.3$), although persistent sources also exist \\citep{reig99}. Persistent \\bex\\ differ from transient \\bex\\ in that they display much less X-ray variability (no large outbursts are detected), lower luminosity ($L_x \\simless 10 ^{35}$ erg s$^{-1}$), contain slowly rotating neutron stars ($P_{\\rm spin} > 200$ s) and reside in wider orbit systems ($P_{\\rm orb} > 200$ d). Table \\ref{bexlist} lists the confirmed \\bex. This table includes only those systems with identified optical counterparts and whose X-ray and optical/IR variability has been shown to be typical of \\bex\\ (see Sects.~\\ref{optir} and \\ref{xrayvar}). Table \\ref{stat} gives the number of various types of X-ray binaries. There are more than 300 bright X-ray sources with fluxes well above $10^{-10}$ erg cm$^{-2}$ s$^{-1}$ in the energy range 1--10 keV \\citep{liu06,liu07,bird10}. The distribution of these sources shows a clear concentration towards the Galactic center and also towards the Galactic plane, indicating that the majority belong indeed to our Galaxy. %---------------------------------------------------------------------------- \\begin{deluxetable}{lllllccccccc} \\tabletypesize{\\scriptsize} \\rotate \\label{bexlist} \\tablecaption{List of galactic \\bex. Only systems with known optical counterparts and well-established optical/X-ray behaviour are included} \\tablewidth{0pt} \\tablehead{ \\colhead{X-ray name}\t&\\colhead{Optical Ctp.}&\\colhead{RA (J2000)}\t&\\colhead{Dec (J2000)} &\\colhead{Spec. type}\t&\\colhead{V}\t&\\colhead{J}\t\t&\\colhead{E(B-V)} &\\colhead{$P_{\\rm spin}$ (s)}\t &\\colhead{$P_{\\rm orb}$ (d)} &\\colhead{e} &\\colhead{d} (kpc)} \\startdata %4U 0054+60\t&$\\gamma$ Cas\t&00 56 42.53\t&+60 43 00.3\t &B0.5 IVe\t&1.6-3.0 &2.04 \t&0.05\t &--\t &203.59 &0-0.26 &0.19 \\\\ 4U 0115+634\t&V635 Cas\t&01 18 31.90\t&+63 44 24.0\t &B0.2 Ve \t&14.8-15.5&10.8-12.3\t&1.55\t &3.6\t &24.3 &0.34 &8 \\\\ IGR J01363+6610\t&--\t\t&01 36 18.00\t&+66 10 36.0\t &B1 IV-Ve\t&13.3\t &10\t \t&1.61\t &--\t &--\t &-- &2 \\\\ RX J0146.9+6121 &LS I +61 235\t&01 47 00.17\t&+61 21 23.7\t &B1Ve \t&11.2\t &10\t \t&0.93\t &1412 &330\u00a0 &-- &2.3 \\\\ IGR J01583+6713\t&--\t\t&01 58 18.20\t&+67 13 26.0\t &B2IVe \t&14.4\t &11.5 \t&1.46\t &19692 &--\t &-- &4 \\\\ RX J0240.4+6112\t&LS I +61 303\t&02 40 31.67\t&+61 13 45.6\t &B0.5 Ve \t&10.7\t &8.8\t \t&0.75\t &--\t &26.45 &0.54 &3.1 \\\\ V 0331+530\t&BQ Cam\t\t&03 34 59.89\t&+53 10 23.6\t &O8-9 Ve \t&15.1-15.8&11.4-12.2\t&1.9\t &4.4\t &34.3 &0.3 &7 \\\\ 4U 0352+309 (X Per)&HD 24534\t&03 55 23.08\t&+31 02 45.0\t &O9.5IIIe-B0Ve &6.1-6.8 &5.7-6.5 \t&0.4\t &837\t &250 &0.11 &1.3 \\\\ RX J0440.9+4431\t&LS V +44 17\t&04 40 59.32\t&+44 31 49.3\t &B0.2Ve \t&10.8\t &9.2\t \t&0.65\t &203\t &--\t &-- &3.3 \\\\ 1A 0535+262\t&V725 Tau\t&05 38 54.57\t&+26 18 56.8\t &B0 IIIe \t&8.9-9.6 &7.7-8.5 \t&0.75\t &105\t &111 &0.47 &2.4 \\\\ IGR J06074+2205\t&--\t\t&06 07 24.00\t&+22 05 00.0\t &B0.5V \t&12.3\t &10.5\t \t&1.88\t &--\t &--\t &-- &5 \\\\ XTE J0658-073\t&--\t\t&06 58 42.00\t&--07 11 00.0\t &O9.7 Ve \t&12.1 &9.7\t \t&1.19\t &160.4 &--\t &-- &3.9 \\\\ 4U 0726-260\t&V441 Pup\t&07 28 53.60\t&--26 06 29.0\t &O8-9Ve \t&11.6\t &10.4 \t&0.73\t &103.2 &34.5 &-- &6 \\\\ RX J0812.4-3114\t&LS 992\t\t&08 12 28.84\t&--31 14 52.2\t &B0.2 III-IVe &12.4 &11.2-12.0\t&0.65\t &31.89 &80\t &-- &9 \\\\ GS 0834-430\t&--\t\t&08 35 55.00\t&--43 11 06.0\t &B0-2 III-Ve\t&20.4\t &13.3\t \t&4.0\t &12.3 &105.8 &0.12 &5 \\\\ GRO J1008-57\t&--\t\t&10 09 46.00\t&--58 17 30.0\t &O9e-B1e \t&15.3 &10.9 \t&1.9\t &93.5 &247.5 &0.66 &2 \\\\ RX J1037.5-5647\t&LS 1698\t&10 37 35.50\t&--56 48 11.0\t &B0 III-Ve\t&11.3\t &--\t \t&0.75\t &862\t &--\t &-- &5 \\\\ 1A 1118-615\t&Hen 3-640\t&11 20 57.21\t&--61 55 00.3\t &O9.5 III-Ve\t&12.1\t &9.6\t \t&0.92\t &405\t &--\t &-- &5 \\\\ 4U 1145-619\t&V801 Cen\t&11 48 00.02\t&--62 12 24.9\t &B1 Vne \t&9.3\t &8.7\t \t&0.29\t &292.4 &187.5 &$>$0.5 &3.1 \\\\ 4U 1258-613\t&GX 304-1\t&13 01 17.20\t&--61 36 07.0\t &B2 Vne \t&13.5\t &9.8\t \t&2.0\t &272\t &132.5 &$>$0.5 &2.4 \\\\ 2S 1417-624\t&--\t\t&14 21 12.80\t&--62 41 54.0\t &B1 Ve \t&17.2\t &13.3 \t&2.0\t &17.6 &42.12 &0.45 &10 \\\\ GS 1843+00\t&--\t\t&18 45 36.90\t&+00 51 48.2\t &B0-2 IV-Ve\t&20.9\t &13.7\t \t&2.5\t &29.5 &--\t &-- &$>$10 \\\\ XTE J1946+274\t&--\t\t&19 45 39.30\t&+27 21 55.0\t &B0-1 IV-Ve\t&16.8\t &12.5\t \t&2.0\t &15.8 &169.2 &0.33 &8-10 \\\\ KS 1947+300\t&--\t\t&19 49 30.50\t&+30 12 24.0\t &B0 Ve \t&14.2\t &11.7 \t&1.09\t &18.76 &40.4 &0.03 &9.5 \\\\ EXO 2030+375\t&V2246 Cyg\t&20 32 15.20\t&+37 38 15.0\t &B0e\t \t&19.7\t &12.1 \t&3.8\t &41.8 &46.03 &0.41 &5 \\\\ GRO J2058+42\t&--\t\t&20 59 00.00\t&+41 43 00.0\t &O9.5-B0 IV-Ve &14.9\u00a0 &11.7 \t&1.38\t &192\t &110 &-- &9 \\\\ SAX J2103.5+4545&--\t\t&21 03 35.71\t&+45 45 05.5\t &B0 Ve \t&14.2\t &11.8 \t&1.35\t &358.6 &12.67 &0.4 &6.8 \\\\ 4U 2135+57\t&Cep X-4\t&21 39 30.72\t&+56 59 10.0\t &B1-B2 Ve\t&14.2\t &11.8 \t&1.3\t &66.3 &--\t &-- &3.8 \\\\ SAX J2239.3+6116&--\t\t&22 39 20.90\t&+61 16 26.8\t &B0-2 III-Ve\t&15.1\t &11.5 \t&1.4\t &1247 &262.6 &-- &4.4 \\\\ \\enddata %\\end{center} %\\end{table*} \\end{deluxetable} %---------------------------------------------------------------------------- %---------------------------------------------------------------------------- \\begin{table*} \\begin{center} \\caption{Statistics on HMXBs in the Milky Way} \\label{stat} \\begin{tabular}{@{}lc@{}} \\tableline Number of neutron-star X-ray binaries\\tablenotemark{1}\t\t&327\t\\\\ Number of suspected HMXB\\tablenotemark{1}\t\t\t&131\t\\\\ Number of suspected \\bex\\tablenotemark{2}\t\t\t&63\t\\\\ Number of {\\em confirmed} \\bex\\tablenotemark{3}\t\t\t&28\t\\\\ %Number of {\\em confirmed} \\bex\\ pulsars$^3$\t&\t\\\\ \\tableline \\end{tabular} \\tablenotetext{1}{Sources in the \\citet{liu07} and \\citet{liu06} catalogs plus update from the 4th IBIS/ISGRI soft gamma-ray survey catalog \\citep{bird10}} \\tablenotetext{2}{Sources in the updated on-line version of the \\citet{ragu05} catalog (http://xray.sai.msu.ru/$^\\sim$raguzova/BeXcat/)} \\tablenotetext{3}{Systems with known optical counterpart {\\em and} verified X-ray behaviour (from Table~\\ref{bexlist})} \\end{center} \\end{table*} %---------------------------------------------------------------------------- Until the advent of the INTEGRAL mission in 2002, the number of \\bex\\ was growing fast, while the number of SGXBs had stabilised. During the 80's and 90's the rate of new discoveries was approximately four \\bex\\ for one SGXB. The reason lies in the different origin of the accreted mass (stellar wind versus Be star's envelope), leading to persistent emission in SGXBs and transient emission in \\bex. Since SGXBs are persistent sources, new discoveries come from the improvement of the sensitivity of the X-ray detectors on board space missions. \\bex\\ benefited from the technical advances too but the discovery of new systems is also related to their triggering mechanism being activated. In addition, the evolutionary time scales involved imply that SGXBs are less numerous than \\bex. The accretion-powered phase is relatively short for OB supergiant systems (of the order of 10000 years). %------------------------------FIG 2----------------------------- \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.9\\linewidth]{./preig_fig02.eps} \\caption{$P_{\\rm orb}-P_{\\rm spin}$ diagram. The open triangles correspond to 2S 0114+65 and OAO 1657-41 and the open circle to SAX J2103.5+4545.} \\label{pspo} \\end{center} \\end{figure} %----------------------------------------------------------- The different types of HMXBs occupy well-defined positions in the spin period versus orbital period diagram \\citep{corb86}, which reflects the different types of mass transfer. SGXBs exhibit no correlation at all or an anticorrelation, while \\bex\\ show a positive correlation in this diagram. Figure~\\ref{pspo} displays an updated version of Corbet's diagram. Disc-fed SGXBs (squares) show short orbital periods and short spin periods and display an anticorrelation in the $P_{\\rm orb}-P_{\\rm spin}$ diagram. The small orbital separation and evolved companions make Roche lobe overflow the most likely mass transfer mechanism. Wind-fed SGXBs (triangles) show long spin periods and short orbital periods, occupying a more or less flat region in the $P_{\\rm orb}-P_{\\rm spin}$ diagram. Two systems (open triangles) prevent the region from being a horizontal line: OAO1657-41, which might be making a transition to the disc-fed SGXB and 2S 0114+65, for which the association of the $10^4$ s pulsations with the spin period of the neutron star remains controversial \\citep{koenigsberger06}. The spin and orbital periods of \\bex\\ (filled circles) exhibit a clear correlation. The open circle in Fig.~\\ref{pspo} represents SAX J2103.5+4545, a peculiar system whose X-ray properties are similar to those of wind-fed systems but whose optical/IR emission resemble that of \\bex\\ \\citep{reig10a}. The observed correlation in Be/X systems is explained in terms of the equilibrium period, defined as the period at which the outer edge of the magnetosphere rotates with the Keplerian velocity \\citep{davi73,stel86,wate89}. If the the neutron star (and hence the magnetosphere) rotates faster that the equilibrium period, then matter is spun away by the propeller mechanism; only when the spin period is larger than the equilibrium period can matter be accreted on to the neutron star surface. This results in angular momentum transfer to the neutron star, increasing its rotation velocity (decreasing the spin period). The equilibrium period depends mainly on the mass flux (or accretion rate) because it determines the size of the magnetosphere which is assumed to corotate with the neutron star. In turn, the accretion rate depends on the separation of the two components of the binary systems, hence on the orbital period. The compact object in all confirmed \\bex\\ (Table~\\ref{bexlist}) is a neutron star. In fact, many times the neutron star is taken as a defining property of \\bex. It is common to read in the literature that a \\bex\\ consists of a neutron star and a Be star. However, there seems to be no apparent mechanism that would prevent the formation of Be stars with black holes (BH) or white dwarfs (WD). Surprisingly, not a single \\bex\\ is known to host a black-hole companion in our Galaxy, whereas the interpretation of $\\gamma$ Cas as a Be+WD system still remains very controversial. Two ideas have been put forward to explain this apparent lack of Be/BH binaries. \\citet{zhan04a} extended the application of the viscous decretion disc model of \\citet{okaz01} to compact companions of arbitrary mass and showed that the most effective Be disc truncation would occur in relatively narrow systems. Using the population synthesis results of \\citet{pods03}, which state that binary black holes are most likely to be born in systems with narrow orbits ($P_{\\rm orb}\\simless 30$ days), the reason for the lack of these systems can be attributed to the difficulty to detecting them. Be/BH binaries are expected to be X-ray transients with very long quiescent states. In contrast, \\citet{belc09} showed that Be/BH binaries do not necessarily have narrow orbits. These authors argued that the predicted ratio of Be/NS binaries to Be/BH binaries ($F_{\\rm NS-to-BH}\\sim 30-50$) in our Galaxy, based on current population synthesis models and evolutionary scenarios, is consistent with the observations: there are 60 known Be/NS binaries, hence one would expect 0--2 Be/BH binaries, consistent with the observed galactic sample. Both \\citet{zhan04a} and \\citet{belc09} agree in that Be/BH binaries should exist in the Galaxy. \\subsection{Other types of high-mass X-ray binaries} The traditional picture of two classes of HMXBs, namely SGXB --- subdivided into low-luminosity or wind-fed systems and high-luminosity or disc-fed systems --- and \\bex\\ --- transient and persistent --- is giving way to a more complex situation where newly discovered systems may not fit in these categories. \\subsubsection{Low eccentricity \\bex} \\label{low-e} There is a group of so far five \\bex\\ (X Per, GS 0834--430, KS 1947+300, XTE J1543--568, and 2S 1553-542\\footnote{The optical counterparts of XTE J1543--568 and 2S 1553-542 are not known, hence they do not appear in Table~\\ref{bexlist}.}), characterised by $P_{\\rm orb} \\simmore$ 30 d and very low eccentricity ($e\\simless 0.2$). Their low eccentricity requires that the neutron star received a much lower kick velocity at birth than previously assumed by current evolutionary models \\citep{pfah02}. Most popular models for neutron star kicks involve a momentum impulse delivered around the time of the core collapse that produced the neutron star. These models assume that neutron stars are born with speeds in excess of 100-200 km s$^{-1}$. Such velocities imply a probability close to zero that the post-supernova eccentricity is less than 0.2. Hence, the discovery of these low-eccentricity \\bex\\ with wide orbits was unexpected. Tidal circularizacion is ruled out since this mechanism requires that the star almost fills its Roche lobe and that there be an effective mechanism for damping the tide. Tidal torques should have little effect on the orbit of a HMXB with $P_{\\rm orb}=10$ days, as long as the secondary is not too evolved and the eccentricity is not so large that the tidal interaction is enhanced dramatically at periastron (as it is the case of typical \\bex). \\citet{pfah02} developed a phenomenological model that simultaneously accounts for the long-period ($\\simmore$30 d), low-eccentricity ($\\simless$0.2) HMXBs and which does not violate any previous notions regarding the numbers and kinematics of other neutron star populations. They propose that a neutron star receives a relatively small kick ($\\simless$50 km s$^{-1}$) if the progenitor, i.e., the core of the initially more massive star in the binary, rotates rapidly. This condition may be met when the progenitor star experienced case B$_e$ or C$_e$ mass transfer in a binary system\\footnote{In case B, mass transfer occurs during the shell hydrogen-burning phase, but prior to central helium ignition, while case C evolution begins after helium has been depleted in the core. Cases B and C are naturally divided into an early case (B$_e$ or C$_e$), where the envelope of the primary is mostly radiative, and a late case (B$_l$ or C$_l$), where the primary has a deep convective envelope.}. If the hydrogen-exhausted core of an initially rapidly rotating massive star is exposed following case Be or Ce mass transfer in a binary, then the core is also likely to be a rapid rotator. \\subsubsection{Obscured sources \\& Supergiant fast X-ray transients} Since the launch of {\\it INTEGRAL} in October 2002, the situation is also changing among the SGXB group. INTEGRAL has unveiled a population of highly obscured HMXBs with supergiant companions and a new type of source displaying outbursts which are significantly shorter than typical for \\bex\\ and which are characterized by bright flares with a duration of a few hours and peak luminosities of $10^{36}-10^{37}$ erg s$^{-1}$. These new systems have been termed as Supergiant fast X-ray Transients \\citep[SFXTs,][]{smit06,negu06,walt07,negu08}. Both obscured HMXBs and SFXTs display X-ray and IR spectra typical of SGXBs. In some cases, the X-ray sources are pulsed and orbital parameters typical of persistent SGXBs have been found \\citep[e.g.,][]{boda06,zuri06}. The heavily-absorbed sources had not been detected by previous missions due to high absorption, which renders their spectra very hard. The current understanding is that the entire binary system is surrounded by a dense and absorbing circumstellar material envelope or cocoon, made of cold gas and/or dust \\citep{chaty08}. SFXTs differ from SGXBs because they are only detected sporadically, during very brief outbursts \\citep{romano10}. A promising model to explain SFXTs invokes highly structured (clumpy) stellar winds. The outburst occurs as a result of the accretion of one of the clumps of dense matter from the wind. An alternative model (\"Be-type\" model) assumes a very elliptical orbit for the binary. In this model outbursts are triggered when the compact object travels through its periastron. Other possibilities imply that SFXTs contain strongly magnetised neutron stars. The outbursts result from the overcoming of centrifugal and magnetic barriers \\citep[see][and references therein]{grebenev10}. %---------------------------FIG 3-------------------------------- \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.9\\linewidth]{./preig_fig03.ps} \\caption{Spectral energy distribution of LSI +61 303. From \\citet{chernyakova06}}. \\label{grb_sed} \\end{center} \\end{figure} %----------------------------------------------------------- \\subsubsection{$\\gamma$-ray binaries} \\label{grb} $\\gamma$-ray binaries are HMXBs that emit most of their radiative output in the MeV-TeV range. Currently, only four HMXBs are well-established members of this group of high-energy sources: LS I +61 303, LS 5039, PSR B1259-63, and Cygnus X-3, while other two are firm candidates: Cygnus X-1, and HESS J0632+057\\footnote{The Be star HD 215227, likely counterpart of the gamma-ray source AGL J2241+4454, has been suggested as a new candidate \\citep{williams10}.}. The optical counterpart is either a luminous O-type star (Cygnus X-1, LS 5039), a Be star (LS I +61 303, PSR B1259-63, HESS J0632+057) or a likely Wolf-Rayet star in Cygnus X-3. In addition to the variable $\\gamma$-ray emission these systems share in common a resolved radio counterpart with a jet or jet-like structure, multi-wavelength orbital modulation and spectral energy distribution (Fig.~\\ref{grb_sed}). The wide range of orbital parameters \\citep{paredes08} and the non-unique nature of the compact companion (unknown in most systems but with a confirmed neutron star in PSR B1259-63 and a fairly secure black hole in Cyg X-1) represent a challenge for the theoretical modelling of these systems. Two alternative scenarios may explain the variable $\\gamma$-ray emission: the microquasar or accretion-powered scenario and the pulsar wind scenario. All confirmed $\\gamma$-ray binaries show a jet or jet-like radio structure, which would indicate the presence of relativistic particles. In black-hole binaries the radio jet can account for their broad-band spectrum, from radio to X-rays \\citep{markoff03} as well as for the origin of most of the timing variability \\citep{kylafis08}. Therefore, it is reasonable to think that the jet may also be the origin of the very high-energy $\\gamma$-rays. In the context of relativistic jets, the most efficient gamma-ray mechanism would be inverse Compton scattering, by which relativistic particles collide with stellar and/or synchrotron photons and boost their energies to the VHE range. Two flavors of the microquasar model can be found in the literature depending on whether hadronic or leptonic jet matter dominates the emission at such an energy range. Among leptonic jet models, there are inverse Compton jet emission models in which X-rays and $\\gamma$-rays result from synchrotron self-Compton processes \\citep{atoyan99}, or in which the seed photons come from external sources, i.e., companion star and/or accretion disk photons \\citep{kaufman02,georganopoulos02}. In the hadronic scenario, the gamma-ray emission arises from the decay of neutral pions created in the inelastic collisions between relativistic protons in the jet and either the ions in the stellar wind of the massive companion star \\citep{romero03} or nearby high-density regions (i.e. molecular clouds) \\citep{bosch05}. Alternatively, relativistic particles can be injected in the surrounding medium by the wind from a young pulsar. In the pulsar wind scenario, the rotation of a young pulsar provides stable energy to the nonthermal relativistic particles in the shocked pulsar wind material outflowing from the binary companion. As in the microquasar-jet models, the $\\gamma$-ray emission can be produced by inverse Compton scattering of the relativistic particles from the pulsar wind on stellar photons \\citep[][and references therein]{torres10}. In the pulsar wind scenario the resolved radio emission is not due to a relativistic jet akin to those of microquasars, but arises instead from shocked pulsar wind material outflowing from the binary \\citep{dubus06}. %---------------------------FIG 4-------------------------------- \\begin{figure*}[t] \\begin{center} \\begin{tabular}{cc} \\includegraphics[width=0.4\\linewidth]{./preig_fig04a.eps} & \\includegraphics[width=0.4\\linewidth]{./preig_fig04b.eps} \\\\ \\end{tabular} \\caption{IR excess and H$\\alpha$ emission are the two main observational characteristics of Be stars.} \\label{be} \\end{center} \\end{figure*} %----------------------------------------------------------- \\subsubsection{$\\gamma$-Cas like objects} A growing number of early Be stars discovered in X-ray surveys exhibit strong X/optical flux correlations and X-ray luminosities intermediate between those of normal stars and those of most \\bex\\ in quiescence \\citep{lopes06}. The optical properties are very similar to those of \\bex: {\\em i)} the spectral type is always in the range B0-B1 III-Ve, and {\\em ii)} show \\ha\\ equivalent widths stronger than $-20$ \\AA. However, they differ from the typical \\bex\\ in their X-ray properties: {\\em i)} they show harder X-ray spectra that can be best fitted by a thin thermal plasma with $T\\sim 10^8$ K, rather than a power law as seen typically in HMXBs, {\\em ii)} there is no evidence for coherent pulsations in any of these systems but strong variability on time scales as short as 100 s is usually observed, and {\\em iii)} they do not exhibit large X-ray outbursts. The prototype of this group of sources is $\\gamma$ Cas. Two models have been put forward to explain this type of system: accretion onto a compact object (most likely a white dwarf) and magnetically heated material between the photosphere of the B star and the inner part of its disc \\citep{robinson02}. In the magnetic corona model the hard X-rays result from high-energy particles that are emitted due to magnetic reconnection, while the optical variability is due to changes in the density structure of the inner disc as a consequence of turbulence generated by changes in the magnetic field. In support of this scenario there is the fact that X-ray fluxes show random variations with orbital phase, thereby contradicting the binary accretion model, which predicts a substantial modulation. In favour of the binary model is the similarity of the X-ray spectra with those of cataclysmic variables, the rapid variability and large numbers of Be + white dwarf systems predicted by the theory of evolution of massive binaries. In the remaining part of this report I shall concentrate on the properties and variability of \\bex\\. Only when for the sake of the discussion a comparison with the behaviour of other type of X-ray binaries may be illustrative, will these other type of binaries be mentioned. ", "conclusions": "The traditional picture of two classes of high-mass X-ray binaries, namely supergiant X-ray binaries and Be/X-ray binaries is giving way to a more complex situation where newly discovered systems may not fit in these categories. Superfast X-ray transients, low eccentricity \\bex, $\\gamma$-ray binaries, and $\\gamma$-Cas like objects are among the new type of systems that only recently have begun to emerge. \\bex\\ do not form a homogeneous group either with transient and persistent sources, highly eccentric and nearly circular orbits, and fast (P$_{spin}$ $\\sim$ few seconds) and slow (P$_{spin}$ $\\sim$ few hundred of seconds) rotating neutron stars. Virtually, all \\bex\\ are X-ray pulsars. The study of the variability of the X-ray spectral and pulse timing parameters across giant outbursts is crucial to understand the physics of accretion in strongly magnetised neutron stars. Only recently, attention is being paid to the aperiodic variability and evolution of broad-band noise components across the outburst. The application of diagrams such as hardness-intensity and colour-colour diagrams, so widely used in black-hole and low-mass binaries, to \\bex\\ and the definition of spectral states is a promising way to understand the phenomenology of the evolution of outbursts in these systems. Progress in understanding \\bex\\ is also being made in the optical and infrared bands where the massive companion shines bright. The data gathered over the past ten years indicate that the interaction between the compact object and the Be type star works in two directions: the massive companion provides the source of matter for accretion; the amount of matter captured and the way it is captured (transfer mechanism) change the physical conditions of the neutron star (e.g. by spinning it up or down and shrinking or expanding the magnetosphere). But also, the continuous revolution of the neutron star around the optical counterpart produces observable effects, the most important of which is the truncation of the Be star's equatorial disc. Future missions with improved sensitivities in the X-ray band will help solve remaining open questions such as the formation of accretion discs during X-ray outbursts, the origin of quasi-periodic oscillations, characterisation of the population of persistent \\bex\\ and the triggering mechanism of the giant outbursts." }, "1101/1101.4933_arXiv.txt": { "abstract": "In order to investigate the evolution of E+A galaxies, we observed a galaxy SDSS J160241.00+521426.9, a possible E+A progenitor which shows both emission and strong Balmer absorptions, and its neighbor galaxy. We used the integral field spectroscopic mode of the Kyoto Tridimensional Spectrograph (Kyoto3DII), mounted on the University of Hawaii 88-inch telescope located on Mauna Kea, and the slit-spectroscopic mode of the Faint Object Camera and Spectrograph (FOCAS) on the Subaru Telescope. We found a strong Balmer absorption region in the center of the galaxy and an emission-line region located 2 kpc from the center, in the direction of its neighbor galaxy. The recession velocities of the galaxy and its neighbor galaxy differ only by 100 km s$^{-1}$, which suggests that they are a physical pair and would have been interacting. Comparing observed Lick indices of Balmer lines and color indices with those predicted from stellar population synthesis models, we find that a suddenly quenched star-formation scenario is plausible for the star-formation history of the central region. We consider that star formation started in the galaxy due to galaxy interactions and was quenched in the central region, whereas star formation in a region offset from the center still continues or has begun recently. This work is the first study of a possible E+A progenitor using spatially resolved spectroscopy. ", "introduction": "E+A galaxies are understood as post-starburst galaxies due to the presence of strong Balmer absorption lines (H$\\beta$, H$\\gamma$, H$\\delta$) and the lack of emission lines \\citep{Poggianti:1999, Goto:2005}. Examinations of galaxy morphologies \\citep{Yang:2004, Yang:2008} and companion galaxies statistics \\citep{Goto:2005, Yamauchi:2008} led to the conclusion that local field E+A galaxies are mainly driven by mergers and interactions. Integral field spectrograph (IFS) observations for E+A galaxies indicate disturbed morphologies and significant rotation, which supports the theory that they are produced by gas-rich galaxy mergers and interactions \\citep{Pracy:2009}. Recently, \\citet{Rich:2010} has reported that a starburst galaxy NGC 839 has both A-type stellar population and a galactic wind, and has discussed a link between lower-mass starburst systems and E+A galaxies. An E+A galaxy is one phase in galaxy evolution, and thus, it is important to understand the evolution of E+A galaxies. All these results, however, are based on information obtained from galaxies already in the $post$-starburst phase. In order to investigate the evolution of E+A galaxies, it is necessary to examine pre-E+A galaxies, or E+A progenitors, which would have both starburst and post-starburst regions, and this requires spatially resolved spectroscopic observations. In this work, we observed a possible progenitor of an E+A galaxy, SDSS J160241.00+521426.9 (J1602), using the IFS mode of the Kyoto Tridimensional Spectrograph \\citep[Kyoto3DII:][]{Sugai:2010}, mounted on the University of Hawaii 88-inch telescope on Mauna Kea. The galaxy has an apparent companion galaxy, seen in the Sloan Digital Sky Survey (SDSS) image. The 3$\\arcsec$-fiber spectra from the SDSS show an interesting combination of very strong Balmer absorption lines and emission lines. With the IFS observation on the UH 88-inch telescope as well as slit-spectroscopic observations using the Subaru Telescope Faint Object Camera and Spectrograph \\citep[FOCAS: ][]{Kashikawa:2002}, we investigated the spatial distributions of absorption and emission line regions. Identifying the locations of starburst regions in E+A galaxies provides constrains on E+A galaxy formation theories. The presence of both post-starburst regions and current starburst regions in one galaxy may provide us with an unprecedented opportunity to examine the progenitor of an E+A galaxy. Through detailed analysis, we attempted to understand the underlying physical process that caused the simultaneous presence of post-starburst and current starburst regions in this galaxy. ", "conclusions": "In order to investigate the evolution of E+A galaxies, we observed SDSS J1602 and its neighbor galaxy with Kyoto3DII and FOCAS. These are the first spatially resolved spectroscopic observations of E+A progenitors. We found a post-starburst region in the center of J1602 (PS1) and a starburst region located 2 kpc from the center (SB1), in the direction of the neighbor galaxy. The fact that this galaxy has both starburst and post-starburst regions indicates that it is in a critical phase of evolution. The recession velocities of J1602 and its neighbor galaxy differ only by $\\sim$ 100 km s$^{-1}$. Thus, they are a physical pair and are considered to have been interacting, probably in a retrograde encounter. A local velocity field of 80 km s$^{-1}$ is detected in SB1. Comparing the observed Lick indices of the H$\\gamma_A$ and H$\\delta_A$ and color indices to those predicted from stellar population synthesis models, we find that a suddenly quenched star-formation scenario is plausible for the star-formation history of PS1. From our results, we suggest a history for J1602. Starbursts occurred within 2 kpc from the center of J1602 probably due to galaxy interaction, then about 200 Myr ago starbursts were suddenly quenched, whereas the starburst in SB1 may have continued or have occurred recently. The SB1 starburst will eventually stop, and all of J1602 will become a post-starburst. Follow-up observation, for example molecular gas mapping of this system, will further elucidate the evolution process of E+A galaxies." }, "1101/1101.3664_arXiv.txt": { "abstract": "{} { We present $VRi$ photometric observations of the quadruply imaged quasar \\he, carried out with the Danish $1.54\\meter$ telescope at the La Silla Observatory. Our aim was to monitor and study the magnitudes and colors of each lensed component as a function of time. } { We monitored the object during two seasons (2008 and 2009) in the $VRi$ spectral bands, and reduced the data with two independent techniques: difference imaging and PSF (Point Spread Function) fitting. } { Between these two seasons, our results show an evident decrease in flux by $\\approx$ 0.2--0.4 magnitudes of the four lensed components in the three filters. We also found a significant increase ($\\approx$ 0.05--0.015) in their $V-R$ and $R-i$ color indices. } { These flux and color variations are very likely caused by intrinsic variations of the quasar between the observed epochs. Microlensing effects probably also affect the brightest ``A'' lensed component.} ", "introduction": "In the framework of the MiNDSTEp (Microlensing Network for the Detection of Small Terrestrial Exoplanets) campaign \\citep{dominik10}, which has as a main target the systematic observation of bulge microlenses, we developed a parallel project concerning photometric multi-band observations of several lensed quasars\\footnote{\\he, UM673/\\object{Q0142--100}, \\object{Q2237+0305}, \\object{WFI2033---4723} and \\object{HE0047--1756}}. In the present paper we focus on \\he (see~Fig.~\\ref{fig:mediumfield}), a QSO discovered by \\cite{wisotzki00} in the course of the Hamburg/ESO digital objective prism survey, and confirmed to be a quadruply imaged quasar by \\cite{wisotzki02}. The lensing galaxy was initially identified as an elliptical with a scale length of $\\approx12\\kilo\\rm pc$ at a redshift in the range $z=0.3$--$0.4$. The time delays between the four images (labeled ``A'', ``B'', ``C'', ``D'', starting from the brighter one and proceeding clockwise) of the quasar were estimated around 10 days, and the quasar itself showed some signs of intrinsic variability \\citep{wisotzki02}. \\noindent More recently, the value of the redshift for the lensing galaxy was estimated as $z=0.44\\pm 0.20$, and the quasar redshift was confirmed to be $z=1.6895\\pm 0.0005$, with a $\\Delta z$ between the components of $\\approx 0.0015$ rms \\citep{wisotzki03}. These spectrophotomeric observations showed some possible signature of microlensing effects in the continuum and in the spectral emission lines for the ``D'' component. \\begin{figure}[t] \\centering \\includegraphics[width=8.7cm]{16188fg1.eps} \\caption{Zoom of a DFOSC $i$ filter image showing the four components of the lensed quasar and four nearby stars. The components are labeled following the notation of \\cite{wisotzki02}: ``A'' for the brighter component, ``B'', ``C'' and``D'' clockwise. The stars ``R'', ``S'', ``T'' and ``U'' were used to search for a suitable reference star. The ``R'' star was finally chosen. The contrast of the displayed image, on a negative scale, was selected to improve the visibility of the lensed components. The image is a median of the three CCD frames collected on 2008 August 8.} \\label{fig:mediumfield} \\end{figure} \\begin{figure}[t] \\centering \\includegraphics[width=8.7cm]{16188fg2.eps} \\caption{Zoom of a DFOSC $i$ filter image showing the galaxy environment near \\he. The objects are labeled following the notation of \\cite{morgan05}. The contrast, on a negative scale, was selected to improve the visibility of the galaxies. The image is a median of the three CCD frames collected on 2008 August 8. } \\label{fig:bigfield} \\end{figure} \\noindent \\cite{morgan05} provided milliarcsecond astrometry, revised the value of the lens redshift at $z=0.4546\\pm 0.0002$ with the Low-Dispersion Survey Spectrograph 2 (LDSS2) on the Clay $6.5\\meter$ telescope, and studied the galaxy environment of the lens, because it is located in a dense galaxy field. The results do not show any evidence of a cluster for the considered galaxies. However, the nearest galaxies (G20, G21, G22, G23, and G24 in Fig.~\\ref{fig:bigfield}) whose redshifts were not measured, left this scenario open. Nevertheless, the results of a deep investigation concerning the direction of an external shear in the gravitational field of the lens do not show an evident correlation with the position of the near galaxies. As a remaining explanation, \\cite{morgan05} suggested the presence of substructures in the lensing galaxy. \\noindent The first systematic monitoring, which was performed in the $R$ filter and covered the years between 2003 and 2005, was carried out by \\cite{kochanek06}: that paper provided astrometric measurements compatible with the previous works, measured the time delays between the images ($\\Delta t_{AD}= -14.37$, $\\Delta t_{AB}= -8.00$, and $\\Delta t_{AC}= -2.10$ days, with errors respectively of $6\\%$, $10\\%$ and $35\\%$), and finally confirmed the lensing galaxy as an elliptical with a rising rotation curve. \\noindent Furthermore, \\cite{mediavilla09} observed \\he in the framework of a monitoring of 29 lensed quasars, and attributed eventual microlensing events to the normal stellar populations, while \\cite{blackburne10a} focused on the quasar itself, applying a model with a time-variable accretion disk to the object. \\cite{mosquera10} found clear evidence of chromatic microlensing in the ``A'' component, and provided an estimate of the disk size in the $R$ band in agreement with the simple thin-disk model. \\cite{blackburne10b} used the chromatic microlensing to model the accretion disk, and \\cite{courbin10} recalculated the time delays with N-body realizations of the lensing galaxy, which he thought to belong to the ``B component'' ($\\Delta t_{BA}= 8.4$, $\\Delta t_{BC} =7.8$ and $\\Delta t_{BD}= \u22126.5$ days with errors of $25\\%$, $10\\%$, and $11\\%$ respectively). Considering multi-color observations of other lensed quasars, a single-epoch multi-band photometry was used on \\object{MG0414+0534} to constrain the accretion disk model and the size of the emission region in the continuum \\citep{bate08a,bate08b,floyd08}. \\noindent A multi-epoch multi-band photometry, carried out during several years, was used for the quasar \\object{Q2237+0305} by \\cite{kopletova06}, who observed the object during five years (1995--2000) in the $VRI$ bands. \\cite{anguita08} combined these data with OGLE observations. \\cite{mosquera09} monitored the object in eight filters and found evidence for microlensing in the continuum, but not in the emission lines. \\noindent Furthermore, \\object{Q2237+0305} was the object of deep studies focused on the lens galaxy \\citep{poindexter10a}, and on the inclination of the accretion disk \\citep{poindexter10b}. Another example of multi-epoch multi-band observations is given by UM673/\\object{Q0142--100}, observed in the Gunn $i$ and Cousins $V$ filters between 1998 and 1999 \\citep{nakos05} and in the $VRI$ bands between 2003 and 2005 \\citep{kopletova09}. \\noindent Unlike for these objects, no \\emph{systematic} multi-band photometry has ever been carried out for \\he. Here, we present two periods of multi-band photometric observations of \\he, performed in the $VRi$ spectral bands with the Danish $1.54\\meter$ telescope at the La Silla Observatory. \\noindent In Sect.~\\ref{sec:obs} we explain how the observations were carried out; in Sect.~\\ref{sec:red} we focus on the data reduction, and we describe the two independent techniques: difference imaging and PSF (Point Spread Function) fitting, that were used to construct the light curves. In Sect.~\\ref{sec:res} we present the results. Finally, in Sect.~\\ref{sec:conc} we summarize the conclusions. ", "conclusions": "\\label{sec:conc} Systematic multi-spectral band photometry of the quadruply imaged quasar \\he, carried out during two seasons (2008 and 2009), shows a significant decrease in flux of the four lensed components between the two epochs. \\noindent The drop in flux observed for the four lensed components between 2008 and 2009 is very likely caused by a change in the intrinsic luminosity of the quasar. This corroborates the previous studies of the \\he gravitational lens system by \\cite{kochanek06} and \\cite{courbin10}. \\noindent Concerning the color variations, the intrinsic reddening of a quasar when it becomes fainter in luminosity is an effect already observed in previous studies \\citep{pereyra06}, and the same trend probably accounts for the similar observed changes in the colors of the four lensed components of \\he. \\noindent This hypothesis is enforced if we suppose that the intrinsic photometric quasar variations in the different colors are not synchronized, which provides an explanation for the differences in scatter between the epochs (see Fig.~\\ref{fig:shifted}). Microlensing probably provides the additional effect necessary to account for the higher flux variation observed for the ``A'' component. \\noindent The presented observations also show that a well sampled multi-band photometry can help in distinguishing the nature of the variability of multiply imaged objects, in particular gravitationally lensed quasars. We suggest to couple this technique in the future with integral field spectroscopy to provide an additional way for an even more detailed investigation of the observed phenomena." }, "1101/1101.4875_arXiv.txt": { "abstract": "We show that the r-mode instability can generate strong toroidal fields in the core of accreting millisecond quark stars by inducing differential rotation. We follow the spin frequency evolution on a long time scale taking into account the magnetic damping rate in the evolution equations of r-modes. The maximum spin frequency of the star is only marginally smaller than in the absence of the magnetic field. The late-time evolution of the stars which enter the r-mode instability region is instead rather different if the generated magnetic fields are taken into account: they leave the millisecond pulsar region and they become radio pulsars. ", "introduction": "Since the first paper in which r-modes in rotating neutron stars were shown to be unstable with respect to the emission of gravitational waves \\cite{Andersson:1997xt}, it was recognized their relevance in Astrophysics to explain the observed distribution of the rotation frequency of stars in Low-Mass-X-Ray-Binaries (LMXBs) (see \\cite{Andersson:2000mf} for a review). On the other hand, the instability triggered by r-modes is actually damped to some extent by the viscosity of the matter composing the star: the higher the viscosity, the higher is the frequency of rotation of the star. This fact opens the possibility of investigating the internal composition of compact stars by studying the viscosity of the different possible high density phases which can appear in these stellar objects. A number of papers are currently present in the literature about the shear and bulk viscosities of nucleonic matter \\cite{Sawyer:1989dp,Haensel:1992zz,Haensel:2000vz,Haensel:2001mw,Benhar:2007yj}, hyperonic matter\\cite{Lindblom:2001hd,Haensel:2001em,Chatterjee:2006hy,Chatterjee:2007iw,Gusakov:2008hv,Sinha:2008wb,Jha:2010an}, kaon condensed matter \\cite{Chatterjee:2007qs,Chatterjee:2007ka}, pure quark phases \\cite{Madsen:1992sx,Madsen:1999ci,Alford:2006gy,Sa'd:2006qv,Sa'd:2007ud,Alford:2007rw,Blaschke:2007bv,Sa'd:2008gf,Mannarelli:2008je,Jaikumar:2008kh,Alford:2009jm}, and mixed phases \\cite{Drago:2003wg}. Interestingly, by studying the so called ``window of instability'' of the r-modes, which is determined by the shear and the bulk viscosity of the matter, it was pointed out in \\cite{Drago:2007iy} that a detection of a sub-millisecond rotating star would indicate the existence of a very viscous phase in the star with only quark or hybrid stars as possible candidates. Also the temporal evolution of the spin frequency of a star under the effect of r-modes instability has been studied for different types of composition and different physical systems: newly born compact stars and old stars in binaries \\cite{Lindblom:1998wf,Owen:1998xg,Andersson:2001ev,Wagoner:2002vr,Drago:2004nx,Drago:2007iy}. While neutron stars could be very powerful gravitational waves emitter only during their first months after birth \\cite{Owen:1998xg}, hyperonic and quark or hybrid stars could turn into steady gravitational waves sources if present in LMXBs \\cite{Andersson:2001ev,Wagoner:2002vr,Reisenegger:2003cq}. R-modes are also responsible for differential rotation in the star which in turn generates a toroidal magnetic field: besides the viscosity of matter, the production of this magnetic field represents a very efficient damping mechanism for the r-modes. This effect has been proposed and investigated in \\cite{Rezzolla2000ApJ,Rezzolla:2001di,Rezzolla:2001dh} for the case of a newly born neutron star and it has been recently included within the r-mode equations of the neutron stars in LMXB \\cite{Cuofano:2009xy,Cuofano:2009yg}: by calculating the back-reaction of the magnetic field on the r-modes instability, it has been proven that magnetic fields of the order of $10^{15}$ G can be produced. Remarkably, this mechanism could be at the origin of the enormous magnetic field of magnetars. In this paper we extend the calculations of \\cite{Cuofano:2009xy,Cuofano:2009yg} to the case of quark stars in LMXB. The motivation for this investigation is that the evolution of accreting stars and their internal magnetic field depends strongly on the r-modes instability window which, for quark stars, is qualitatively different from the one of neutron star. Indeed, due to the large contribution to the bulk viscosity of the non leptonic weak decays occurring in strange quark matter, the window of instability splits into two windows, a small one at large temperatures (which is actually irrelevant for the evolution) and a big one at low temperatures \\cite{Madsen:1999ci,Drago:2007iy}. Moreover, quark stars do not have a crust (or just a very thin crust) and, as we will discuss, this prevents the trapping of the internal magnetic field developed during the evolution with possible observable signatures. The paper is organized as follows: in Sec.~II we introduce the system of equations which provide the temporal evolution of a compact star in a LMXB. In Sec.~III we show the results of our numerical calculations and finally in Sec.~IV we present of conclusions. ", "conclusions": "We have shown that strong toroidal fields can be generated in the core of an accreting millisecond quark star which enters the r-mode instability window. Tayler instability sets in when the generated toroidal fields exceed the critical value $B_{tor}^{cr}\\sim 10^{12}$~G and a new poloidal component of similar strength is produced. Our results show that the maximum spin frequency for quark stars does not change significantly when taking into account the internal generated magnetic fields. \\\\ The scenario after the development of the Tayler instability depends on the presence and the properties of a possible crust. If the crust is not present, the generated large poloidal component diffuses quickly outside the core and prevents the further accretion of mass on the star. On the other hand, if a highly conductive crust is present, it could screen to some extent the internal magnetic field. However, taking into account the ambipolar diffusion, which is, in this case, the dominant dissipation mechanism, the star could expel the internal magnetic field in a few millions years, which would then stop the accretion. In both cases the quark star evolves into the region of radio pulsars, as shown in Fig.\\ref{evolution}: this represents a new possible scenario for the formation of radio pulsars." }, "1101/1101.4044_arXiv.txt": { "abstract": "The young high-eccentricity binary DQ~Tau exhibits powerful recurring millimeter-band (mm) flaring attributed to collisions between the two stellar magnetospheres near periastron, when the stars are separated by only $\\sim 8$~R$_{\\star}$. These magnetospheric interactions are expected to have scales and magnetic field strengths comparable to those of large X-ray flares from single pre-main-sequence (PMS) stars observed in the Chandra Orion Ultradeep Project (COUP). To search for X-rays arising from processes associated with colliding magnetospheres, we performed simultaneous X-ray and mm observations of DQ~Tau near periastron phase. We report here several results. 1) As anticipated, DQ~Tau was caught in a flare state in both mm and X-rays. A single long X-ray flare spanned the entire 16.5 hour $Chandra$ exposure. 2) The inferred morphology, duration, and plasma temperature of the X-ray flare are typical of those of large flares from COUP stars. 3) However, our study provides three lines of evidence that this X-ray flare likely arises from colliding magnetospheres: the chance of capturing a large COUP-like flare within the span of our observation is small; the relative timing of the X-ray and mm flares indicates the Neupert effect and is consistent with a common coronal structure; the size of the emitting coronal structure ($4-5$~R$_{\\star}$) inferred from our analysis (which is admittedly model-dependent and should be considered with caution) is comparable to half the binary separation. 4) The peak flare X-ray luminosity is in agreement with an estimate of the power dissipated by magnetic reconnection within the framework of a simple model of interacting magnetospheres. ", "introduction": "} T-Tauri stars generally show highly elevated levels of X-ray activity arising mostly from violent magnetic reconnection events \\citep[e.g.][]{Feigelson99}. This strong X-ray emission has far-reaching implications for the physical processes in the circumstellar environment, the formation of planetary systems, and the evolution of protoplanetary atmospheres \\citep[e.g.][]{Glassgold05,Feigelson09}. Recent X-ray surveys of nearby pre-main-sequence (PMS) stellar populations give detailed insights into T-Tauri magnetic flaring; these include the Chandra Orion Ultradeep Project \\citep[COUP;][]{Getman05} and the XMM-Newton Extended Survey of Taurus \\citep[XEST;][]{Gudel07}. Astrophysical studies of the properties of individual flares and statistical studies of many flares, from both the COUP and XEST observations, reveal that most events are similar to solar magnetic flaring, but with X-ray luminosities up to $10^3-10^5$ times higher than seen in the Sun and plasma temperatures up to $50$ times higher \\citep[e.g.][]{Favata05, Wolk05, Flaccomio05, Stassun06, Maggio07, Caramazza07, Franciosini07}. Accretion shocks may contribute only a small fraction to the total X-ray emission from T-Tauri stars, in the form of soft X-ray excess emission \\citep[e.g.][]{Telleschi07b, Gudel07b}. COUP, the deepest and longest (13.2 continuous days) X-ray observation of a young stellar cluster, provided a unique opportunity to study relatively rare (typically $1$ flare per week per star) big X-ray flares from T-Tauri stars \\citep{Favata05}. More recently, a detailed and systematic study of $>200$ big flares from $>150$ bright PMS stars detected in the COUP observation by \\citet{Getman08a,Getman08b} [hereafter G08a and G08b] shows that they are the most powerful, longest, and hottest stellar flares corresponding to the largest known stellar X-ray coronal structures, reaching up to several stellar radii in both disk-bearing (Class~II) and diskless (Class~III) systems. The associated large-scale magnetic fields (assuming a dipolar geometry) have an equipartition strength $B = 0.05-0.3$~kG in the outer loop region, consistent with optical Zeeman measurements of surface field strengths around $2-3$~kG in magnetically active T-Tauri stars \\citep[e.g.][]{Johns-Krull07,Donati08}. G08ab also provide observational evidence for magnetospheric truncation by a disk in a Class~II system, and for the ability of X-ray loops to withstand centrifugal forces in rapidly rotating Class~III systems. G08ab propose the COUP sample of flares as possible enhanced analogues of very rare Solar Long Decay events (LDEs) associated with X-ray arches and streamers. While observations of PMS stars at millimeter wavelengths are typically used to study steady thermal emission from dust in their protoplanetary disks, transient gyrosynchrotron and synchrotron continua from flares with spectral peaks in the GHz-THz range can also be seen \\citep{Priest02, Kaufmann86}. Long-term radio variability in older Class~III PMS stars has been known for some time \\citep{Garay87} and generally does not show correlations with X-ray variability \\citep[e.g.][]{Forbrich07}. Short-term (hours- or day-long) radio outbursts are occasionally seen from PMS systems, for example: a remarkably powerful millimeter flare from a heavily absorbed Class~III system GMR-A \\citep{Bower03}; recurring flares from the Class~III binary V773 Tau A \\citep{Massi08} and the Class~II binary system DQ~Tau \\citep{Salter08,Salter10}; IRS~5b and IRS~7A Class~I protostars in the Corona Australis cloud \\citep{Choi09}; and a poorly characterized system in Orion's embedded BN/KL star forming region \\citep{Forbrich08}. Prior to the current study, only one of these cases, GMR-A, was simultaneously observed with a modern X-ray telescope, by a coincidence with the {\\it Chandra}-COUP observation. The GMR-A mm flare was associated, though not exactly simultaneous, with several days of complex big X-ray flares \\citep{Bower03, Furuya03, Favata05}. The GMR-A star is believed to be a single star with a strong coronal magnetic field, while V773 Tau A and DQ Tau are close binary (or triple) systems with component separations at periastron of $30$~R$_{\\star}$ and $8$~R$_{\\star}$, respectively. For these multiple systems the magnetic reconnection events have been attributed to interacting magnetospheres. The DQ~Tau binary is especially useful in regards to possible X-ray emission. This is a non-eclipsing, double-lined spectroscopic binary, comprised of two relatively equal-mass (equal-radius) PMS stars of $M \\sim 0.65$~M$_{\\odot}$ ($R \\sim 1.6$~R$_{\\odot}$) with spectral types in the range of K7 to M1, a rotational period of $P \\sim 3$~days for both stars, and a robust orbital period of 15.804 days \\citep{Mathieu97,Basri97}. Its highly eccentric orbit ($e=0.556$) exhibits a periastron separation of only $\\sim 8$~R$_{\\star}$ ($\\sim 13$~R$_{\\odot}$). The spectral energy distribution (SED) of DQ~Tau is fairly typical of a Class~II system, fit by a large circumbinary disk of about $0.002 - 0.02$~M$_{\\odot}$\\citep{Mathieu97}. For more than 65\\% of periastron encounters, the system experiences optical brightenings as a result of variable and irregular accretion \\citep{Mathieu97,Basri97}. The periastron separation is expected to induce magnetospheric interactions at scales and magnetic field strengths comparable to those inferred for the COUP sample of big flares (G08b). We thus performed simultaneous X-ray and mm observations of the orbital segment of DQ~Tau around the peaks of the previously detected mm flares \\citep{Salter08} to search for X-ray emission arising from processes associated with colliding magnetospheres. We report here a $Chandra$ detection of a long X-ray flare accompanied by mm activity close to the periastron passage of DQ~Tau. The mm observations are discussed in more detail in \\citet{Salter10}. The $Chandra$ data are described in \\S \\ref{sec_chandra_reduction}; and the treatment of mild photon pile-up in the observation is given in \\S \\ref{sec_pileup_analysis}. Archived {\\it XMM-Newton} X-ray observations of DQ~Tau at an orbital phase away from periastron are presented in \\S \\ref{sec_xmm_reduction}. Time-integrated $Chandra$ and $XMM-Newton$ spectra are compared in \\S \\ref{sec_chandravsxmm}. The $Chandra$ flare analyses and the derived flare loop length and loop thickness (within the framework of a single-loop model) are presented in \\S \\ref{sec_trs} and \\S \\ref{sec_single_loop_modeling}. A comparison of the $Chandra$ flare with the coincident mm flare is provided in \\S \\ref{sec_chandra_iram}. A comparison with the COUP sample of big flares is given in \\S \\ref{sec_coup_star_comparison}. We end in \\S \\ref{sec_discussion_conclusions} with a discussion of the applicability of the single-loop approach (including information from Appendix ~\\ref{sec_solar_flares} on an X-ray analysis of the multi-loop solar X-class flares), our observational findings, and their implications for the origin of the X-ray emission, energetics, and loop geometry. ", "conclusions": "} \\subsection{Single Loop versus Multiple Loop Scenarios \\label{sec_single_vs_multiple}} Large flares on the Sun often involve arcades of dozens or hundreds of sequentially reconnected magnetic loops, which are often observed to have different temperatures with the cooler loops lying below the hotter ones \\citep[e.g.][and references therein]{Reeves02}. Despite the morphological complexity of such flares, their X-ray light curves often have simple exponential rise and decay shapes.\\footnote{Images of complex solar flaring loop arcades spatially resolved by {\\it TRACE} \\citep[][]{Handy99} at Extreme Ultraviolet (EUV) wavelengths and their associated X-ray light curves obtained by the Geostationary Operational Environmental Satellites \\citep[{\\it GOES},][and references therein]{Garcia94} can be found and inspected at the {\\it TRACE} flare catalog web site {\\url{http://hea-www.harvard.edu/trace/flare\\_catalog/index.html}} (see also Figure~\\ref{fig_solar_flares}).} In view of this solar analogy, the concept of the single-loop approach and its application to powerful spatially unresolved stellar flares might be questionable. R97 have applied their single-loop method to flare decays of 20 solar M- and C-class flares monitored by the Soft X-ray Telescope \\citep[SXT,][]{Tsuneta91} on board the {\\it Yohkoh} solar observatory satellite \\citep{Ogawara91}. Loop sizes derived using the method were shown to agree with the length of X-ray structures measured from direct inspection of SXT images (Figure~6 in R97). In Appendix~\\ref{sec_solar_flares} we apply the approach of R97 to 5 solar X-class flares that are clearly associated with arcades of multiple loops seen in EUV {\\it TRACE} images. In contrast to the general view that single loop models tend to overestimate flaring loop sizes of complex flare events, for all of our 5 testbed flares the equations of R97 yield a single-loop length comparable to or shorter than the lengths of the individual EUV loops measured from the {\\it TRACE} images. Here we are unable to give a definite answer to the question of why the loop lengths of multi-loop X-class flares derived using the R97 approach can be comparable to the observed loop lengths, considering that the application of the single-loop approach to multi-loop flares is a priori incorrect. The consistency may be by chance. For instance, for all 5 flares, the inferred slope $\\zeta$ in the $\\log T$~--~$\\log \\sqrt{EM}$ diagram is found to be close to $\\sim 0.4$. Let us ignore for a moment the fact that the physical meaning of the slope $\\zeta$ for multi-loop flare arcades is generally different from that of a single loop (i.e. different plasmas in different loops heated and cooled sequentially versus a single plasma in a single loop heated and cooled). Within the framework of the single-loop model of R97, $\\zeta \\sim 0.4$ is close to the lowest values allowed, and the equations of R97 produce smaller loop lengths at smaller slopes. In fact, in the regime of low $\\zeta$ ($\\zeta < 0.7$), the loop length is steeply decreasing with decreasing $\\zeta$, by a factor of 10 when $\\zeta$ changes from 0.7 to 0.4. Contrary to the ``by chance'' explanation above, loop length numbers can be comparable due to the presence of a single loop (or localized multiple loops ignited simultaneously and undergoing similar heating and cooling processes) that dominates the flare X-ray emission. For instance, an arcade-like structure with a single primary loop dominating the rise and early decay phases of a flare was proposed (based on a detailed modeling) for the complex flare on Proxima Centauri \\citep{Reale04}. At least for the best studied of the 5 solar X-class flares considered here, the Bastille Day flare, we do no find compelling reasons for the presence of such a dominant X-ray emitting structure (Appendix~\\ref{sec_solar_flares}). Clearly, even if the loop lengths inferred via the single-loop method happen to be comparable to the observed values for such complex flare arcade events, the method would not predict a correct loop geometry and would likely underestimate the emitting volume by at least a factor comparable to the number of instantaneously heated loops in the arcade, $N_{loop}$, and would overestimate the average electron density of emitting plasma by $\\sim \\sqrt{N_{loop}}$. For instance, during the Bastille Day flare arcade, $N_{loop} \\sim 20$ out of the $\\sim 100$ observed loops fired near the flare peak time \\citep{Aschwanden01}. For all 5 solar X-class flares considered here, the inferred slope $\\zeta$ in the $\\log T$~--~$\\log \\sqrt{EM}$ diagram is found to be close to the value of $\\sim 0.4$. Observationally, such a value is among the lowest values seen in solar flares \\citep[e.g.][]{Sylwester93}. Such a shallow slope indicates continued heating during the decay phase of the X-ray emission. Continued heating is also what one would expect for multi-loop two-ribbon flares where the decay phase could be entirely driven by the heating released through sequential reconnection in individual loops \\citep{Kopp84}. The slope $\\zeta \\ga 1$ in the $\\log T$~--~$\\log \\sqrt{EM}$ diagram for the DQ~Tau flare is significantly larger than the $\\zeta \\sim 0.4$ found for complex powerful solar flares (Appendix~\\ref{sec_solar_flares}), suggesting heating behaviour different from the 5 solar X-class flares. The possibility that the DQ~Tau flare could involve an arcade-like structure with a single loop dominating the rise and early decay phases, as was proposed for the complex flare on Proxima Centauri \\citep{Reale04}, can not be excluded. On the other hand, within the framework of a model involving explosive chromospheric evaporation into a single loop \\citep[e.g.][]{Reale07}, an evaporation time of only $3-5$~ks is needed for a plasma with temperature $T_{obs} \\sim 60-90$~MK (\\S \\ref{sec_trs}) to propagate with an isothermal sound speed $v_s$$\\rm{[cm~s^{-1}]}$$ = \\sqrt{(5/3 k_b T_{obs})/(\\mu m_p)} = 1.5 \\times 10^4 \\sqrt(T_{obs}\\rm{[K]})$ \\citep[e.g.][]{Aschwanden00} and to fill a DQ Tau flaring loop of a semi-length $L=4-5$~R$_{\\star}$. The observed rise timescale of the DQ Tau light curve, $\\tau_{rise} = 26.3$~ks, is much longer than $3-5$~ks and thus might indicate a complex history of heating during the early phase of the flare, e.g., due to multiple loops. It is also worth mentioning that the time delay between temperature and emission measure peaks observed in the DQ Tau flare (\\S \\ref{sec_trs}) can not be used to argue in favor of a single loop event. This effect is commonly seen in multi-loop flares and can be explained by invoking the principle of linear superposition (Appendix~\\ref{sec_solar_flares}). Likewise, the Neupert-like effect observed in the DQ Tau flare (\\S \\ref{sec_chandra_iram}) can not be used to argue in favor of a single loop event. If the Neupert effect occurred in each loop of a multi-loop system, then an observation that did not spatially resolve the loops would show the Neupert effect. Low-resolution observations of multi-loop solar flares commonly exhibit the Neupert effect \\citep{Veronig02}. After all, in view of the fact that the observed timescales, temperatures, and X-ray luminosities of PMS stars (DQ~Tau and ONC stars) are much higher than those of typical solar flares (\\S \\ref{sec_coup_star_comparison}), one could argue that the solar analogy might not be directly applicable. PMS and magnetically active older stars possess stronger, than solar, surface and global magnetic fields and larger volumes for magnetic fields to interact, so the extreme flaring behaviour, beyond solar analogy, is expected \\citep[e.g. \\S 2.2 in][and \\S \\ref{sec_origin} in this work]{Benz10}. But then the single-loop approach developed on solar analogy might not be applicable either. We conclude that we can neither prove the presence of loop arcades analogous to the solar cases, nor refute the presence of single flaring loops longer than any seen on the Sun. The results from the single-loop X-ray modeling presented in \\S \\ref{sec_single_loop_modeling} and \\S \\ref{sec_energy_profiles} should not be treated as definitive and should be considered with caution. \\subsection{Origin of the DQ~Tau Flare \\label{sec_origin}} The origin of the big flares from the COUP sample themselves is unclear. For some of the 32 most powerful COUP flares whose inferred coronal structures reach several stellar radii, \\citet{Favata05} suggest magnetic loops linking the stellar photosphere with the inner rim of the circumstellar disk. For the extended sample of $>200$ flares G08ab find that large (a few to several stellar radii in length based on the single-loop approach of R97) coronal structures are present in both disk-bearing and disk-free stars; however G08ab also find a subclass of super-hot flares with peak plasma temperatures exceeding 100~MK that are preferentially present in highly accreting systems. G08ab further propose that the majority of big flares from the COUP sample can be viewed as enhanced analogs of the rare solar long-decay events (LDEs; see \\S 7.3 in G08b and references therein). LDEs are eruptive solar flare events that produce X-ray emitting arches and streamers with altitudes reaching up to several hundred thousand kilometers ($L \\sim 0.5-1$~R$_{\\odot}$). In panels (b, d, e) of Figure \\ref{fig_comparison_with_coup_oldstars}, representative solar LDEs (gray diamonds; compiled in G08a) are shown to have systematically higher flare durations and coronal lengths than the more typical solar flares (solid and dotted gray loci). We wish to emphasise that the analogies of the big flares from the COUP sample to solar LDEs (``growing'' systems of giant multiple loops) are not based on information about the specific detailed geometry of COUP flares (the geometry is really un-known to us and is only simplistically modeled as a large single loop), but are instead based on the simple fact that the observed flare durations and model-inferred characteristic loop scales for COUP flares are the largest ever reported from PMS stars, provided that these model-inferred loop scales are close to the truth. The most widely accepted model for the origin of LDEs is that the impulsive flare near the solar surface ($L\\la 10^{-2}$~R$_{\\odot}$) blows open the overlying large-scale magnetic field with subsequent reconnection of large-scale magnetic lines through a vertical current sheet. The large-scale magnetic field of PMS stars is likely far stronger than in the Sun, and can sustain giant X-ray arches and streamers with sizes $L\\sim 1-10$~R$_{\\star}$. For the DQ~Tau flare a different mechanism from that of a solar LDE-like reconnection can be considered. There are at least three independent supporting lines of evidence suggesting that the DQ~Tau flare could be produced through a process of colliding magnetospheres. \\begin{enumerate} \\item Although the DQ~Tau flare duration, morphology, and plasma temperature are typical of those of big flares from the COUP sample (\\S \\ref{sec_coup_star_comparison}), the probability of observing a big COUP-like flare is small. Within the COUP observation window ($\\sim 1200$~ks) on average 3 big flares per bright PMS star were detected. This number is derived from the analysis of the 161 brightest X-ray PMS stars (G08a); the frequency of big flares for the remaining 1200 fainter COUP PMS stars is even lower. The DQ Tau $Chandra$ observation was designed for an X-ray flare to appear close to the orbital phase of $\\Phi = 0.98$, which corresponds to the peak of the April~2008 large sub-mm flare \\citep{Salter08}. The detected $Chandra$ flare indeed peaks within $15$~ks of this orbital phase. The probability of detecting a big COUP-like flare within $15$~ks of a pre-specified point in time is small ($Prob < 15 \\times 2 / (1200 / 3) = 0.075$). \\item Due to the lack of systematic monitoring observations in mm-bands, mm flare activity in PMS stars has been rarely reported \\citep[e.g.][]{Bower03, Salter10}. The DQ~Tau X-rays are accompanied by a relatively unique non-thermal mm activity. The X-rays and mm are likely related through a Neupert-like effect (\\S \\ref{sec_chandra_iram}). The observed re-appearance of this unique mm activity near periastron is proposed to be associated with processes caused by the interacting magnetospheres \\citep{Salter08,Salter10}. \\item It is natural to expect large-scale magnetic structures to be involved in a process of colliding magnetospheres. The loop flaring sizes inferred from the X-ray flare analysis ($4-5$~R$_{\\star}$) are comparable to half the separation of the DQ~Tau components near periastron (\\S \\ref{sec_loop_length}). This third item of evidence is completely model-dependent and should be considered with caution (\\S \\ref{sec_single_vs_multiple}). \\end{enumerate} \\subsection{Energetics and Loop Geometry \\label{sec_e_l_b}} Based on the concept of interacting magnetospheres, we can further comment on energetics and possible flare loop geometry. We can show that a crude estimate of the power expected to be dissipated by magnetic reconnection from colliding magnetospheres in DQ~Tau is in agreement with the derived peak flare X-ray luminosity ($L_{X,0.5\\_8} = (4-5) \\times 10^{30}$~ergs~s$^{-1}$; see Table \\ref{tbl_trs}) and the modeled rate of the kinetic energy injected into the chromosphere by non-thermal electrons ($\\alpha \\times F_R(t_1) = (0.6-0.8) \\times 10^{32}$~ergs~s$^{-1}$; \\S \\ref{sec_chandra_iram}). For a large scale dipolar topology and a magnetic field $B(L) \\simeq B_{ph}/(L/R_{\\star}+1)^3$ with photospheric field strengths in the range $1-6$~kG consistent with measurements of Zeeman broadening and circular polarization of photospheric lines in PMS stars \\citep[e.g.][]{Johns-Krull07,Donati08}, the field strength at distances of $4-5$~R$_{\\star}$ from the stellar surface is expected to be $B(4-5 \\rm{R}_{\\star}) \\simeq 5-50$~G. Within the framework of the simple magnetic reconnection model shown in Figure~\\ref{fig_magneto_inter}, the initially dipolar fields of average strength $B = 5-50$~G at $4-5$~R$_{\\star}$ from the stellar surface colliding at a speed of $v \\sim 100$~km/s \\citep{Mathieu97} reconnect with the rate of energy release $E_m/\\tau_R \\sim 10^{30} - 10^{32}$~ergs~s$^{-1}$, which is in agreement with the numbers given above. It is interesting to note that the analogous procedure applied to cases of star-planet magnetic interaction have difficulty explaining an excess X-ray emission associated with stellar chromospheric hot spots rotating synchronously with close-in giant planets \\citep{Lanza09}. The model of interacting magnetospheres in RS~CVn-type binaries by \\citet{Uchida85} predicts the existence of X-ray emitting loop structures connecting the stars. The model of star-planet magnetic interaction by \\citet{Lanza09} expects complex topologies, including loops connecting the planet with the stellar surface. While our X-ray modeling method and many related results are limited to the scenario of a single X-ray emitting loop on a single star, other scenarios for the DQ~Tau loop geometry are possible: multiple X-ray emitting loops on a single star (\\S \\ref{sec_single_vs_multiple}), two or more loops appearing on both stars simultaneously, or a single or multiple loops connecting the two stars. For our modeled scenario, where the loop's footprints are anchored on the surface of a single star, potential destruction of the loop by centrifugal forces is not a serious issue. The derived length for the X-ray emitting structure of $4-5$~R$_{\\star}$ is only comparable to and does not exceed the Keplerian corotation radius $R_{cor} = 4.7$~R$_{\\star}$\\footnote{The Keplerian corotation radius where the centrifugal force balances the gravitational force is $R_{cor} = (G \\times M \\times P^2/4 \\times \\pi^2)^{1/3}$. For the DQ~Tau binary components with similar masses of $M=1.6$~M$_{\\odot}$ and rotational periods of $P=3$~days, $R_{cor} = 4.7$~R$_{\\star}$.}." }, "1101/1101.4473.txt": { "abstract": "%% Text of abstract The advent of the Auger Engineering Radio Array (AERA) necessitates the development of a powerful framework for the analysis of radio measurements of cosmic ray air showers. As AERA performs ``radio-hybrid'' measurements of air shower radio emission in coincidence with the surface particle detectors and fluorescence telescopes of the Pierre Auger Observatory, the radio analysis functionality had to be incorporated in the existing hybrid analysis solutions for fluoresence and surface detector data. This goal has been achieved in a natural way by extending the existing Auger Offline software framework with radio functionality. In this article, we lay out the design, highlights and features of the radio extension implemented in the Auger Offline framework. Its functionality has achieved a high degree of sophistication and offers advanced features such as vectorial reconstruction of the electric field, advanced signal processing algorithms, a transparent and efficient handling of FFTs, a very detailed simulation of detector effects, and the read-in of multiple data formats including data from various radio simulation codes. The source code of this radio functionality can be made available to interested parties on request. ", "introduction": "Forty years after the initial discovery of radio emission from extensive air showers \\citep{JelleyFruinPorter1965}, the CODALEMA \\citep{ArdouinBelletoileCharrier2005} and LOPES \\citep{FalckeNature2005} experiments have re-ignited very active research activities in the field of radio detection of cosmic ray air showers. Nowadays, the field is in a phase of transition from first-generation experiments covering an area of less than 0.1 km$^2$ to large-scale arrays of tens of km$^2$. In particular, the Auger Engineering Radio Array (AERA) \\citep{HuegePisa2009} will complement the southern site of the Pierre Auger Observatory \\citep{AugerNIM2004} with $161$ autonomous radio detector stations covering an area of $\\approx 20$~km$^2$. One particular merit of the Pierre Auger Observatory is its hybrid mode of observation, which uses coincident detection of extensive air showers with both optical fluorescence telescopes (FD) and surface particle detectors (SD) to gain in-depth information on the measured air showers. Consequently, the analysis software has to support complete hybrid processing and interpretation of the data. This requirement is fulfilled by the Auger Offline software framework \\citep{ArgiroOffline2007}. To take full advantage of the radio data taken in the hybrid environment of the Pierre Auger Observatory, it is clear that also radio analysis functionality, which has so far been existing in a separate software package \\citep{FliescherArena2008}, had to be included in this hybrid analysis framework. In this article, we describe how we have therefore built advanced radio analysis functionality into the Auger Offline software framework. The general structure of the radio implementation in the Offline framework will be discussed in section \\ref{sec:structure}. A number of innovative features have been realized in this context for the very first time. These and other highlights will be discussed in section \\ref{sec:highlights}. Finally, in section \\ref{sec:analysis} we demonstrate how the advanced radio functionality embedded in the Offline framework can be used to carry out a complete detector simulation and event reconstruction on the basis of a simulated radio event. ", "conclusions": "We have implemented a complete set of radio analysis functionality in the Offline software framework of the Pierre Auger Observatory. The radio functionality has been included in a canonical and seamless way in addition to the existing SD and FD functionality. This approach will make the realization of ``radio-hybrid'' analysis strategies in the future straight-forward. Already now, however, the radio functionality in Offline has reached a high degree of sophistication with highlights such as a very fine-grained simulation of detector effects, advanced signal processing algorithms, transparent and efficient handling of FFTs, read-in of multiple file formats for measured and simulated radio data, and in particular the reconstruction of the three-dimensional electric field vector from two-dimensional measurements. Planned improvements encompass the implementation of a curved fit, inclusion of interferometric radio analysis functionality, and the handling of a time-variable detector including a fine-grained treatment of the instrumental calibration. Parties interested in using the functionality are encouraged to contact the corresponding author. The source code can be made available on request." }, "1101/1101.2255_arXiv.txt": { "abstract": "Plasmoid structures in fast reconnection in low-beta plasmas are investigated by two-dimensional magnetohydrodynamic simulations. A high-resolution shock-capturing code enables us to explore a variety of shock structures: vertical slow shocks behind the plasmoid, another slow shocks in the outer-region, and the shock-reflection in the front side. The Kelvin-Helmholtz-like turbulence is also found inside the plasmoid. It is concluded that these shocks are rigorous features in reconnection in low-beta plasmas, where the reconnection jet speed or the upstream Alfv\\'{e}n speed exceeds the sound speed. ", "introduction": "Magnetic reconnection \\cite{sweet,parker,petschek} is a fundamental mechanism to power various events in laboratory, space, solar, and astrophysical plasmas. The reconnection process releases a fast plasma jet by consuming the magnetic energy stored in the upstream region. Such a plasma jet further interacts with the external environments and exhibits complex plasma phenomena. One of the most characteristic phenomena is a ``plasmoid'', a magnetic island with plasmas embedded inside the magnetic loop. Plasmoids are well accepted to explain the satellite observation in the Earth's magnetotail \\cite{hones77,slavin84,slavin03} and morphological features in solar flare and coronal mass ejections \\cite{shibata95,lin08,linton09}. Motivated by the above and many other important reasons, the nonlinear evolution of magnetic reconnection and the associated plasmoid evolution has been studied over many decades. Significant insights have been obtained by magnetohydrodynamic (MHD) simulations in two \\cite{ugai77,forbes83,ugai86,ugai87,scho89,ugai92,ugai95,nitta01,shuei01,bis01,tanuma07,barta08,murphy10,yu11} and three dimensions \\cite{birn81,hesse91,ugai91,ugai96,ugai08}. These simulations showed that the nonlinear evolution of the reconnection-plasmoid system involves complicated structures, such as slow shocks of the Petschek outflow \\cite{petschek} and various shocks around the plasmoid \\cite{forbes83,ugai87,ugai95,shuei01}. Recently, reconnection research has been rapidly developing in relativistic astrophysics, too.\\cite{zeni05b,naoyuki06,zeni09a,zeni10b} In this context \\citet{zeni10b} carried out MHD simulations on relativistic magnetic reconnection. They discussed several shock structures that were not known in the nonrelativistic plasmoid physics: the postplasmoid vertical shocks and the shock-reflection structure in the plasmoid-front. It is known that reconnection is highly influenced by upstream conditions. One key parameter is the Alfv\\'{e}n speed, which usually approximates the outflow jet speed. Another key parameter is the plasma beta, the ratio of the plasma pressure to the magnetic pressure. As will be discussed later, it corresponds to the ratio of the Alfv\\'{e}n velocity and the sound velocity. The beta is typically below the unity in the lower corona \\citep{gary01,aschwanden} and the lobe (upstream region) of the magnetotail plasma sheet \\citep{baum89}. In this paper, we explore fine structures of a plasmoid in fast reconnection when the upstream region consists of low-beta plasmas by using two-dimensional MHD simulations. Employing a high-resolution shock-capturing (HRSC) code, we successfully resolve a variety of shock structures. We confirm that the new shock structures in Ref. \\onlinecite{zeni10b} similarly appear in the nonrelativistic regime. We further discuss the condition for these shocks, and argue that they are common features of plasmoids in the low-beta plasmas. We also report some new features such as the Kelvin-Helmholtz-like turbulence inside the plasmoid and a potential signature of the corrugation instability of the slow shock surface. This paper is organized as follows. We describe our numerical setup in Sec. \\ref{sec:setup}. We present simulation results and then investigate the local structures in depth in Sec. \\ref{sec:results}. Section \\ref{sec:beta} presents a brief parameter survey. We discuss these results in Sec. \\ref{sec:discussion}, particularly focusing on the plasma beta condition. Section \\ref{sec:summary} summarizes this paper. ", "conclusions": "\\label{sec:discussion} In this paper we visited many structures in and around the plasmoid. They are schematically illustrated in Fig. \\ref{fig:illust}. This extends an informative result in an earlier work (i.e., Fig. 10 in Ref.~\\onlinecite{shuei01}). Among these structures, the postplasmoid vertical shocks (Sec. \\ref{sec:shock}) and the shock-reflected diamond-chain structure in the front side (Sec. \\ref{sec:front}) were recently found in the relativistic regime.\\cite{zeni10b} We confirmed them in the nonrelativistic regime as well --- they are ubiquitous, regardless of the relativistic effects. A reverse flow or a backward flow around the plasmoid is often found in recent large-scale simulations\\cite{shuei01,zeni09a}, however, we clearly demonstrated that a slow shock stands at the front of the reverse flow. We also introduced another vertical slow shock outside the plasmoid. Those two vertical shocks are similarly found in other runs in our parameter survey. We think that the vertical slow shocks are logical consequences of the low-beta condition of $\\beta_{up} \\lesssim 1$. From the following relation, \\begin{eqnarray} \\label{eq:beta} ( {c_{s,up}}/{c_{A,up}} )^2 = ({\\Gamma}/{2}) \\beta_{up} , \\end{eqnarray} we see % $c_{s,up} < c_{A,up}$ when $\\beta_{up} \\lesssim 1$. As already discussed, the plasmoid-front propagates at a sizable fraction of the reconnection jet speed $v_{jet}\\sim c_{A,up}$. When $\\beta_{up} \\lesssim 1$, the plasmoid supersonically travels in the system. At some location, the expanding plasmoid structure interacts with the surrounding plasmas at a supersonic speed, and then a shock stands there. Another interpretation is as follows. In a self-similarly expanding coordinate of the plasmoid system,\\cite{nitta01} the surrounding plasmas supersonically flow in the $-x$-direction near the plasmoid-front, while the relative flow is very small near the reconnection point. A shock stands at the point where the supersonic flow slows down to subsonic speed. In summary, when $\\beta_{up} \\lesssim 1$, the plasmoid system will inevitably involve vertical slow shocks. In addition, the compression effects by the plasmoid introduces another pair of vertical shocks outside the plasmoid. As discussed in Sec. \\ref{sec:shock}, the compression invokes the field-aligned flows in the ${\\pm}x$ directions. In addition, the adiabatic heating or cooling weakly modify the local sound speed, $c_s = \\sqrt{\\Gamma s^{\\frac{1}{\\Gamma}}} \\cdot p^{(\\Gamma-1)/{2\\Gamma}} \\propto p^{{1}/{5}}$. In the main run, the relative velocity between the plasmoid system and the upstream plasmas are initially supersonic in the right side. It temporally becomes subsonic at the forward vertical shock $x \\sim 100$ (Sec. \\ref{sec:outer}) due to the compression effects. However, due to the field-aligned acceleration in the $-x$-direction, the flow becomes supersonic again, and then it becomes subsonic at the postplasmoid vertical shock at $x\\sim 51$ (Sec. \\ref{sec:shock}). The positions of the vertical shocks depend on the $\\beta_{up}$ parameter, because it controls the ratio of the sound speed to the Alfv\\'{e}nic plasmoid motion. When $\\beta_{up}$ is lower than the main run ($\\beta_{up}=0.1$), the expanding speed of the entire structure easily exceeds the sound speed. Therefore, the postplasmoid vertical shock is closer to the reconnection point, as demonstrated in Sec. \\ref{sec:beta} (Fig.~\\ref{fig:beta}a). The forward vertical slow shock can be small, when the plasmoid compression effects are not strong enough to change the supersonic condition. Probably the forward vertical shock disappears in the extreme limit of $\\beta_{up} \\ll 1$. On the other hand, when $\\beta_{up}$ is higher ($\\beta_{up} = 0.5$), the vertical slow shocks are nearer to the front-side of the plasmoid: the postplasmoid and forward ones are located just behind and in front of the plasmoid, respectively. As $\\beta_{up}$ increases, the postplasmoid slow shock moves frontward even in the backward side of the plasmoid outer shell, and then it will eventually disappear. The forward slow shock will disappear when the Alv\\'{e}n speed is no longer supersonic, $\\beta_{up} \\sim 2/\\Gamma = 1.2$. Depending on $\\beta_{up} (\\lesssim 1)$, we expect one or two pairs of vertical slow shocks around the plasmoid. Our discoveries can readily be applied to various numerical results in low-beta plasmas. \\cite{shuei01,tanuma07,zeni10b,yu11} \\citet{shuei01} carried out large-scale MHD simulation of a plasmoid in a similar configuration. While they did not mention it in the text, one can see a shock-like structure outside the plasmoid in RUN 2 with $\\beta_{up}=0.25$ (the top panel in Fig. 9 in Ref. \\onlinecite{shuei01}). We think that this is a forward vertical slow shock. \\citet{zeni10b} recognized a small shock outside the plasmoid (indicated by an arrow in Fig. 2c in Ref. \\onlinecite{zeni10b}) in the relativistic regime with $\\beta_{up}=0.1$. Inspired by this work, we carry out further analysis and confirm that it was a forward vertical slow shock. In multiple plasmoid systems, \\citet{tanuma07} found ``bow shocks'' during the nonsteady reconnection under an anomalous-type resistivity with $\\beta_{up}=0.2$ (Fig. 1d in Ref. \\onlinecite{tanuma07}). Judging from displacements, their plasmoid velocities are sub-Alfv\\'{e}nic, on an order of ${\\sim}c_{s,up} < c_{A,up}$. In such a developing stage, the forward vertical shock is located near the front, and so it would be a forward slow shock. Finally, a recent work by \\citet{yu11} reported a steepening structure of slow-mode wave behind the primary plasmoid (Fig. 15a in Ref.~\\onlinecite{yu11}). It would be related to a postplasmoid slow shock, as we found it in our relevant run with $\\beta_{up} =0.5$ (Sec. \\ref{sec:beta}; Fig. \\ref{fig:beta}b). The shock-reflection or the diamond-chain structure is also a characteristic in the low-beta condition. Dropping a factor of order unity, the requirement (Eq. \\ref{eq:diamond}) reads \\begin{eqnarray} \\Big( \\frac{c_{s,cs}}{c_{A,up}} \\Big) \\sim \\Big( \\frac{c_{s,cs}}{c_{s,up}} \\Big) \\sqrt{ \\beta_{up} } < 1. \\end{eqnarray} Although we still have a free parameter $( {c_{s,cs}}/{c_{s,up}} )$, the condition can be easily satisfied in the low-$\\beta_{up}$ regimes. Let us discuss some limitations of our simulation model. First, we assumed that the evolution in the bottom quadrant ($z<0$) is symmetric to that in the upper quadrant ($z>0$). We do not think the quadrant assumption affects the main shock structures, because Ref. \\onlinecite{zeni10b} was carried out without the assumption. On the other hand, we found the Kelvin-Helmholtz-like turbulence in the internal region (Sec. \\ref{sec:internal}). Since this class of instabilities may prefer odd-parity perturbation across the neural line ($z=0$), the central Harris sheet may flap in very late stages. This is left for future investigation. Second, in the three dimensions, it is known that quasiperpendicular slow shocks are unstable to the corrugation instability.\\cite{stone95} We do see a potential signature of the instability even in the two dimensions (Sec. \\ref{sec:beta}). In the real world, the vertical slow shocks may be violently folded and then we will see turbulent transition layers instead. Regarding the shock-reflected diamond-chain structure, the oblique fast shocks will be stable\\cite{gardner64}, but it is not clear how stable the intermediate shocks are. It is interesting to see how those structures are modulated in the three dimensions in future. We carried out HRSC MHD simulations to study the fine structures of a plasmoid. We introduced new shock structures: (1) postplasmoid vertical slow shocks, (2) forward vertical slow shocks, and (3) shock-reflection structures, some of which were recently reported in Ref. \\onlinecite{zeni10b}. We further found new features such as (4) the flapping motion of the reflected jet and (5) the finger-like instability of the slow shock surface. We interpreted that the vertical slow shocks and the shock-reflection are consequences of the Alfv\\'{e}nic plasmoid motion, which is supersonic in low-beta plasmas. We argue that these shocks are rigorous features of a reconnection system in low-beta plasmas." }, "1101/1101.2125_arXiv.txt": { "abstract": "The \\emph{AGILE} $\\gamma$-ray satellite provides large sky exposure levels ($\\geq$10$^9$ cm$^2$ s per year on the Galactic Plane) with sensitivity peaking at $E\\sim$100 MeV where the bulk of pulsar energy output is typically released. Its $\\sim$1 $\\mu$s absolute time tagging capability makes it perfectly suited for the study of $\\gamma$-ray pulsars. \\emph{AGILE} collected a large number of $\\gamma$-ray photons from EGRET pulsars ($\\geq$40,000 pulsed counts for Vela) in two years of observations unveiling new interesting features at sub-millisecond level in the pulsars' high-energy light-curves, $\\gamma$-ray emission from pulsar glitches and Pulsar Wind Nebulae. \\emph{AGILE} detected about 20 nearby and energetic pulsars with good confidence through timing and/or spatial analysis. Among the newcomers we find pulsars with very high rotational energy losses, such as the remarkable PSR\\,B1509--58 with a magnetic field in excess of 10$^{13}$ Gauss, and PSR\\,J2229+6114 providing a reliable identification for the previously unidentified EGRET source 3EG\\,2227+6122. Moreover, the powerful millisecond pulsar B1821--24, in the globular cluster M28, is detected during a fraction of the observations. ", "introduction": "Poor $\\gamma$-ray pulsar statistics has been a major difficulty in assessing the dominant mechanism which channels pulsar rotational energy into high energy emission as well as understanding the sites where charged particles are accelerated. The large field of view of the AGILE Gamma-Ray Imaging Detector (GRID) \\cite{tavani09} allows long uninterrupted observations and simultaneous monitoring of tens of nearby radio pulsars belonging to the ``$\\gamma$-ray pulsar region'' of the $P$--$\\dot{P}$ diagram characterized by $B>2\\times10^{11}$ G and spin-down energy $\\dot{E}_{\\rm{rot}}>1.3\\times10^{33}$ erg s$^{-1}$ \\cite{pellizzoni04}. Here we present the results of three years of pulsar observations with \\agile. ", "conclusions": "" }, "1101/1101.1917_arXiv.txt": { "abstract": "We examine the wavelength dependence of flux ratios for six gravitationally lensed quasars using $K$ and $L'$ images obtained at the Gemini North 8m telescope. We select lenses with source redshifts $z_s < 2.8$ so that $K$-band images probe rest-frame optical emission from accretion disks, while $L'$-band images probe rest-frame near-infrared flux emitted (in part) from the more extended surrounding torus. Since the observations correspond to different source sizes, the $K$ and $L'$ flux ratios are sensitive to structure on different scales and may be useful for studying small-structure in the lens galaxies. Four of the six lenses show differences between $K$ and $L'$ flux ratios. In HE 0435$-$1223, SDSS 0246$-$0825, and HE 2149$-$2745 the differences may be attributable to known microlensing and/or intrinsic variability. In SDSS 0806+2006 the wavelength dependence is not easily attributed to known variations, and may indicate the presence of substructure. By contrast, in Q0142$-$100 and SBS 0909+523 the $K$ and $L'$ flux ratios are consistent within the uncertainties. We discuss the utility of the current data for studying chromatic effects related to microlensing, dust extinction, and dark matter substructure. ", "introduction": "\\label{sec:irintro} While the cold dark matter (CDM) paradigm for structure formation successfully describes cosmological observations on large (CMB and cluster) scales, there is notable disagreement with small-scale observations. Among other issues, $N$-body simulations (e.g., Via Lactea -- Diemand et al. 2008; Aquarius -- Springel et al. 2008) predict the existence of numerous CDM subhalos, with masses $M \\sim 10^4$--$10^9~M_\\odot$, embedded in galaxy-scale dark matter halos. This has proved troubling observationally, because there are many fewer dwarf galaxies in our own Milky Way than predicted by CDM. Since the discrepancy may be due to baryon stripping from subhalos \\citep[e.g.,][]{2008ApJ...679.1260M, 2010MNRAS.402.1995M}, we need ways to probe dark matter substructure directly, regardless of the presence of baryonic material. Gravitational lensing provides a unique way to detect CDM substructure in distant galaxies \\citep[e.g.,][]{2001ApJ...563....9M, 2002ApJ...565...17C,2002ApJ...572...25D}.\tStars and CDM substructure perturb the lens potential on micro- to milli-arcsecond scales, which can have dramatic effects on the properties of lensed images. Most notably, lensing from stars and dark matter substructure can alter the flux ratios from those of smooth mass distributions. As shown by \\citet{2006MNRAS.365.1243D}, lens flux ratios depend on the size of the source compared to the size of the perturber. When the source is very small, it is effectively a point source for the lens substructure, so it feels an additional magnification boost ($B_{\\rm sub} \\equiv \\mu_{\\rm sub}/\\mu_{\\rm smooth}$) over and above the macroscopic lens properties. As the source increases in size, $B_{\\rm sub}$ may increase or decrease depending on the source's location relative to the substructure and the parity of the ``macro'' magnification. As the source grows still larger, it ceases to be affected by the substructure and $B_{\\rm sub} \\rightarrow 1$. This phenomenon implies that by measuring flux ratios at different wavelengths, corresponding to different source sizes, we may be able to map substructure on a variety of scales \\citep[also see][]{2003MNRAS.339..607M}. Heuristically, a quasar emitting region of size $R_S$ is significantly affected by a subhalo with Einstein radius $R_E$ only if $R_S \\lesssim R_E$. For typical lens and source redshifts (e.g., $z_l=0.5$ and $z_s=2.0$), the Einstein radius of a subhalo of mass $M$ is $R_E\\sim 10^{16}~\\mbox{cm}~(M/M_\\odot)^{1/2}$. Since the optically emitting regions of quasars have $R_S \\sim 10^{15}$--$10^{16}$ cm \\citep{2000MNRAS.315...62W,2007ApJ...661...19P,2010ApJ...712.1129M}, optical lens flux ratios are sensitive to both microlensing by stars and millilensing by CDM substructure. By contrast, the more extended infrared emitting regions with $R_S \\gtrsim 1$ pc \\citep{2005ApJ...627...53C, 2009ApJ...697..610M, 2009ApJ...697.1010A} can only be affected by relatively massive subhalos. Comparing lens flux ratios at different wavelengths therefore makes it possible to constrain the amount of micro- and milli-lensing present in the system, as well as the sizes of the perturbers. Previous studies have used mid-IR observations to probe rest-frame IR emission, yielding evidence for subhalos with masses $\\gtrsim 10^5 M_\\Sun$ in the lenses B1422+231 and MG 0414+0534 \\citep{2005ApJ...627...53C, 2009ApJ...697..610M}, constraints on the mass of a luminous companion galaxy in H1413+117 \\citep{2009ApJ...699.1578M},\\footnote{\\citet{2009MNRAS.394..174M} also constrained the mass of a luminous satellite, but using radio observations.} and null detections in several other systems \\citep[][]{2005ApJ...627...53C, 2009ApJ...697..610M, 2009ApJ...697.1010A}. Here we extend the study of wavelength-dependent flux ratios by using $K$ ($2.2\\,\\mu$m) and $L'$ ($3.8\\,\\mu$m) images of six lenses obtained with Gemini North during the 2008B semester. For source redshifts $z_s < 2.8$, the $L'$-band images correspond to rest-frame emission at $>1\\,\\mu$m where $\\sim$20--100\\% of the flux is thermal emission from the inner dusty torus \\citep{1995MNRAS.272..737R, 2004ApJ...600L..35M, 2008ApJ...685..160N, 2009ApJ...697.1010A}. By contrast, the $K$-band flux comes mostly from the smaller accretion disk. Thus, comparing $K$ and $L'$ flux ratios may provide a sufficient source size baseline to identify substructure. To be sure, there are phenomena besides CDM substructure that may cause lens flux ratios to vary with wavelength. Optical observations probe rest-frame UV emission from the accretion disk of the AGN, so they are sensitive to microlensing by stars in the lens galaxy \\citep[e.g.,][]{2009ApJ...693..174C, 2010ApJ...712.1129M}. Optical flux ratios can also be altered by differential extinction from dust in the lens galaxy \\citep[e.g.,][]{2006ApJS..166..443E}. Finally, intrinsic variability in the source coupled with the lens time delay can cause flux ratios to vary with time, and the variability may be chromatic.\\footnote{Photometric monitoring can be used to quantify the variations, but has only been done for certain lenses \\citep{2006A&A...451..747E,2006ApJ...640...47K,2007A&A...464..845V,2008A&A...488..481V,2008ARep...52..270K,2008NewA...13..182G}. Four of our targets have been monitored at optical wavelengths: Q0142$-$100, HE 0435$-$1223, SBS 0909+523, and HE 2149$-$2745.} All three effects should be attenuated at near-IR wavelengths, because the effective size of the accretion disk is larger and some of the flux originates in the larger dusty torus, but they may not be entirely absent. In particular, the importance of microlensing and intrinsic variability for $L'$-band observations will depend on the relative strengths of the accretion disk (0.01-- 0.1 pc, $10^{-6}$--$10^{-5}$ arcsec) and dusty torus (0.5--5 pc, $10^{-4}$--$10^{-3}$ arcsec) emission. Any significant wavelength-dependence in lens flux ratios is interesting, whether related to CDM substructure or not, so at minimum our observations can highlight systems that warrant further study. Most lensing constraints on CDM substructure have relied on the identification of flux ratio anomalies, which is best done using four-image lenses. The reason is that identifying such anomalies requires either well-constrained lens models \\citep[e.g.,][]{2002ApJ...572...25D, 2002ApJ...567L...5M} or universal magnification relations that apply only to certain four-image configurations \\citep{2003ApJ...598..138K, 2005ApJ...635...35K}. However, the search for wavelength dependence in flux ratios is a purely observational task that may provide model-independent evidence for substructure. We are therefore able to analyze two-image lenses for the presence of substructure for the first time. Some care is needed when interpreting the results (see \\S \\ref{sec:irdiscuss}), but the observations are still valuable. ", "conclusions": "We have used a multi-wavelength analysis of lens flux ratios to search for effects of dark matter substructure, along with quasar variability, microlensing, and dust extinction. We have presented new $K$ and $L'$ images of six lensed quasars, which were selected to have source redshifts $z_s<2.8$ so that $L'$ observations probe rest-frame wavelengths $>1.0\\,\\mu$m. Some of the $L'$ flux should therefore originate from the extended torus of gas surrounding the central accretion disk, possibly providing the conditions for chromatic millilensing \\citep[see][]{2006MNRAS.365.1243D}. We find strong differences between the $K$ and $L'$ flux ratios for two lenses. In HE 0435$-$1223 the $L'$ measurement of the B/C flux ratio is consistent with results at shorter wavelengths, but the $K$-band value is enhanced by $\\sim$30\\%. We argue in a separate analysis (Fadely \\& Keeton, in prep.) that the enhancement may be attributable to microlensing. To test that hypothesis, future observations are needed to look for variability in the $K$-band and compare it with variability at shorter wavelengths. In SDSS 0806+2006 we do not detect image B in $L'$ observations, with an upper limit of $F_B/F_A<0.164$ ($3\\sigma$), even though the flux ratio is $F_B/F_A>0.4$ at shorter wavelengths. It would be very interesting to use deeper observations at $L'$ to find image B, to measure the flux ratio at other long wavelengths, and to combine all the observations with lens modeling to ascertain whether the $L'$ anomaly may be caused by dark matter substructure or some other interesting effect. We find a $\\sim$30\\% difference between the $K$ and $L'$ flux ratios for SDSS 0246$-$0825, and a smaller but statistically significant difference for HE 2149$-$2745. Multi-epoch observations of both systems have shown variations in the flux ratios at shorter wavelengths, which are presumably associated with quasar variability and/or microlensing. It is possible that the same phenomena affect the $K$-band flux ratios and explain the NIR discrepancies. In that case, the flux ratios in $L'$ or other longer-wavelength bands would be useful for limiting effects of quasar variability and microlensing in order to probe dark matter substructure. Alternatively, it remains possible that chromatic millilensing causes the NIR discrepancies. To test both hypotheses, future observations are needed to quantify variability in multiple optical and NIR bands. Finally, we detect no difference between the $K$ and $L'$ flux ratios for Q0142$-$100 and SBS 0909+523. These results are interesting null detections, indicating that there is no substructure close to the images with masses in the right range to produce chromatic effects. Such results offer constraints on dark matter substructure that may be mild but nevertheless interesting for future statistical studies of substructure populations." }, "1101/1101.3912_arXiv.txt": { "abstract": "{} { We aim to study excitation of the observed $\\sim$ $5$-min oscillations in the solar corona by localized pulses that are launched in the photosphere. } {We solve the full set of nonlinear one-dimensional Euler equations numerically for the velocity pulse propagating in the solar atmosphere that is determined by the realistic temperature profile. } {Numerical simulations show that an initial velocity pulse quickly steepens into a leading shock, while the nonlinear wake in the chromosphere leads to the formation of consecutive pulses. The time interval between arrivals of two neighboring pulses to a detection point in the corona is approximately $5$ min. Therefore, the consecutive pulses may result in the $\\sim$ $5$-min oscillations that are observed in the solar corona.} {The $\\sim$ 5-min oscillations observed in the solar corona can be explained in terms of consecutive shocks that result from impulsive triggers launched within the solar photosphere by granulation and/or reconnection. } \\titlerunning{$5$-min oscillations in the solar corona} \\authorrunning{T.V. Zaqarashvili et al.} ", "introduction": "Propagating acoustic waves are frequently seen in the solar corona as periodic variations of spectral line intensity (De Moortel et al. \\cite{de Moortel00,de Moortel02}, Marsh et al. \\cite{marsh03}, Lin et al. \\cite{lin2005,lin2006}, Srivastava et al. \\cite{srivastava08}, Wang et al. \\cite{wang09}). As these waves are often observed within the frequency range corresponding to the acoustic waves in the solar photosphere/chromosphere, this logically leads to the idea of penetration of the photospheric acoustic oscillations into the corona. However, the photospheric $5$-min oscillations are evanescent in the gravitationally stratified solar atmosphere as their frequency is lower than the cut-off frequency (Lamb \\cite{lamb1908}, Roberts \\cite{roberts}, Musielak et al. \\cite{musielak2006}). Bel \\& Leroy (\\cite{bel77}) suggested that the cut-off frequency of the magnetic field-free atmosphere is lower for waves propagating obliquely to the vertical direction. De Pontieu et al. (\\cite{dep05}) proposed that p-modes may be channeled into the solar corona along inclined magnetic field lines as a result of the decrease of the acoustic cut-off frequency. McIntosh \\& Jefferies (\\cite{mcintosh06}) found the observational justification of the modification of the cut-off frequency by inclined magnetic field. As the magnetic field of active region loops is predominantly vertical in the photosphere/chromosphere, it is unclear how p-modes may penetrate into the coronal regions. There are two different types of drivers in the highly dynamic solar photosphere: oscillatory (e.g. p-modes) and impulsive (e.g. granulation and/or explosive events due to magnetic reconnection). Both types of drivers may be responsible for the observed dynamical phenomena in upper atmospheric regions. As a result of the rapid decrease of mass density with height, finite-amplitude high-frequency photospheric oscillations can quickly grow in their amplitudes and steepen into shocks, which by energy dissipation can lead to the chromospheric heating (Narain \\& Ulmschneider \\cite{narain,narain96}, Carlsson \\& Stein \\cite{carlson97}, Ruderman \\cite{ruderman}). Lower-frequency waves, those with $\\sim$5 min period, are not good candidates for the chromospheric heating (Narain \\& Ulmschneider \\cite{narain}). It was found by Hollweg (\\cite{hol82}) that a localized pulse that is launched initially sets up a nonlinear wake which results in a trail of consecutive shocks. Such shocks were called rebound shocks by Hollweg (\\cite{hol82}). The time interval between consecutive shocks is close to the period of the nonlinear wake. In the linear case, the wake oscillates with the acoustic cut-off frequency of the stratified atmosphere. Nonlinearity modifies the wave period of the wake, with many features of spicules which exhibit the periodicity of about $5$-min (Murawski \\& Zaqarashvili \\cite{murawski2010}). Then, these quasi-periodic shocks may lead to the oscillatory dynamics of coronal plasma, which is observed as intensity oscillations in coronal spectral lines. The aim of this paper is to study the role of rebound shocks, which are formed by an impulsive perturbation, on the observed $\\sim 5$-min oscillations in the solar corona. Here we consider the simplest hydrodynamic case, which can be developed to more realistic magnetohydrodynamic model in future studies. This paper is organized as follows. The basic equations and the atmospheric model are described in Sect.~\\ref{sect:num_model}. The numerical model and results of numerical simulations of impulsive photospheric driver are discussed in Sect.~\\ref{sec:num_res}. This paper is concluded by a summary of the main results in Sect.~\\ref{sec:sum}. ", "conclusions": "\\label{sec:sum} Frequently observed $\\sim 5$-min oscillations in the solar corona are often explained by leakage of photospheric p-modes along inclined magnetic field. Vertically propagating acoustic waves with 5-min period are evanescent due to the stratification of the solar atmosphere as chromospheric acoustic cut-off period is $\\sim$ 3 min (Roberts \\cite{roberts}). But, acoustic-gravity waves have a smaller cut-off frequency when they propagate with the angle about the vertical. This may allow the photospheric 5-min oscillations to channel along a stratified chromosphere and penetrate into the corona (De Pontieu et al. \\cite{dep05}, Erd{\\'e}lyi et al. \\cite{erd2007}, Fedun et al. \\cite{fedun09}). In order to increase the cut-off period from 3-min up to 5-min, the propagation angle (or magnetic field inclination) should be $\\theta$ $\\sim$ 50$^0$ ($\\cos\\theta \\approx$ 3/5). Therefore, the leakage of p-modes may take place only in particular regions of the solar atmosphere, where the magnetic field is significantly inclined in the chromosphere. In this paper, we suggest an alternative mechanism to explain the observed oscillations in the solar corona, which is based on the rebound shock model of Hollweg (\\cite{hol82}). We numerically solved the full set of Euler equations for the realistic VAL-C temperature profile and for a Gaussian velocity pulse launched within the photosphere. We found that velocity pulses, originating from granules or magnetic reconnection in the lower regions, lead to different responses of the chromosphere/transition region than the periodic acoustic waves resulting from p-modes. It must be mentioned, that the energy of granular motions is higher than the energy of p-modes, therefore the impulsively triggered waves should have more power than those triggered by a periodic driver. The numerical simulations show that as a result of the rapid decrease of the equilibrium mass density the initial velocity pulse quickly steepens into a shock. The shock propagates into the corona, while the nonlinear wake is formed in the chromosphere due to the atmospheric stratification. This nonlinear wake leads to consecutive shocks as was first shown by Hollweg (\\cite{hol82}). The interval between the arrival times of two consecutive shocks depends on the amplitude of the initial pulse; a stronger pulse leads to longer intervals. The initial pulse with the granular velocity of $1$ km s$^{-1}$ leads to $\\sim$ $5$-min intervals between consecutive shocks. Therefore, the quasi-periodic arrival of consecutive shocks in the solar corona may cause the intensity oscillations with a period close to the interval between the shocks. We implemented a simple 1D analytical model in order to avoid the propagation of acoustic oscillations with the angle to the vertical. We showed that purely vertically propagating pulses may lead to quasi $5$-min oscillations in the corona due to consecutive shocks. Therefore, it is not necessary to invoke the propagation along inclined magnetic field in order to explain the observed $5$-min periodicity in the corona. One dimensional propagation for acoustic waves is justified for purely vertical magnetic field. In this case, acoustic waves are in fact slow magneto-acoustic waves for low plasma $\\beta \\sim c_s^2/v_A^2 <1$ and fast magneto-acoustic waves for high plasma $\\beta >1$. Here $v_A$ is the Alfv\\'en speed. In the solar photosphere $\\beta$ is larger than unity, but it rapidly decreases due to the mass density fall off with height (and consequent increase of the Alfv\\'en speed). It becomes smaller than unity in the chromosphere (Gary \\cite{gary2001}) and tends to unity somewhere between the photosphere and the chromosphere and this surface should be thinner than the width of the chromosphere. The linear fast and slow magneto-acoustic waves are coupled near the level of $\\beta \\sim 1$ when they propagate obliquely to the magnetic field (Rosenthal et al. \\cite{rosenthal}, Bogdan et al. \\cite{bog2003}). However, these waves remain purely acoustic for parallel propagation unless the tube dispersive effects are taken into account. Therefore, a pulse propagating along the vertical magnetic field may not feel the $\\beta \\sim 1$ surface. On the other hand, the acoustic wake, which is formed behind the pulse and oscillates along the magnetic field, may lead to the non-linear energy transfer into Alfv\\'en waves near the $\\beta \\sim 1$ region (Zaqarashvili and Roberts \\cite{zaqarashvili}, Kuridze and Zaqarashvili \\cite{kuridze08}). This effect may be revealed as far as one includes the magnetic field into numerical simulations. Then, a part of oscillation energy may be transformed into transverse oscillations near this region, but most of energy will remain in longitudinal oscillations due to thinness of $\\beta \\sim 1$ region. Therefore, $\\beta \\sim 1$ region may not significantly affect the formation of rebound shocks in the chromosphere and consequently quasi-periodic acoustic oscillations in the lower corona. However, a two-dimensional consideration and inclusion of magnetic field are necessary for the complete understanding of the proposed scenario. The first step in this direction was already done by Murawski and Zaqarashvili (\\cite{murawski2010}) in the modeling of spicule formation, but they considered a simple temperature profile covering only the chromosphere-corona. The plasma $\\beta$ was considered less than unity along the whole simulation region, therefore the effects of the $\\beta \\sim 1$ region were absent in these simulations. We intend to consider the rebound shock model in the case of magnetic field and realistic temperature profile in the future. An important consequence of the rebound shock model is that the interval between consecutive shocks depends on the initial amplitude of pulse (see Fig. 2). The smaller amplitude pulses lead to shorter interval between consecutive shocks with lower limit of 3-min, which is the linear acoustic cut-off period of the solar atmosphere. The interval is longer for stronger initial pulses being $\\sim$ 5-min for $1$ km s$^{-1}$ in the photosphere. Then the rebound shock model predicts the longer period oscillations above the regions of strong granular power. The granulation is suppressed in strong magnetic field regions of the photosphere (e.g. sunspots), therefore the initial pulses should have smaller amplitudes there, leading to the oscillation at near cut-off frequency. Indeed, the strong magnetic field regions (sunspots, magnetic network cores) show predominantly 3-min oscillations in the chromosphere. Note that some observations also show the 3-min oscillations in the solar corona above sunspots (De Moortel et al. \\cite{de Moortel02}) and other magnetic structures (Lin et al. \\cite{lin2005}). These oscillations can be excited as a consequence of consecutive shocks due to chromospheric 3-min oscillations. As a result, we expect that $\\sim 5$-min oscillations can be detected above the regions where the granular energy is significant i.e. the quiet Sun and surroundings of active regions/magnetic network. This is consistent with observations. On the other hand, the detection of 5-min oscillations in the quiet corona is not an easy task due to a relatively weak intensity of coronal lines. However, a careful analysis still can be performed. But, one should keep in mind that the dynamics of acoustic waves in the field-free regions and along the vertical magnetic field could be quite different as it was discussed above. An initial acoustic pulse may be spread horizontally in field-free regions, while almost the whole energy would be guided along the field lines in magnetic structures. Therefore, the amplitude of intensity oscillations could be smaller in the quiet corona than near active regions and chromospheric network cores. However, the divergence of magnetic field and increased thermal conduction may significantly weaken the amplitude of coronal oscillations, which seem to be quite strong and non-linear on Figs. 1 and 2. Then, the strong slow wave pulses may become almost linear once they penetrate into the corona as it is seen by observations. Our numerical model then requires the inclusion of magnetic field and the thermal conduction (at least, in the coronal part of the atmosphere). This will be done in future studies. It should be noted that the granular velocities may take values between 0.5-2 km s$^{-1}$ with peak on 1 km s$^{-1}$. As the interval between consecutive shocks strongly depends on the amplitude of the initial pulse, then the resulted coronal oscillations may take values between 4-7 min with a peak on 5-min. Indeed, the wavelet analysis of coronal line images obtained by Hinode/EIS show that the oscillation power of coronal oscillations is concentrated at the period in a range of 4-6 min (Wang et al. \\cite{wang09}), which is fully consistent with our theory. Our simulations were performed for an isolated pulse in order to show clearly the effect of rebound shocks. However, the solar photosphere is very dynamic, hence the initial pulse probably is followed by other pulses. The subsequent pulse coming from the photosphere may also trigger the consecutive shocks in the chromosphere. The interaction between rebound shocks formed by different photospheric pulses may set up complex dynamics in the chromospheric plasma with immediate influence on the lower corona. Therefore, the coronal oscillations may have broad spectrum as we already discussed in the previous paragraph. On the other hand, if the mean interval between subsequent initial pulses is close to the mean interval between consecutive shocks, then a resonance may occur in the chromosphere. This could be subject of future study. The excitation of coronal oscillations due to leakage of p-modes may occur only along significantly inclined magnetic field in the chromosphere (in order to reduce the cut-off frequency). The magnetic field lines, which are significantly inclined from the vertical in the chromosphere, may not reach the corona at all. Therefore, the p-mode leakage may have problems in real geometry of active region magnetic field. On the contrary, the rebound shock mechanism may work in any geometry of magnetic field including the purely vertical field lines. Therefore, the excitation of 5-min oscillations in the solar corona by photospheric impulsive drivers has larger area of application than that of by p-modes. We believe that the future sophisticated models may shed light on the excitation of coronal acoustic waves. Our conclusions are: \\begin{itemize} \\item [(a)] a velocity pulse that is initially launched at the photospheric level (due to granules or reconnection) quickly steepens into a shock and can penetrate into the corona, while a nonlinear wake that is formed behind this shock leads to consecutive shocks in the chromosphere; \\item [(b)] for the initial photospheric pulse amplitude of $1$ km s$^{-1}$ the time interval between two consecutive shocks is $\\sim 5$-min; the consecutive shocks propagate upwards and may cause the observed $5$-min intensity oscillations in the solar corona; \\item [(c)] the final conclusion is that the observed $\\sim$ 5-min oscillations in the solar corona could be caused by impulsive photospheric perturbations (convection, reconnection) not necessarily by p-modes. \\end{itemize} {\\it Acknowledgements:} The authors express their thanks to the unknown referee for his/her constructive comments. The work of TZ and MK was supported by the Austrian Fond zur F\\\"orderung der Wissenschaftlichen Forschung (project P21197-N16). The work of KM was supported by the Polish Ministry of Science (the grant for years 2007-2010). TZ was also supported by the Georgian National Science Foundation grant GNSF/ST09/4-310. The software used in this work was in part developed by the DOE-supported ASC / Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago." }, "1101/1101.1060_arXiv.txt": { "abstract": "Large-amplitude Alfv\\'en waves are ubiquitous in space plasmas and a main component of magnetohydrodynamic (MHD) turbulence in the heliosphere. As pump waves, they are prone to parametric instability by which they can generate cyclotron and acoustic daughter waves. Here, we revisit a related process within the framework of the multi-fluid equations for a plasma consisting of many species. The nonlinear coupling of the Alfv\\'en wave to acoustic waves is studied, and a set of compressive and coupled wave equations for the transverse magnetic field and longitudinal electric field is derived for waves propagating along the mean-field direction. It turns out that slightly compressive Alfv\\'en waves exert, through induced gyro-radius and kinetic-energy modulations, an electromotive force on the particles in association with a longitudinal electric field, which has a potential that is given by the gradient of the transverse kinetic energy of the particles gyrating about the mean field. This in turn drives electric fluctuations (sound and ion-acoustic waves) along the mean magnetic field, which can nonlinearly react back on the transverse magnetic field. Mutually coupled Alfv\\'en-cyclotron--acoustic waves are thus excited, a nonlinear process that can drive a cascade of wave energy in the plasma, and may generate compressive microturbulence. These driven electric fluctuations might have consequences for the dissipation of MHD turbulence and, thus, for the heating and acceleration of particles in the solar wind. ", "introduction": "Large-amplitude Alfv\\'en waves are ubiquitous in space plasmas, and particularly prominent in the solar wind \\citep{tu95, bruno05}. They are an essential component of magnetohydrodynamic (MHD) turbulence in the heliosphere and known to originate mainly in the solar coronal holes \\citep{cranmer09}. As has been shown in the ample literature, an Alfv\\'en mother (pump) wave is prone to parametric instability \\citep{stenflo76, derby78, goldstein78, longtin86, brodin88, hollweg94, wong86, vinas91b, vinas91a, stenflo00, ruderman04, stenflo07} by which it can generate cyclotron and acoustic daughter waves that may undergo kinetic effects \\citep{araneda98} and collisionless Landau damping \\citep{inhester90, araneda07}. The continuous and wide interest in these waves also comes from their astounding properties, namely that Alfve\\'n-cyclotron waves, like parallel magnetosonic-whistler waves, are nonlinear eigenmodes \\citep{sonnerup67, stenflo76} of the MHD, and multi-fluid equations as shown below, for propagation along the mean magnetic field. Nonlinearly excited \\citep{spangler89} acoustic waves appear to be common in space plasmas as well, and density fluctuations \\citep{tu95, bruno05} are observed everywhere in the solar wind, although at a comparatively low fluctuation level of merely a few percent. However, since compressive fluctuations can be damped through kinetic effects, like Landau damping on the thermal ions and electrons, they can provide an effective dissipation mechanism for the nonlinear damping \\citep{medvedev97} of Alfv\\'en waves. Consequently, the understanding of the coupling between Alfv\\'enic wave activity and density or charge-density fluctuations is of paramount interest and importance in basic plasma physics, but alike in its applications to nonlinear processes in space \\citep{stenflo07} and astrophysical plasmas. As will be shown in this paper, a coupled set of nonlinear second-order wave equations for the transverse magnetic field, the transverse gyromotion of any particle species in the multicomponent plasma considered, and the related longitudinal electric field can be derived, which together describe the wave--wave interactions and their mutual forcing. These equations provide a physically and intuitively clear picture of the field and particle/plasma dynamics and allow us to understand the results of recent hybrid simulations of the parametric decay of Alfv\\'en waves and their effects on the plasma particles better. The main aim of this work is to provide algebraic derivations and physical explanations. A numerical treatment of the full equations to be derived subsequently appears promising, yet is beyond the scope of this work. In analytical \\citep{araneda07}, hybrid-simulation, and other numerical simulation \\citep{araneda08} studies of the parametric instabilities of Alfve\\'n-cyclotron waves, it became obvious that ion trapping \\citep{araneda08, araneda09} in the nonlinearly driven ion-acoustic waves and pitch-angle scattering by the transverse daughter waves were found to cause anisotropic heating of the proton core velocity distribution, and simultaneously to create a proton beam along the mean field \\citep{araneda08,valentini09}. These numerical results are in close agreement with observed kinetic features in the solar wind and support the observation that pitch-angle scattering \\citep{heuer07, marsch01} is the key to understand the kinetic characteristics of thermal solar-wind protons. But only recently convincing evidence has been found for ion-cyclotron waves \\citep{jian09} to exist in the solar wind. Also simulations of electric field spectra \\citep{valentini08} have shown that the short-scale termination of solar wind turbulence is characterized by the occurrence of longitudinal electrostatic fluctuations. The spectra thus obtained seem to be consistent with the electrostatic waves actually measured in the solar wind \\citep{bale05} close to the Earth's bow shock, and in particular in the ion-cyclotron range \\citep{kellogg06}. The present study will provide the foundation for insight into and further study of the processes occurring at macroscopic and microscopic scales in solar-wind turbulence, and thus will throw light on the related dissipation processes through kinetic cascades and wave--particle interactions \\citep{marsch06}. The nonlinear equations derived here are used to describe an elliptically polarized Alfv\\'en wave as a simple but nontrivial example of their application. ", "conclusions": "Starting from the multi-fluid equations of a warm plasma, we have derived and investigated the coupled wave equations for the particles' gyromotions about the mean field and for the transverse magnetic field and longitudinal electric field. It is a natural outcome of the electromotive forces arising from compressible Alfv\\'en-cyclotron waves and can be derived from a potential that is just the kinetic energy associated with the gyromotion in the electromagnetic wave. Electric waves are thus excited, which can react back on the pump wave by nonlinear effects through terms in its own wave equation that contains the electric field explicitly. Known limiting cases are reproduced, such as the standard linear electric waves like the ion-acoustic or Langmuir waves of course, and for the transverse magnetic field the usual two branches of Alfv\\'en-cyclotron and magnetosonic-whistler waves in case of a two-component electron-proton plasma, or many similar related branches in the case of a multi-ion plasma. The main result of this paper is the closed set of second-order wave equations (\\ref{eq.76}), (\\ref{eq.79a}), and (\\ref{eq.79b}), from solutions of which the transverse electric field and charge densities of each species can be derived as auxiliary quantities. To study these wave equations in more detail and to find their nonlinear solutions is left as a future task, which will presumably require a numerical treatment. The structure of our equations already permits to derive some qualitative conclusions and to treat some simple applications (like the effect of elliptical polarization) analytically. Further study is certainly required to corroborate them quantitatively. Apparently, the weakly compressible large-amplitude Alfv\\'en-cyclotron waves can drive electric fluctuations, essentially of the ion-acoustic type, along the mean field, and thus will naturally produce an electric field that can accelerate particles and will lead to heating via Landau damping in a kinetic Vlasov description. By excitation of acoustic waves, the amplitude of the driver wave will be diminished until a dynamic wave--wave equilibrium is reached. Similar processes are clearly found in the direct numerical simulations \\citep{araneda08, araneda09, valentini08,valentini09}. The third-order coupling terms in Eqs.~(\\ref{eq.76}) and (\\ref{eq.79b}) correspond to such three-wave processes in Fourier space, and therefore will lead to cascading of spectral energy and broadening of the original spectrum of the pump wave, which need not be monochromatic. This way a new path towards micro- and macro-turbulence could be opened, and a non-MHD cascade is rendered possible by these compressive Alfv\\'en-cyclotron--acoustic wave interactions. \\begin{acknowledgment} D.~V. appreciates financial support by the International Max Planck Research School (IMPRS) on Physical Processes in the Solar System and Beyond. \\end{acknowledgment}" }, "1101/1101.3779_arXiv.txt": { "abstract": "We present new high signal$-$to$-$noise spectroscopic data on the M31 globular cluster (GC) system, obtained with the Hectospec multifiber spectrograph on the 6.5m MMT. More than 300 clusters have been observed at a resolution of 5\\AA \\ and with a median S/N of 75 per \\AA , providing velocities with a median uncertainty of 6 \\kms. The primary focus of this paper is the determination of mean cluster metallicities, ages and reddenings. Metallicities were estimated using a calibration of Lick indices with [Fe/H] provided by Galactic GCs. These match well the metallicities of 24 M31 clusters determined from HST color$-$magnitude diagrams, the differences having an rms of 0.2 dex. The metallicity distribution is not generally bimodal, in strong distinction with the bimodal Galactic globular distribution. Rather, the M31 distribution shows a broad peak, centered at [Fe/H]=$-1$, possibly with minor peaks at [Fe/H]=$-1.4$, $-0.7$ and $-0.2$, suggesting that the cluster systems of M31 and the Milky Way had different formation histories. Ages for clusters with [Fe/H] $> -1$ were determined using the automatic stellar population analysis program {\\it EZ\\_Ages}. We find no evidence for massive clusters in M31 with intermediate ages, those between 2 and 6 Gyr. Moreover, we find that the mean ages of the old GCs are remarkably constant over about a decade in metallicity ($-0.95 \\simless $ [Fe/H]$ \\simless 0.0$). ", "introduction": "Star clusters in the Andromeda galaxy have been studied since \\citet[][]{hubble}, and have long been known to comprise ages from young to old. It is the latter which are the topic of this paper. Initially, cataloging was all that was possible for clusters, work which indeed continues to this day with searches for distant clusters \\citep{huxor}. But spectroscopic studies as a means of measuring stellar populations and group kinematics also began early, starting with \\cite{vdb}, with substantial contributions by \\cite{huchra82}, \\cite{burstein}, \\cite{trip}, \\cite{huchra}, and \\cite{federici}. \\cite{barmby} added much spectroscopy and photometry to the sample, and presented the largest study before the use of multiobject spectrographs on M31. At the time of the Barmby et al (2000) work, it was generally thought that the old M31 globular clusters (GCs) numbered several hundred, spanned a range of metallicity (as deduced from colors and absorption line strengths) equal to that of the galactic GCs, and likewise did not have a simple single Gaussian distribution of metallicities. The metal$-$poor clusters were thought to be roughly spherically distributed and to show a small amount of systemic rotation, while the most metal$-$rich clusters were thought to be confined to the disk of M31, and to show more systemic rotation, although the specifics of those and subsequent results will be further refined in this paper and other papers in this series. Two large multifiber studies of M31 clusters have been presented in the last decade: \\cite{perrett} who used WYFFOS on the WHT, and \\cite{kim} who used Hydra on the WIYN telescope. The former paper produced more than 200 new velocities, and also found a bimodality of the old cluster metallicities, using spectral indices rather than colors. A subsequent kinematic analysis of those data by \\cite{morrison} suggested that the globulars could be explained using two kinematic components: a thin, cold rotating disk and a higher velocity dispersion component whose properties resemble M31's bulge. The second study \\citep[where the kinematic results were presented in][]{lee} provided spectra for an additional 150 objects. However, in both studies, the kinematic analysis suffered from the inclusion of young disk clusters, which in particular led to statements that there were metal$-$poor clusters with thin disk kinematics \\citep{morrison} or that the metal$-$poor clusters showed strong systemic rotation \\citep{lee}, in conflict with what \\cite{huchra} had reported, and what we will also conclude in a subsequent paper (A. Romanowsky et al., in preparation). The important task of distinguishing M31's young clusters from old was one of the topics in \\citet[Paper I]{PaperI}, made possible with new spectroscopy taken on the 6.5m MMT and using the Hectospec multi$-$fiber spectrograph \\citep{fab}. Ages and masses for more than 140 young clusters were determined by comparison with models, finding clusters with masses as great as $10^4 M_\\sun$, and a median cluster age of 0.25 Gyr. Table 1 of that paper also listed all the clusters, regardless of age, for which we added new information to the long$-$standing cataloging efforts by \\cite{galleti2}. That new information included revised coordinates, magnitudes, reddenings, and cluster classifications based on images and spectroscopy. With the distinction between the young and old clusters now better clarified, we present here the first results of our high signal$-$to$-$noise (S/N) spectroscopic study of the old clusters in M31 (where we define old to be those with ages greater than 6 Gyr). This paper presents the velocities, ages and metallicities for the old clusters, and discusses the spatial, abundance and age properties of these clusters. The improved data allow us to revisit the topics addressed by previous authors. Where possible, we present comparisons with metal abundances derived from cluster color$-$magnitude diagrams (CMDs) using {\\it HST} imaging, some of which is presented here. Subsequent papers will discuss kinematics and abundance ratios of these clusters. We assume a distance of 770 kpc throughout \\citep{freedman}. ", "conclusions": "\\label{s:ages} We now discuss the issue of age variations among the M31 globular clusters, comparing the results found using EZ\\_Ages and the statements of previous authors that a number of M31 clusters have intermediate ages. Based on archival spectra, \\cite{beasley} stated that, in particular, B158-G213 and B337-G068 had similar metallicities but different ages. They also stated that five M31 clusters were intermediate in age, with ages between 2 and 5Gyr: B126-G184, B301-G022, NB16, NB67 and also, B337-G068. \\cite{brodie} compared the spectrum of NB67 with two other M31 clusters and concluded that it was an intermediate age cluster. Our response to those claims follows. From the Hectospec data in Table \\ref{main}, we found that (a) B158-G213 and B337-G068 are old and have dissimilar metallicities ($-0.8$ and $-1.2$, respectively), (b) B126-G184, B301-G022, NB16, and B337-G068 are all older than 9 Gyr, with [Fe/H] values ranging from $-0.8$ to $-1.5$. As it turns out, NB67 is a foreground F star \\citep{PaperI}. Using at the time fresh data, \\cite{burstein2} additionally claimed that B232-G286 and B311-G033 had ages of 5 Gyr. We found that both of these clusters are old, and simply very metal$-$poor ( [Fe/H] = $-2.0$ and $-1.9$, respectively). In general, previous authors have mistaken lower metallicity clusters for younger ones. To demonstrate our claim graphically, Figure \\ref{notyoung} highlights the six purportedly intermediate age clusters in the $\\langle$Fe$\\rangle$$-$H$\\beta $ index diagram, similar to the top panel of Figure \\ref{hbhdfem}. If these clusters were substantially younger than the mean cluster age at any given metallicity (as measured by $\\langle$Fe$\\rangle$), we would expect them to have H$\\beta $ indices stronger than the mean H$\\beta $ index. Such is not the case. \\begin{figure} \\vspace{1.0cm} \\plotone{f20.eps} \\caption{$\\langle$Fe$\\rangle$ (average of Lick indices Fe5270 and Fe5335) plotted against H$\\beta$. All of the old clusters are plotted in small, gray filled circles. The six clusters previously reported to be intermediate age are shown as large black circles, with their names attached. None of these shows any evidence for enhanced H$\\beta$ strength, and hence none has evidence for an intermediate age. Average error bars for all the clusters are again shown in the bottom left corner. \\label{notyoung}} \\end{figure} A series of papers using the BATC photometric system combined with other photometry has produced several tables of ages for clusters, young and old \\citep{fan2, ma2, wang}. The ages were derived using SSP models, based on Padova isochrones. There is very little correlation of the ages in those papers and those we have reported here and in Paper I. The main source of the discrepancy is that many of the clusters identified by the cited papers as being young are in fact old and metal$-$poor. Of the 77 clusters in common with \\cite{wang}, 50 clusters stated to be younger than 5 Gyr are older than 10 Gyr based on our analysis. Returning to our own age determinations, we note again that from the EZ\\_Ages analysis, some clusters marked as old in Paper I were realized to be younger than 2 Gyr after publication of that paper. These are all disk clusters; Table \\ref{revised} lists these clusters with their revised ages, all under 2 Gyr, as derived from the method of Paper I and confirmed by EZ\\_Ages. Aside of those, there are a small number of clusters (12) with ages younger than 8 and older than 2 Gyr, but six of these have abundances close to the problem [Fe/H] value of $-0.95$ and whose ages are thus suspect (see above). Thus only six are worth further consideration with regard to intermediate ages. These are B015, B071, B138, B140, B268, and AU010. Their ages are all around 7 Gyr, and all but B015 are within 2 kpc of M31's center. These have masses between $10^5$ and $4 \\times 10^5 M_\\sun$, close to the median mass for all the M31 GCs. Interestingly, they are all metal$-$rich (five out of the six have [Fe/H] $> -0.2$, and have very strong CN bands. \\begin{figure} \\vspace{1.0cm} \\plotone{f21.eps} \\caption{Histogram of ages. Ages for clusters older than 1 Gyr were determined via EZ\\_Ages, ages for younger clusters come from \\cite{PaperI}. Ages for old clusters with [Fe/H]$<-0.95$ were set to 14 (1.15 in the log). \\label{age_hist}} \\end{figure} To conclude the age discussion, we find no evidence for any massive clusters in M31 with intermediate ages, those between 2 and 6 Gyr. Figure \\ref{age_hist} shows the age histogram, including the young clusters whose ages were determined in Paper I. For this diagram, we assumed that all clusters with [Fe/H] $< -0.95$ (whether we determined the metallicity here or others did so elsewhere) have ages of 14 Gyr. We also required the clusters to have masses greater than $5 \\times 10^3 M_\\sun$. This diagram clearly shows the gap in ages between 2 and 6 Gyr. Moreover, we have found that the mean ages of the old GCs in M31 seem to be remarkably constant over about a decade in metallicity ($-0.95 \\simless $ [Fe/H]$ \\simless 0.0$). Using the high$-$quality spectra reported here, we have provided new homogeneous estimates of the metallicities and ages for more than 300 globular clusters in M31. Again, within a radius 21 kpc we observed 94\\% of the old clusters. We note that only 13\\% of MW GCs lie beyond that limit of 21 kpc. The search for outer M31 clusters, presumably metal$-$poor, is still underway \\cite[e.g.,][]{huxor}, but it is unlikely that enough clusters will be discovered to change our basic comparison of metal$-$rich and metal$-$poor clusters in M31 We find no evidence for any massive clusters in M31 with intermediate ages, those between 2 and 6 Gyr. The metallicities span the range of those found in the MW, with a few clusters perhaps extending beyond the most metal$-$rich galactic globulars. The M31 cluster metallicity distribution is quite different from that of the MW, not showing a strong bimodality as does the MW. However, there are hints of multi$-$modality. Since it is likely that GCs were deposited in M31 in a hierarchical fashion from merging and accretion of a number of smaller galaxies, the lack of simple bimodality in the M31 cluster [Fe/H] distribution is perhaps not surprising. Our new data confirm about 320 true GCs. To that we may add another 50 clusters which we did not observe, but which have been confirmed as globulars by various sources, but most significantly by \\cite{barmby} and \\cite{huxor}. Thus the total number of known M31 GCs is about 370, compared with 150 for the MW \\citep{harris}. We close this paper with some other comparisons between the two galaxies. Roughly, in M31 the metal$-$poor and metal$-$rich groups have equal numbers of clusters. There are 333 clusters with [Fe/H] measured here or in the literature. If we divide them at [Fe/H]=$-1$, there are 160 above that with a combined mass of $1 \\times 10^8 M_\\sun$ and 173 below it with a mass of $8 \\times 10^7 M_\\sun$. If we place all the clusters with unmeasured [Fe/H] into the metal$-$poor group (there are 39 of these, mostly in the outer halo), the result changes only slightly. There would be 218 clusters in the metal$-$poor group, but the total mass increases only to $9 \\times 10^7 M_\\sun$ since most of the added clusters are low mass. In the MW, that same dividing point in [Fe/H] divides the clusters into 1/3 metal$-$rich and 2/3 metal$-$poor, with the former having a total mass of $9 \\times 10^6 M_\\sun$ and the latter with $3 \\times 10^7 M_\\sun$. The relative total masses of M31 and the MW are a topic of current interest \\citep{reid}, but we can limit the discussion to the bulge masses, the ratio of which is about a factor of 2 \\citep{vdk}. Therefore, the ratio of number of GCs of all metallicities to the parent's bulge mass does appear to be similar in the two galaxies. However, the ratio of the numbers of metal$-$rich GCs to bulge mass, or equivalently, the ratio of metal$-$poor clusters to the bulge mass is not similar, due implicitly to the different ratio of metal$-$rich to metal$-$poor clusters between M31 and the MW. The specific frequency of GCs is defined as the ratio of the number of clusters of all metallicity to the galaxy luminosity in units of $M_{\\rm V}=-15$. Since M31 has about 370 GCs, and has $M_{\\rm V}=-21.2$ \\citep{vdb2}, we then find a specific frequency of 1.2. Isolating the bulge luminosity (taken to be 30\\% of the total light), we derive a bulge specific frequency of 4. These numbers can be compared with the values of 1 and 2 for the MW (the latter for the bulge specific frequency. Thus, M31 has more clusters per unit luminosity of old stars. In M31, 50\\% of the known GCs lie within 5.1 kpc of the Galactic center, remarkably close to the half$-${\\it total number} radius of 4.8 kpc for the MW. (We calculated the projected radius of the MW GCs in the $YZ$ plane). However, the half$-${\\it mass} radius of the M31 GC system is smaller, 3.7 kpc, because the more distant clusters are less massive in the mean than the inner ones. The half$-${\\it mass} radius for the MW GC system is about 5.5 kpc, slightly larger than the half$-${\\it total number} radius. Succeeding papers in this series will concentrate on the abundance ratios of the clusters, the relation of kinematics and abundances in the cluster system, the horizontal branch morphologies, and the $M/L$ ratios of the clusters as derived from high dispersion spectroscopy and {\\it HST} imaging." }, "1101/1101.5255_arXiv.txt": { "abstract": "We perform a cosmological-model-independent test for the distance-duality (DD) relation $\\eta(z)=D_L(z)(1+z)^{-2}/D_A(z)$, where $D_L$ and $D_A$ are the luminosity distance and angular diameter distance respectively, with a combination of observational data for $D_L$ taken from the latest Union2 SNe Ia and that for $D_A$ provided by two galaxy clusters samples compiled by De Filippis {\\it et al.} and Bonamente {\\it et al.}. Two parameterizations for $\\eta(z)$, i.e., $\\eta(z)=1+\\eta_0z$ and $\\eta(z)=1+\\eta_0z/(1+z)$, are used. We find that the DD relation can be accommodated at $1\\sigma$ confidence level (CL) for the De Filippis {\\it et al.} sample and at $3\\sigma$ CL for the Bonamente {\\it et al.} sample. We also examine the DD relation by postulating two more general parameterizations: $\\eta(z)=\\eta_0+\\eta_1z$ and $\\eta(z)=\\eta_0+\\eta_1z/(1+z)$, and find that the DD relation is compatible with the results from the De Filippis {\\it et al.} and the Bonamente {\\it et al.} samples at $1\\sigma$ and $2\\sigma$ CLs, respectively. Thus, we conclude that the DD relation is compatible with present observations. ", "introduction": "The distance-duality (DD) relation~\\citep{Etherington} between the luminosity distance $D_L$ and the angular diameter distance (ADD) $D_A$, i.e., \\begin{equation} \\frac{D_L}{D_A}(1+z)^{-2}=1, \\end{equation}where $z$ is the redshift, plays an important role in modern observational cosmology~\\citep{Schneider, Cunha, Mantz, Komatsu}, and, actually, it has heretofore been applied to all analysis of the cosmological observations without any doubt. However, in reality, it is possible that one of the requirements in obtaining the DD relation may be violated. A violation of the DD relation may even be considered as a signal of the breakdown of physics on which the DD relation is based upon \\citep{Csaki, Bassetta,Bassettb}). Thus, it is desirable to perform a validity check on the DD relation by the astronomical observations. In this regard, \\citet{Uzan} have tested it by using the observations from the Sunyaev-Zeldovich effect (SZE) and X-ray surface brightness from galaxy clusters, and found that the DD relation is consistent with observations at $1\\sigma$ confidence level (CL). With a different galaxy clusters sample provided by \\citet{Bonamente}, \\citet{Bernardis} also obtained a non-violation of the DD relation. In addition, by combining the Union Type Ia supernovae (SNe Ia) \\citep{Kowalski} with the latest measurement of the Hubble expansion at redshifts between $0$ and $2$ \\citep{Stern}, \\citet{Avgoustidis} discussed this relation and obtained that it is consistent with observations at $2\\sigma$ CL. Recently, by assuming that the DD relation satisfies the following expression \\begin{equation} \\frac{D_L}{D_A}(1+z)^{-2}=\\eta(z), \\end{equation} where $\\eta(z)$ is parameterized as $\\eta(z)=1+\\eta_0z$ and $\\eta(z)=1+\\eta_0z/(1+z)$, \\citet{Holandaa} discussed the validity of the DD relation with the ADD $D_A$ measurements from galaxy clusters provided by the \\citet{Filippis} (elliptical $\\beta$ model) sample and the \\citet{Bonamente} (spherical $\\beta$ model) sample, and the luminosity distance $D_L$ given in the context of $\\Lambda$CDM. Here, the elliptical and spherical $\\beta$ models are two different geometries used to describe the galaxy clusters. Their results showed that the elliptical model is more compatible with no violation of the DD relation. However, all the aforementioned analyses are model dependent since a cosmic concordance model ($\\Lambda$CDM) is assumed in their discussions. It is worth noting that \\citet{Bernardis} have in fact tried to test the DD relation in a model-independent way. In their method, the ADD is given from galaxy clusters and the luminosity distance is from SNe Ia. To obtain the values of the ADD and the luminosity distances at the same redshift, \\citet{Bernardis} binned their data and found that the DD relation is not violated at $1\\sigma$ CL. However, when they determine the ADD from observations, the relation $D_A^{cluster}(z)=D_A(z)$ is used, which holds under the condition with no violation of the DD relation. Recently, a consistent cosmological-model-independent test for the reciprocity relation was proposed by \\citet{Holandab}. The main idea is to test the DD relation directly with the observed luminosity and ADDs, provided by SNe Ia and galaxy clusters samples, respectively. They considered two specific different redshift-dependent parameterizations for $\\eta(z)$: $\\eta(z)=1+\\eta_0z$ and $\\eta(z)=1+\\eta_0z/(1+z)$. The data sets used were given from the Constitution SNe Ia \\citep{Hicken} and two ADD samples \\citep{Filippis, Bonamente} from galaxy clusters obtained through SZE effect and X-ray measurement with different geometry descriptions to the cluster: the elliptical $\\beta$ model \\citep{Filippis} and the spherical $\\beta$ model \\citep{Bonamente}. They found that the result from the elliptical model is consistent with the DD relation at $2\\sigma$ CL, while for the spherical model the relation is clearly incompatible with observations. However, we find that six and twelve ADD data points should be removed, respectively, for the De Filippis et al. and Bonamente et al. samples, instead of only three ADD data points that were discarded for both samples in {Holanda} {\\it et al.} (2010b). So, in the present Letter, we first redo the same analysis as \\citet{Holandab} but with more data points removed to see how this would affect the result, and we obtain a more serious violation of the DD relation. We then consider the effect of the errors of SNe Ia which were neglected in \\citet{Holandab} and give a comparison of the results obtained with and without the errors of SNe Ia. More importantly, we retest the DD relation using the latest Union2 SN Ia data \\citep{Amanullah}. Compared with the Constitution set used by \\citet{Holandab}, the Union2 has the following advantages: (1) the selection criteria ($\\Delta z <0.005$) can be satisfied for all data points of two ADD samples except for the cluster CL J1226.9+3332 ($z=0.890$), which corresponds to $\\Delta z =0.005$, from the Bonamente et al. (2006) sample, (2) the values of $z_{SNe~Ia}-z_{Cluster}$ are more centered around the $\\Delta z=0$ line. So, the results from Union2 may be more reliable. Finally, we test the DD relation by assuming two more general parameterizations: $\\eta(z)=\\eta_0+\\eta_1z$ and $\\eta(z)=\\eta_0+\\eta_1z/(1+z)$, and find that in this case the consistencies between the observations and the DD relation are improved markedly for both samples of galaxy clusters. ", "conclusions": "In this Letter, we test the DD relation by considering the ADDs given by two samples of galaxy clusters together with the luminosity distances provided by sub-samples of SNe Ia picked from the Constitution and the latest Union2 data sets. The Constitution sample has already been discussed by \\citet{Holandab}, and they found that for both ADD samples three data points should be removed with the selection criteria ($\\Delta z= \\left|z_{Cluster}-z_{SNe~Ia}\\right|<0.005$). However, we find that, with the same selection criteria, the data points that have to be removed are actually six and twelve respectively for the \\citet{Filippis} sample and the \\citet{Bonamente} sample. A re-analysis with more data points discarded suggests a violation of the DD relation stronger than that given in \\citet{Holandab}. In order to obtain a more reliable result, we investigate the DD relation by considering the latest Union2 SNe Ia. It is worthy to note that with the Union2 SNe Ia all ADD data can be retained and the differences of the redshifts between ADD from the galaxy cluster and the associated luminosity distance from SNe Ia are much smaller. Thus the accuracy of our test should be improved. Our results then show that the DD relation can be accommodated at $1\\sigma$ CL. for the elliptical $\\beta$ model \\citep{Filippis} and at $3\\sigma$ C. L. for the spherical $\\beta$ model \\citep{Bonamente}. Finally, we examine the DD relation by postulating two more general parameterization forms: $\\eta(z)=\\eta_0+\\eta_1z$ and $\\eta(z)=\\eta_0+\\eta_1z/(1+z)$, and we find that the consistencies between the observations and the DD relation are improved markedly for both samples of galaxy clusters. The DD relation is compatible with \\citet{Filippis} sample and \\citet{Bonamente} sample at $1\\sigma$ and $2\\sigma$ CL., respectively. Furthermore, with the inclusion of the errors of SNe Ia, the results become more consistent with the DD relation. Therefore, our results suggest that the DD relation is compatible with the observations. This differs from what is obtained by \\citet{Holandab}, where the results from the \\citet{Bonamente} sample give a clear violation of the DD relation." }, "1101/1101.2643_arXiv.txt": { "abstract": "We investigate the feasibility of representing a structured \\multidimensional\\ stellar atmosphere with a single \\onedimensional\\ average stratification for the purpose of spectral diagnosis of the atmosphere's average spectrum. In particular we construct four different \\onedimensional\\ stratifications from a single snapshot of a magneto-hydrodynamic simulation of solar convection: one by averaging its properties over surfaces of constant height, and three different ones by averaging over surfaces of constant optical depth at 500 nm. Using these models we calculate continuum, and atomic and molecular line intensities and their center-to-limb variations. From analysis of the emerging spectra we identify three main reasons why these average representations are inadequate for accurate determination of stellar atmospheric properties through spectroscopic analysis. These reasons are: non-linearity in the Planck function with temperature, which raises the average emergent intensity of an inhomogeneous atmosphere above that of an average-property atmosphere, even if their temperature-optical depth stratification is identical; non-linearities in molecular formation with temperature and density, which raise the abundance of molecules of an inhomogeneous atmosphere over that in a \\onedimensional\\ model with the same average properties; the anisotropy of convective motions, which strongly affects the center-to-limb variation of line-core intensities. We argue therefore that a one-dimensional atmospheric model that reproduces the mean spectrum of an inhomogeneous atmosphere necessarily does not reflect the average physical properties of that atmosphere, and are therefore inherently unreliable. ", "introduction": "\\label{sec:introduction} Stellar spectra contain rich information about the physical conditions of the stars from which they originate, but this information can only be extracted meaningfully with sufficiently realistic stellar model atmospheres and line formation theory. Traditionally, applicable models have consisted of \\onedimensional\\ plane-parallel stratifications in which the run of temperature with height is determined from flux conservation, the mixing length formalism to represent convective energy transport, and hydrostatic equilibrium. This homogeneous and static picture, however, is very much at odds with high-resolution observations of the surface of our closest star, the Sun, which appears structured even at the smallest observable spatial scales, and dynamic down to the shortest temporal scales detectable with current instrumentation, most particularly at wavelengths that coincide with spectral lines. Building \\onedimensional\\ atmospheric representations that account for this spatial structure and temporal variation is not straightforward, even if the intent of the model is to reproduce the average spectrum of such an atmosphere. First, the model's stratification must represent in some way the horizontally averaged thermodynamic properties of the \\threedimensional\\ atmosphere. Secondly, it should represent convective motions and resulting Doppler shifts and their effect on line formation, and finally, if it is to be physically self-consistent, it must allow a complex structure supported by (magneto-)hydrodynamic forces to be represented by hydrostatic equilibrium, supplemented perhaps by some form of turbulent pressure. These approximations necessarily involve free parameters, for instance in the form of a mixing length parameter, micro- and macro turbulent broadening, and collisional line broadening, to match line widths, all of which may be freely adjusted to accurately reproduce certain observables for a given model. As a result measurement and model are not independent, and this draws into question the uniqueness of the measurement process. Finally, adiabatic cooling resulting from rapid expansion in the upper layers of true convective models is difficult to represent adequately in \\onedimensional\\ average models. In particular, in metal-poor stellar atmospheres this cooling leads to substantially cooler upper layers than predicted by radiative equilibrium, whereas in the case of more solar like metallicities the effect is much less pronounced as the adiabatic cooling is compensated by radiative heating in the multitude of weak spectral lines \\citep{Asplund2005}. In the upper photosphere of such metal-poor stars the temperatures may be as much as several hundred to one thousand degrees lower than predicted by a \\onedimensional\\ radiative equilibrium model constructed with the same stellar parameters. Self-consistent theoretical \\onedimensional\\ models are not the only ones used for spectroscopic diagnostics. In the special case of modeling the solar atmosphere the availability of high resolution spectra and information on limb darkening, allow for an alternative to the constraint of flux conservation, namely a semi-empirical determination of the temperature stratification, with the successful model by \\citet{HOLMUL} perhaps the best known example. Semi-empirical models have no requirement for physical self-consistency, but as they are constructed to serve a specific set of observables, the question arises if they can equally reproduce others than those on which they are based. \\citet{Holweger+Heise+Kock1990} argue, that the answer to this question is affirmative when considering line strengths used for abundance determinations, because the strengths of different lines vary in similar fashion, depending mostly on the local temperature gradient, with lines being stronger in steep gradients, and weaker in shallow ones, despite differences in excitation and ionization. To better account for the inhomogeneous and dynamic nature of stellar atmospheres than possible with \\onedimensional\\ modeling, more realistic models have been introduced over the last two decades that solve the equations for (magneto-)hydrodynamic forces in a gravitationally stratified atmosphere, consistently with radiative transfer \\citep[see][for an overview and references therein]% {Stein+Nordlund+Asplund-lrsp}, eliminating most of the free parameters that plague \\onedimensional\\ atmospheric representations. Many sophisticated codes now exist for these simulations \\citep[e.g.,][]{Stein+Nordlund1998,Freytag+Steffen+Dorch2002,% Schaffenberger_etal2005,Voegler_etal2005,Abbett2007,Jacoutot+Kosovichev+Wray+Mansour2008,% Hayek_etal2010,Muthsam_etal2010}. These numerical simulations of solar magneto-convection have been highly successful in reproducing the morphology of granules \\citep{Stein+Nordlund1989,Stein+Nordlund2000}, the prediction of solar $p$-modes \\citep{Nordlund+Stein2001,Stein+Nordlund2001}, and in particular in the reproduction of the space- and time-averaged shapes of photospheric absorption lines \\citep{Asplund+others2000,Asplund+Nordlund+Trampedach+Stein2000}. They have, however, also stirred a controversy in the determination of solar abundances, prompting a significant downward revision of the abundances of oxygen \\citep{AllendePrieto+Lambert+Asplund2001,% Asplund+Grevesse+Sauval+AllendePrieto+Kiselman2004}, and carbon \\citep{AllendePrieto+Lambert+Asplund2002,Asplund+Grevesse+Sauval+AllendePrieto+Blomme2005} by almost a factor of two, in sharp contradiction to values determined from \\onedimensional\\ modeling, and more importantly, from helioseismology \\citep[see][for a recent overview]{Serenelli+Basu+Ferguson+Asplund2009}, eliciting the question if even these \\threedimensional\\ convection simulations need further refinement. With currently available computer resources it is, however, not yet practical to employ self-consistent \\threedimensional\\ (magneto-) hydrodynamic simulations for all stellar spectroscopic analysis. The computational task is simply too vast. Ultimately though, many observations will have to be analyzed in the context of such modeling, or the validity of much less demanding \\onedimensional\\ modeling will have to be more firmly established by comparing with the more realistic \\threedimensional\\ solutions. Meanwhile, (semi-empirical) \\onedimensional\\ models continue to be employed and are in many instances still at the forefront of stellar spectroscopic analysis, simply because no applicable self-consistent models exist yet. Examples of the latter are analysis of chromospheric spectra \\citep{SocasNavarro+Uitenbroek2004, SocasNavarro2007,% Avrett+Loeser2008, Centeno+Trujillo+Uitenbroek+Collados2008,% Grigoryeva+Teplitskaya+Ozhogina2009, Ermolli_etal2010}, irradiance variations resulting from large scale magnetic phenomena, spectral irradiance in the UV \\citep{Fontenla_etal2009,Shapiro_etal2010}, and cases where the curvature of the atmosphere plays a role, like interpretation of spectra of giants and supergiants, and spectra taken close to the limb in smaller stars. In an early paper \\citet{Wilson+Williams1972} discuss the effect of inhomogeneities on continuum intensity of an inhomogeneous atmosphere in the context of modeling the atmosphere of a sunspot umbra. Several contributions heretofore \\citep{Kiselman+Nordlund1995,Steffen+Ludwig+Freytag1995,Shchukina+TrujilloBueno+Asplund2005,% Ayres+Keller+Plymate2006,Scott+Asplund_etal2006,% Pereira+Asplund+Kiselman2009,Ramirez+etal2009,% TrujilloBueno+Shchukina2009,Caffau+etal2010} have spectroscopically compared the results from \\threedimensional\\ simulations with those from established \\onedimensional\\ models. \\citet{Koesterke+AllendePrieto+Lambert2008} compared the center-to-limb variation (CLV) of intensity in continua and lines computed from snapshots of a \\threedimensional\\ hydrodynamic simulation to values obtained from a \\onedimensional\\ model derived by spatially and temporally averaging their snapshots, and compared their behavior with observed CLVs, and found that the continuum intensities from their \\threedimensional\\ snapshots vary stronger with heliocentric angle than those from the average model (their figure 2). \\citet{Kiselman+Nordlund1995} compared oxygen abundances derived from two \\threedimensional\\ HD snapshots with those derived from the \\citet{HOLMUL} model and spatially averaged \\onedimensional\\ models derived from their snapshots. They found that molecular OH lines in all of their \\onedimensional\\ models were weaker than the spatially averaged ones from their simulation. \\citet{Scott+Asplund_etal2006} compared carbon abundances derived from CO modeling in \\threedimensional\\ and average \\onedimensional\\ models and found that the latter require slightly higher carbon abundance. An intermediate approach, between one- and \\threedimensional\\ modeling was presented by \\citet{Ayres+Keller+Plymate2006} in an attempt to circumvent the numerical burden of full-blown convection simulations, but still account for thermal variations in the atmosphere. They employed a so-called 1.5D transfer model in which properly weighted intensity contributions from 5 \\onedimensional\\ atmospheres with perturbed temperature stratifications with respect to an average model are added, such that the CLV of continuum intensity matches observed behavior. When comparing the oxygen abundance derived from matching CO line equivalent widths they found that the single averaged model predicts higher oxygen abundance than the 5 component model, by about 14\\% (their table 4). In this paper we further investigate the usability of \\onedimensional\\ atmospheric representations by comparing theoretical spectra of several different spectral lines and continua and their CLV behavior for models that have the same average stratification. Our goal is to identify the reasons why the spectra differ (or are similar) between the one and \\threedimensional\\ versions of the same average stratification, so that the nature of the errors are better understood when employing the less computationally demanding \\onedimensional\\ approach. In this approach we limit ourselves to the simplest case of \\onedimensional\\ models directly derived from a \\threedimensional\\ MHD snapshot by averaging. \\citet{Atroshchenko+Gadun1994} remark that such an average model is in general, not a physically consistent model, as it does not satisfy vertical pressure equilibrium, nor does it have a proper equation of state, since dynamical forces that support material in the vertical direction in the \\threedimensional\\ simulation are obviously neglected in the \\onedimensional\\ average. Like semi-empirical models, our derived models are thus not physically self-consistent, although perhaps not to the same degree, as semi-empirical models at least fulfill hydrostatic equilibrium. To prove that analysis by one-dimensional modeling is problematic there is in our view no need to discuss all possible one-dimensional models that reproduce significant parts of the average spectrum. That is a search without end because such models are not sufficiently constrained. Instead, with the proof that the one and only one-dimensional model that has, on average, the same exact thermal stratification as our three-dimensional snapshot \\emph{does not} produce the same average spectrum, we can easily turn our reasoning around and conclude that any one-dimensional model that does reproduce significant parts of that spectrum \\emph{necessarily} has a different thermal stratification and, therefore, is bound to fail in the reproduction of other parts of the spectrum. In Section~\\ref{sec:modelatmospheres} we describe and discuss the employed model atmospheres and spectral synthesis. Resulting spectra are compared in Section~\\ref{sec:spectra}, and discussion and conclusions are given in Section~\\ref{sec:conclusions}. ", "conclusions": "} The analysis of stellar spectra would greatly benefit if it would be possible to extract many physical parameters accurately through \\onedimensional\\ modeling. The radiative transfer in this type of models is orders of magnitude less computationally demanding than in self consistent and more realistic \\threedimensional\\ (M)HD models. In this paper we have attempted to further clarify if such a less numerically demanding approach is feasible, and if not, what the physical reasons for failure could be. We have approached this question in the simplest possible fashion by comparing spectra in several continua and lines from \\onedimensional\\ atmospheres that were derived from a \\threedimensional\\ MHD snapshot by straight averaging over equal geometric heights and over surfaces of equal optical depth at 500 nm. In this way we assured as much as possible that the derived \\onedimensional\\ atmospheres had the same average properties as the original snapshot, so that possible differences between the spectra are mostly the result of the inhomogeneities in the \\threedimensional\\ model. As a result we have identified several mechanisms that affect the average spectrum of the inhomogeneous \\threedimensional\\ model that cannot be adequately represented in derived \\onedimensional\\ models, and thus would lead to wrong estimates of the physical parameters extracted from the spectrum, if that spectrum were to be interpreted in the context of simpler \\onedimensional\\ models. In particular, we constructed two types of \\onedimensional\\ spatially averaged atmospheric stratifications, 1DZ, and 1DTAU and 1DT4, respectively, by averaging the thermodynamic properties of a single snapshot, MHD30G, from simulation of solar magneto-convection, over surfaces of equal geometric height and equal optical depth at 500 nm. In addition, we created a model, 1DHSE, by taking model 1DTAU and allowing its stratification to be determined by hydrostatic equilibrium. We then calculated the intensities in continua at 400, 500, 800 nm, and 4.7 $\\mu$m, two \\ion{Fe}{1} lines at 525.0 and 525.3 nm, and a CO molecular line at 4663 nm from all five models, compared the spatially averaged spectra from the \\threedimensional\\ model at different viewing angles with those from the average-property \\onedimensional\\ models at the same angles. From the calculated spectra it is immediately clear that the geometrically averaged model 1DZ provides a particularly poor representation of the average properties of the \\threedimensional\\ model in terms of spectroscopic diagnostics. The geometric averaging results in a temperature--$\\tau$ stratification that is much shallower than the stratification of average temperature in the inhomogeneous atmosphere, and thus results in much lower continuum intensities at disk center, and a much shallower CLV. Only very close to the limb where intensities emanate from relatively higher layers of the atmosphere that are less affected by granulation and thus more homogeneous, and where the obliquity of the rays naturally averages over horizontal inhomogeneities, the intensities from this model converge on the averaged ones from MHD30G. It is no surprise that intensities derived from models 1DTAU and 1DT4 are much closer to the spatially averaged ones from model MHD30G, as radiation at a given wavelength naturally comes from a layer with limited range in optical depth around unity. Yet, even in continua the average intensity from model MHD30G lies well above those of the tau-averaged models, in particular at shorter wavelengths, and closer to disk center. This underestimate of continuum intensities by the \\onedimensional\\ models is the result of the non-linear mapping of temperature inhomogeneities into intensity fluctuations by the Planck function. In Section~\\ref{sec:spectra} we show that the average intensity of an atmosphere with temperature fluctuations is always higher than the intensity of a homogeneous atmosphere with the same average temperature stratification, because of the strict positivity of the second derivative of the Planck function with temperature. For large temperature and/or long wavelengths the second derivative of $B_{\\lambda}$ vanishes, and temperature inhomogeneities are mapped linearly into intensity variations, so that differences between the \\threedimensional\\ and derived \\onedimensional\\ models disappear, as is evident in the continuum at 4.7 $\\mu$m (Figure~\\ref{fig_int_CLV}). A corollary of the non-linear mapping of the Planck function is that the interpretation of an average spectrum from a horizontally inhomogeneous atmosphere, in terms of a one-dimensional model will lead to an overestimate of the average temperature gradient in the atmosphere. Molecular lines are often used in abundance determinations, because they are generally weak, less affected by thermal broadening because of their mass, and plentiful. However, their concentration is sensitive to thermal conditions. To investigate how temperature and density inhomogeneities affect the average concentration molecules and emergent spectra of their lines, we calculated the CO number densities in the inhomogeneous model MHD30D and averaged models 1DTAU and 1DZ, as well as the emergent spectrum of the CO 7-6 R68 rotation-vibration line for different viewing angles. In Section~\\ref{sec:moldensity} we show that the presence of both temperature and density inhomogeneities to first order always lead to an \\textit{increase} of the average molecular density of a diatomic molecule over the values in a homogeneous atmosphere with the same average temperature and density stratification. As an example we show the calculated number density of the CO molecule in models MHD30G and 1DTAU as a function of average optical depth in Figure~\\ref{fig:COnum}. The underestimate of molecular densities by a one-dimensional model that is used to represent the average of an inhomogeneous atmosphere will lead to an overestimate of the abundances of the constituent atoms of the molecule if molecular lines are used for determination of these abundances. It is obvious that \\onedimensional\\ average-property models cannot represent the complex mass motions that result from convection. One effect that the average models necessarily fail to incorporate is the convective blue shift and line characteristic ``C-shaped'' asymmetry of the bisectors of spectral lines that results from the correlation of brightness and upflows and asymmetry between up- and downflows in the convective flows. We identify one other important aspect of the convective flows that \\onedimensional\\ models lack, namely the broadening and line weakening that horizontal motions produce for moderate viewing angles ($\\mu \\approx 0.7$). The weakening results from the shift of opacity out of the line core by horizontal motions, both to the red and to the blue, increasing formation height and raising the central intensity. This effect leads to an initial increase in the line-core intensity as function of increasing viewing angle (decreasing $\\mu$), and a much shallower decline of the residual line-core intensity as is clearly shown in Figure~\\ref{fig:linecores}. The horizontal velocity effect is apertly absent in the CLVs obtained from the \\onedimensional\\ models. It thus also affects the disk-integrated line profile shape that will therefore be wrongly interpreted if analyzed in the context of such average-property models. Observational evidence for the the shallow behavior of the residual line-core intensity is seen in \\citet{RodriguezHidalgo+Collados+Vazques1994} who determined the CLV of this property in the \\ion{Fe}{1} 593.02 nm line, among others. Curiously, the \\ion{Mn}{1} 539.47 nm line shows behavior more in line with our results from the \\onedimensional\\ models. This is perhaps the result of the large hyperfine structure broadening of manganese lines, which makes them less susceptible to convective velocity broadening, as explained by \\citet{Vitas+Viticchie+Rutten+Voegler2009}. In summary, our comparison of spectra, and analysis of the physical effects of temperature and density inhomogeneities make clear that extreme caution has to be taken when the average spectrum of a horizontally inhomogeneous atmosphere is interpreted in the context of a \\onedimensional\\ atmospheric representation. Not only the lack of incorporation of convective motions into \\onedimensional\\ models, but also the non-linearities involved in establishing molecular equilibrium and level populations as function of temperature and density, and the non-linearities of the Planck function as function of temperature likely lead to misinterpretation of the observed spectrum. In the work here we have concentrated on spatial averaging, but our derivations and contentions hold equally well in case of temporal averaging. In addition, we suspect that non-linearities introduced by non-LTE conditions, which we have not addressed in this paper, have to be treated with even more caution. The realism of chromospheric modeling, which still relies heavily on one-dimensional models, therefore needs careful examination in this respect. A very illustrative example is given by the simulations presented by \\citet{Carlsson+Stein1994}, who simulate acoustic shock propagation in the solar atmosphere. The shocks heat the atmosphere, but their main effect is an enhancement of chromospheric emission. When the resulting time-averaged spectrum is modeled with a \\onedimensional\\ hydrostatic semi-empirical model, the resulting atmosphere is required to have a chromospheric temperature rise. This temperature rise is very much at odds with the mean temperature stratification of the simulations which shows a monotonic decline with height. Since the 1DTAU model we employed exactly matches the average thermal stratification of the \\threedimensional\\ snapshot, but produces different intensity values in lines, and even continua, we conclude that a one-dimensional model that reproduces the average spectrum of an inhomogeneous atmosphere necessarily has a different average stratification than this atmosphere. It is therefore in our opinion not possible to produce a one-dimensional atmospheric model that is at the \\textit{same time} spectroscopically equivalent \\textit{and} has matching average physical properties." }, "1101/1101.0646_arXiv.txt": { "abstract": "To better understand long-term flare activity, we present a statistical study on soft X-ray flares from May 1976 to May 2008. It is found that the smoothed monthly peak fluxes of C-class, M-class, and X-class flares have a very noticeable time lag of 13, 8, and 8 months in cycle 21 respectively with respect to the smoothed monthly sunspot numbers. There is no time lag between the sunspot numbers and M-class flares in cycle 22. However, there is a one-month time lag for C-class flares and a one-month time lead for X-class flares with regard to sunspot numbers in cycle 22. For cycle 23, the smoothed monthly peak fluxes of C-class, M-class, and X-class flares have a very noticeable time lag of one month, 5 months, and 21 months respectively with respect to sunspot numbers. If we take the three types of flares together, the smoothed monthly peak fluxes of soft X-ray flares have a time lag of 9 months in cycle 21, no time lag in cycle 22 and a characteristic time lag of 5 months in cycle 23 with respect to the smoothed monthly sunspot numbers. Furthermore, the correlation coefficients of the smoothed monthly peak fluxes of M-class and X-class flares and the smoothed monthly sunspot numbers are higher in cycle 22 than those in cycles 21 and 23. The correlation coefficients between the three kinds of soft X-ray flares in cycle 22 are higher than those in cycles 21 and 23. These findings may be instructive in predicting C-class, M-class, and X-class flares regarding sunspot numbers in the next cycle and the physical processes of energy storage and dissipation in the corona. ", "introduction": "Since the 11-year period of the solar cycle was discovered by Schwabe (1844), the long-term activity of the solar cycle has become a hot issue in the field of solar physics. Sunspots are taken as the most famous and typical indicators of solar activity. Sunspot activity has complex spatial and temporal behavior. Carrington (1858, 1859) investigated a drift latitude of sunspot motion towards the equator and a variation of the rotation rate of the Sun. Hathaway et al. (2003) examined the drift of the centroid of the sunspot area toward the equator in each hemisphere from 1874 to 2002 and found that the drift rate slows as the centroid approaches the equator. The distribution of sunspots and flares in a solar cycle exhibits a ``butterfly diagram\" (Carrington, 1858, Maunder 1904, 1913; Garcia 1990; Li et al. 2003) and the cycle appears uniformly in both hemispheres on average (Newton \\& Milson 1955; White \\& Trotter 1977). Sunspot activity is described by sunspot numbers and sunspot areas in general. Consequently, the sunspot numbers and sunspot areas are used to investigate long-term solar activity (Li et al. 2009). Because of the periodicity of sunspot activity (sunspot numbers and sunspot areas), the associated eruptions (flares, CMEs) appear to have the similar periodicity to that of sunspot activity (Storini \\& Hofer 1999). In addition, the soft X-ray flares were significantly delayed with respect to sunspot numbers, with a time lag of two to three years between the peak times in solar cycle 21 (Wagner 1988; Aschwanden 1994, Bromund et al. 1995). However, Wilson (1993) did not find evidence for a time lag between the maxima of the rates of optical flares and X-ray background flux in solar cycle 22. Wheatland \\& Litvinenko (2001) found that there is an average delay of about 6 months considering cycles 21 and 22 together by using correlation analysis. For the cycle 23, Tan (2010) compared the sunspot numbers with the distributions of the appearance rate of solar GOES flares. They found that there is a time lag between the maximum value of sunspot numbers and the maximum values of the annual numbers of C-class, M-class, and X-class flares. The annual averaged relative sunspot number reaches its maximum in about 2000 and the annual numbers of C-class, M-class, and X-class flares reach their maxima in about 2001, 2001, and 2002 respectively. Temmer et al. (2003) found that there is a characteristic time lag between flare activity and sunspot activity in the range of 10-15 months for solar cycles 19, 21, and 23 by using the number of the flares in each month. In order to investigate all kinds of long-term flare activity, the flare index calculated by T. Atac and A. Ozguc from Bogazici University Kandilli Observatory and the daily soft X-ray flare index defined by Antalov$\\acute{a}$ (1996) are two main flare indices in the present research. For instance, Li et al. (2010) found that the northern-hemispheric flare activity should lead the southern-hemispheric flare activity for low-frequency components by using the former flare index. Joshi et al. (2004) reported a real north-south asymmetry during solar minimum and obtained the significant periods of approximately 28.26 days, 550.3 days, and 3.72 years by using the latter flare index. In this paper, we focus on the three solar cycles (21, 22, and 23) to investigate the phase relation between sunspot numbers and flare activity (C-class, M-class, and X-class flares) by using a flare index, which is similar to Antalov$\\acute{a}$ (1996). ", "conclusions": "In this paper, we present a statistical study on three types of soft X-ray flares from May 1976 to May 2008. We use the data of the smoothed monthly peak fluxes of C-class, M-class, and X-class flares and the smoothed monthly sunspot numbers. The main results are as follows: 1. The smoothed monthly peak fluxes of C-class, M-class, and X-class flares have a very noticeable time lag of 13, 8, and 8 months respectively with respect to the smoothed monthly sunspot numbers in cycle 21. 2. There is a one-month time lag for the smoothed monthly peak fluxes of C-class flares, a one-month time lead for the smoothed monthly peak fluxes of X-class flares and no time lag for the smoothed monthly peak fluxes of M-class flares with respect to the smoothed monthly sunspot numbers in cycle 22. 3. The smoothed monthly peak fluxes of C-class, M-class, and X-class flares have a very noticeable time lag of one month, 5 months, and 21 months in cycle 23 respectively with respect to the smoothed monthly sunspot numbers. 4. If we take the three types of flares together, we find that the soft X-ray flares have an obvious time lag of 9 months in cycle 21, no time lag in cycle 22 and a characteristic time lag of 5 months in cycle 23 with respect to the smoothed monthly sunspot numbers. 5. The correlation coefficients of the smoothed monthly peak fluxes of M-class and X-class flares and the smoothed monthly sunspot numbers are higher in cycle 22 than those in cycles 21 and 23. 6. The correlation coefficients between the three kinds of soft X-ray flares in cycle 22 are higher than those in \\textbf{21} and 23. In this paper, we adopted the method developed by Antalov$\\acute{a}$ (1996) and used the smoothed monthly peak fluxes of soft X-ray flares. We divided the soft X-ray flares into three groups according to the classification of soft X-ray flare, which is slightly different from Temmer et al. (2003). When the numbers of flares are multiplied by the peak fluxes, there is a great change in value pertaining to the original flare numbers. From table 1, one can see that the numbers of M-class, X-class flares are small but they have large monthly peak fluxes. The peak fluxes of flares can be approximately taken as the energy released by flares. As M-class and X-class flares have released more energy than that of B-class and C-class flares, we think that the peak fluxes of flares taken as the flare index may be closer to the energy released by the flares. That is why we used this method to analyze the phase relation. The activity of the C-class, M-class and X-class flares exhibited an obvious time lag behind sunspot activity in odd-numbered cycles, while in even-numbered cycles M-class flares do not. Also, the C-class and X-class flares show reversed behavior in the even cycle. If we consider the three types of flares together, we find that there is an obvious time lag in the odd solar cycles and no time lag in the even solar cycle. Though different sorts of flares show different time lags or time leads with respect to sunspot numbers, there is a regular time delay in the odd solar cycles and no time delay in the even solar cycle between soft X-ray flares and sunspot numbers after we sum the three types of flares together. The value of summing the monthly peak fluxes of three types of flares approximately denotes the total energy released by flares. These results imply that the flare activities are related to the 22-year magnetic cycle of the Sun. We need forthcoming data to confirm these results and model the process of flare energy storage and dissipative mechanism. \\begin{table*}{} \\begin{center} TABLE 1 \\\\ The numbers and peak fluxes of three types of soft X-ray flares (C-class, M-class, and X-class flares) in cycles 21, 22, and 23. The values in parentheses denote the peak fluxes of the defined flares. ``Total\" denotes the total numbers and the total peak fluxes in the selected period from May 1976 to May 2008. The unit of flare peak flux is erg$\\cdot$$cm^{-2}$$\\cdot$$s^{-1}$.\\\\ \\vspace{0.5cm} \\begin{tabular}{llllllll} \\tableline\\tableline \\label{tbl-1} Flare class &21 cycle &22 cycle & 23 cycle &Total \\\\ \\tableline C &14576 (46.76)&12430(41.79) &13148 (39.39)&40154(127.94)\\\\ M& 2175(50.22)&2019(48.15)&1439(34.37)&5633(132.74)\\\\ X&165(37.59)&152(35.46)& 126(31.08)&443(104.13) \\\\ \\tableline \\end{tabular} \\end{center} \\end{table*} \\begin{table*}{} \\begin{center} TABLE 2 \\\\ Correlation coefficients between the smoothed monthly peak fluxes of C-class flares, M-class flares, X-class flares and the smoothed monthly sunspot numbers in cycle 21 (the values in parentheses denote correlation coefficients in cycles 22 and 23.). ``SSN\", ``SFC\" ,``SFM\", and ``SFX\" indicate the smoothed monthly sunspot numbers, the smoothed monthly peak fluxes of C-class, M-class, and X-class flares respectively. \\\\ \\vspace{0.5cm} \\begin{tabular}{lllllll} \\tableline\\tableline \\label{tbl-1} & SFC & SFM & SFX \\\\ \\tableline SSN &0.751(0.974, 0.984) &0.825(0.950, 0.908) & 0.778(0.921, 0.457) \\\\ SFC && 0.951(0.954, 0.948)&0.874(0.908, 0.478) \\\\ SFM &&& 0.959(0.974, 0.655)\\\\ \\tableline \\end{tabular} \\end{center} \\end{table*} \\begin{figure*} \\centering \\includegraphics[width=12cm]{fig1.eps} \\caption{The distribution of the monthly numbers of three types of flares from May 1976 to May 2008 respectively.} \\label{} \\end{figure*} \\begin{figure*} \\centering \\includegraphics[width=12cm]{fig2.eps} \\caption{The distribution of the monthly peak fluxes of C-class (a), M-class (b), X-class (c) flares, total peak fluxes of three types of flares(d), the monthly sunspot numbers(e) (blue lines) and their 13-point smoothed values (red lines) superimposed.}% \\end{figure*} \\begin{figure*} \\centering \\includegraphics[width=12cm]{fig3.eps} \\caption{Cross-correlation coefficients between the smoothed monthly peak fluxes of C-class ($a_1$-$a_3$), M-class ($b_1$-$b_3$), X-class ($c_1$-$c_3$) flares, the total peak fluxes of the three types of flares ($d_1$-$d_3$) and the smoothed monthly sunspot numbers in cycles 21, 22, and 23. The abscissa indicates the shift of the smoothed monthly sunspot numbers with respect to the smoothed monthly peak fluxes of C-class, M-class, and X-class flares, with negative values representing backwards shifts.}% \\end{figure*}" }, "1101/1101.4854_arXiv.txt": { "abstract": "{The number of strong iron $K_\\alpha$ line detections in Seyfert AGN is clearly growing in the {\\it{Chandra, XMM-Newton}} \\rm{and} {\\it{Suzaku}} \\rm{era}. The iron emission lines are broad, some are relativistically blurred. These relativistic disk lines have also been observed for galactic black hole X-ray binaries. Thermal components found in hard spectra were interpreted as an indication for a weak inner cool accretion disk underneath a hot corona.} {Accretion in low-mass X-ray binaries (LMXB) occurs during phases of high and low mass accretion rate, outburst and quiescence, soft and hard spectral state, respectively. After the soft/hard transition for some sources a thermal component is found, which can be interpreted as sustained by re-condensation of gas from an advection-dominated flow (ADAF) onto the disk. In view of the similarity of accretion flows around stellar mass and supermassive black holes we discuss whether the broad iron emission lines in Seyfert 1 AGN (Active Galactic Nuclei) can be understood as arising from a similar accretion flow geometry as in X-ray binaries.} {We derive accretion rates for those Seyfert galaxies for which broad iron emission lines were observed, the ``best candidates'' in the investigations of Miller (2007) and Nandra et al. (2007). For the evaluation of the Eddington-scaled rates we use the observed X-ray luminosity, bolometric corrections and black hole masses from the literature most values taken from the investigation of Fabian \\& Vasudevan (2009).} {The accretion rates derived for the Seyfert galaxies in our sample are less than 0.1 of the Eddington rate for more than half of the sources. For $10^7$ to $10^8 M_\\odot$ black holes in Seyfert 1 AGN this limit corresponds to 0.01 to 0.2 $M_\\odot/{\\rm{yr}}$. This documents that the sources probably are in a hard spectral state and iron emission lines can arise from an inner weak accretion disk surrounded by an ADAF as predicted by the re-condensation model. Some of the remaining sources with higher accretion rates may be in a spectral state that is comparable to the ``very high'' state of LMXBs.} {Our investigation shows that in quite a number of Seyfert AGN the broad iron emission lines may indeed originate in a weak inner disk below the ADAF, close to the black hole, indicating the same accretion flow geometry as recently found for LMXBs. For the accretion history one then concludes that the accretion rates were higher in the outer radii at some earlier time .} ", "introduction": "The physical processes of accretion onto galactic and supermassive black holes and the accretion flow geometry in the innermost regions are key features for modeling relativistic spectral lines and for estimates of the black hole spin. Observations in recent years with the {\\it{Chandra X-ray Observatory}}, the {\\it{X-ray Multi-Mirror Mission-Newton}} and {\\it{Suzuka}} reveal for some AGN broad, often relativistic iron emission lines from the innermost regions close to the black hole. Iron emission lines are the most obvious response of an accretion disk to irradiation by an external source of hard X-rays. When X-rays from a hot coronal flow fall on optically thick cool material, they induce fluorescence and are backscattered, resulting in a Compton reflection spectrum with a prominent emission line of iron K$\\alpha$. The accretion geometry strongly depends on the mass accretion rate. For LMXBs it is well known that on the one hand for high rates an optically thick, geometrically thin, radiatively efficient Shakura-Sunyaev accretion disk reaches inward to the innermost stable circular orbit (ISCO), or on the other hand for low accretion rates (lower than the critical rate of spectral state transition $\\dot m_{\\rm{crit}}$), a more spherical optically thin, hot advection-dominated flow fills the inner region (ADAF; Narayan \\& Yi 1994, 1995, Abramowicz et al. 1995; see Narayan 2005, Yuan 2007, Narayan \\& McClintock 2008 for reviews). At larger distance an outer Shakura-Sunyaev disk exists, whose truncation radius recedes as the mass accretion rate decreases. These ``high'' and ``low'' states are documented by a soft and hard spectrum. The critical accretion rate $\\dot m_{\\rm{crit}}$ lies around a few to 10 percent of the Eddington accretion rate (Esin et al. 1997). For supermassive black holes, especially low-luminosity AGN, truncated thin disks were found, along with a hard spectrum (Narayan et al. 1998). The power-law spectrum arising from the coronal emission is scale-invariant (accretion rate in units of Eddington accretion rate), while the disk emission of AGN peaks in the UV range, in contrast to the emission in soft X-rays from disks around stellar mass black holes. Vasudevan \\& Fabian (2007) discuss the similarity of accretion states in AGN and galactic black hole sources and point out that the radiation in the UV is an important contribution to the bolometric luminosity. They derived bolometric corrections for AGN using recent results from the {\\it{ Far Ultraviolet Spectroscopic Explorer (FUSE)}}, together with data from the Optical Monitor (OM) archives and X-ray data from the {\\it{XMM-Newton}} archive. They expanded the investigation and for the first time took into account simultaneous observations (Vasudevan \\& Fabian 2009). The accretion geometry in the two spectral states, either an ADAF (or one of its variants) in the inner regions or a disk reaching inward to the ISCO, seemed to exclude each other for a long time. But the recent observations of LMXBs seem to indicate that during in intermediate state, the brightest hard state (after soft/hard transition in outburst decay), an ADAF in the inner region and a weak innermost disk can both be present at the same time (see Figure 1). These co-existing hot and cool gas flows inward toward the black hole clearly interact. The interaction causes mass and angular momentum exchange between disk and ADAF, either a mass flow from the disk into the ADAF, evaporation of the disk, or a reverse mass flow, re-condensation of gas from the ADAF onto the disk (Liu et al. 2006, Meyer et al. 2007). Indications for a disk near the black hole during a canonical low-hard state were found for several X-ray binaries, with clearest hints in GX 339-4, SWIFT J1753.5-0127 and XTE J1817-330. The re-condensation model allows us to understand the observed mass flow rates in the inner disk (Liu et al. 2007, Taam et al. 2008). Observational evidence for the possible presence of a thermal component during the hard state in eight sources, broad skewed Fe K$\\alpha$ lines in half of the sample, were shown in the recent work by Reis et al. (2010). The authors interprete the reflection features as caused by illumination of a more-or-less permanent disk by the hard, power-law component of a jet. We note that an extremely skewed relativistic Fe K$\\alpha$ emission line was also found in the spectrum of GX 339-4 during the bright phase of its 2002-2003 outburst, in a ``very high'' state. As Miller (2007) pointed out, this observation is of interest for understanding the relativistic iron lines observed in Seyfert 1 AGN. We discuss these observations in connection with the detection of broad iron K and L line emission in the narrow-line galaxy 1H0707-495. The existence (at least for some time) of a weak disk in the innermost region around stellar mass black holes during the hard spectral state strengthens the expectation to find the same phenomenon for accreting supermassive black holes as well. Seyfert 1 AGN are the class of supermassive black holes that offer the best view of the innermost accretion region. Tanaka et al. (1995) observed the first asymmetric disk line profile in the Seyfert 1 AGN MCG-6-30-15 using the ASCA/SIS. Now {\\it {XMM-Newton, Chandra}} and {\\it{ Suzaku}} allow us to detect and measure relativistic disk lines in Seyferts. The Fe K-shell emission lines are the strongest X-ray emission lines both in AGN and X-ray binaries. In a detailed review of relativistic X-ray emission lines from inner accretion disks around black holes Miller (2007) discusses observational and theoretical developments. The review of Nandra et al. (2007) focuses on broad iron lines in Seyfert galaxies observed by {\\it{XMM-Newton}} and raises the question how frequent this broadening is, whether it originates in an accretion disk, and how robust the evidence for an accretion disk is. At first glance it is not clear whether the lines originate in a weak disk below a prominent hard coronal X-ray flux or are connected with accretion via an untruncated disk during the ``very high'' state. \\begin{figure} \\begin{center} \\includegraphics[width=8.8cm]{15478fg1.eps} \\caption{\\label{schematic} Change of the geometry of the accretion flow with decreasing mass accretion rate scaled to the Eddington rate $\\dot m$, ($\\dot m_c$ critical rate for which the state transition happens): (1) soft state, (2) transition to the hard state begins, formation of a gap where evaporation is most effective (3-4) hard states with disk truncation receding outward, the weak inner disk disappears (Fig.1, Meyer-Hofmeister et al.2009).} \\end{center} \\end{figure} The aim of our paper is to analyze under which circumstances the broad iron emission lines in Seyfert galaxies may originate in a weak disk in the innermost region around the black hole, similar to the situation of stellar mass black holes (and neutron stars). In Section 2 we refer to the process of disk evaporation, which for low accretion rates leads to the truncation of the standard Shakura-Sunyaev disk at a certain radius. In Section 3 we discuss how changes of the mass flow rate result in distinct spectral states for both LMXBs and AGN. The mass flow toward the inner regions can be modulated by the ionization instability and magnetorotational and gravitational instabilities (Siemiginowska et al. 1996, Menou \\& Quataert 2001). In the Appendix we discuss how the ionization instability can modulate the mass flow toward the inner regions. As a consequence of the disk evaporation process a new picture arises for the presence of these disk instabilities. These mass flow variations can cause changes between hard and soft spectral states and lead to the temporary existence of a weak inner disk. We discuss in Section 4 for which rates the appearance of an inner disk can be expected as a transient phenomenon in Seyfert galaxies. In Section 5 we derive accretion rates from observations for the best candidates of broad iron emission lines in the samples of Seyfert galaxies of Miller (2007) and Nandra et al. (2007). In Section 6 we discuss a different accretion flow geometry: broad emission lines from an untruncated disk in bright sources in the ``very high'' state. We note that a different picture is also discussed for the accretion flow geometry, an always present accretion disk together with a jet. For supermassive black holes, Sgr A* and low-luminosity AGN such a jet-dominated situation was suggested (Falcke \\& Biermann 1999, Falcke \\& Markoff 2000, Falcke et al. 2000) as an alternative to the ADAF solution (Narayan et al. 1998, Di Matteo et al. 2000, 2003). The iron emission lines could be caused by illumination from a non-thermal jet. ", "conclusions": "For LMXBs broad iron emission lines found in an intermediate hard spectral state were interpreted as originating from a weak inner disk, which is left over from a disk reaching inward to the ISCO at earlier time. The accretion flow geometry in the thin disk+ADAF model scales with the black hole mass. We discussed whether such a model for the interpretation of iron emission lines is also applicable to Seyfert galaxies. We used recent observations and interpretations of broad, partly relativistically blurred iron emission lines in Seyfert 1 AGN. We considered sources in a hard spectral state. For the best candidates in the samples of Miller (2007) and Nandra et al. (2007) we evaluated accretion rates (scaled to the Eddington accretion rate) to see whether these rates point to a hard spectral state. As a limiting Eddington-scaled upper accretion rate for sources in the hard spectral state we chose the value 0.1, taking into account the observational results for the state transition luminosities of LMXBs. It is difficult to judge the uncertainty of the accretion rates evaluated for the Seyfert galaxies. Therefore we cannot definitely determine the spectral state. As dicussed in Sect.3, it is not possible to derive firm information on the spectral state from the photon index of the spectral energy distribution. For the majority of the sources considered here we found accretion rates $\\le$ 0.1 (see Table 1), agreeing with the theoretical expectation of a weak inner accretion disk below the ADAF during the low/hard state for AGN as well. The appearance of iron K$\\alpha$ lines signifies the presence of cool material close to the accretng black hole, irradiated by a power-law continuum from a corona above. In this state the inner disk carries only a fraction of the total accretion flow, the major part flows through the corona and provides the hard luminosity. This situation is desribed in the modeling of the ``truncated disk + corona + re-condensed inner disk'' geometry. Disk truncation is an inherent feature of disk evaporation that explains the appearance of the hot, optically thin flow for low accretion rates (and also the hysteresis in transition luminosity between soft and hard spectral states). Besides the Seyferts with accretion rates below the critical value there is a significant number of sources in our sample (40\\%) that show accretion rates above this limit, though not as high as in some LMXBs in high state with iron emission lines (a significant fraction of the Eddington value). These LMXB sources probably are related to the so-called ``very high'' state, where the disk material is illuminated by a strong corona on top of an inner disk (Done, Gierli\\'nski and Kubota 2007). In the LMXB case a clear gap exists in the accretion rate distribution between the two groups with low and high rates, no sources with rates slightly above 0.1 are known. The galactic sources XTE J1650-500 and GX 339-4 were observed in such a very high state, and extremely skewed relativistic Fe K$\\alpha$ emission lines were detected (Miller et al. 2002a, Miller et al. 2004). Miller (2007) pointed out that the spectral and variability phenomena then closely resemble the behavior seen in Seyfert AGN. The narrow-line Seyfert 1 galaxy 1H0707-495 seems to be a particularly bright source of this kind, with an accretion rate just below the Eddington rate. The Seyfert galaxies with accretion rates somewhat about the critical value (with no analogous stellar sources) deserve particular attension because the existence of these sources might indicate a difference between the accretion flow geometry around stellar mass and supermassive black holes. Theory predicts an untruncated disk for sources with accretion rates above 0.1. The reflection features then indicate the presence of significantly stronger coronae at comparable accretion rates in these AGN than in LMXBs. If strong magnetic flux is the cause of this phenomenon, these coronae might well also be the base of a jet issuing from them. Such differences between accretion in stellar mass and supermassive black holes certainly deserve further work. Weak inner disks below an ADAF-type accretion flow in the inner region of an accreting black hole can be considered as the natural remainder of an originally standard Shakura-Sunyaev disk that had reached inward all the way to the ISCO at high accretion rates and has now, when the accretion rate has dropped below its critical value, become truncated. This could therefore be a hint to a higher accretion rate in the past. Moderate accretion rate variations caused by the ionization instability in AGN disks could well allow the inner disk to exist for long times (as in the case of the galactic black hole binary Cyg X-1). Seyfert 1 galaxies with $10^7$ to $10^8 M_\\odot$ black holes and accretion rates of about 0.01 to 0.1$M_\\odot/{\\rm{yr}}$ seem specially suited to display this weak inner disk accretion flow geometry." }, "1101/1101.4253_arXiv.txt": { "abstract": "We investigate the efficiency of interstellar polarization $p_\\lambda/A_\\lambda$, where $p_\\lambda$ is the fractional linear polarization and $A_\\lambda$ is extinction, in 14 lines of sight as a function of wavelength $\\lambda$. We have used the data of lines of sight to the Pleiades cluster obtained with the low-dispersion spectropolarimeter HBS as well as those in literature. It is found that the polarization efficiency $p_\\lambda/A_\\lambda$ is proportional to $\\exp(-\\beta/\\lambda)$ in wavelength $\\lambda \\approx 0.4\\sim0.8 \\micron$, where $\\beta$ is a parameter which varies from 0.5 to 1.2 $\\micron$. We find that $\\beta$ is negatively correlated with the dust temperature deduced from infrared data by Schlegel et al., suggesting that the polarization efficiency is higher in short wavelength for higher temperature. According to the alignment theory by radiative torques (RATs), if the radiation is stronger, RATs will make small grains align better, and the polarization efficiency will increase in short wavelength. Our finding of the correlation between $\\beta$ and the temperature is consistent with what is expected with the alignment mechanism by RATs. ", "introduction": "Observed linear polarization in the light from distant stars, i.e., often called as \"interstellar polarization\", is interpreted as a phenomenon of dichroic extinction, and shows that interstellar grains are optically anisotropic and aligned, although the mechanism of the alignment has been still on debate (\\cite{L07}, for a recent review). The alignment of grains had been explained with the paramagnetic relaxation of thermally spinning grains that obtain angular momentum by collisions with gas particles (\\cite{DG51}, hereafter DG). However, the DG mechanism is not efficient, and it cannot explain the interstellar polarization quantitatively. For more efficient alignment, \\citet{Purcell79} assumed a spin-up of grain by the ejection of molecular hydrogen from grain surface (the \"pinwheel mechanism\"), though there still remain problems in quantitative explanations \\citep{LD99,RL99}. \\citet{DM76} first pointed out that irregularly shaped grains that have \"helicity\" can spin up by radiative torques (hereafter RATs). More recently, \\citet{DW96} showed that RATs are very effective to align grains. Since magnetic moments within rotating grains are induced by the Barnett effect, the grains precess around the magnetic field, and the direction of alignment is usually parallel to the interstellar magnetic field (e.g. \\cite{DW97}; \\cite{LH07}), i.e., the same direction as that by the DG mechanism. The alignment by RATs can be more efficient if grains have superparamagnetic inclusions \\citep{LH08} or if the pinwheel mechanism is working with RATs \\citep{HL09}. The efficiency of the RATs alignment varies with strength and spectral energy distribution of the radiation field, and thus the size of aligned grains should vary accordingly (\\cite{DW96}; \\cite{ChoLaz07}). Observationally, the maximum wavelength $\\lambda_{\\rm max}$ of polarization in dark clouds was shown to be correlated with extinction $A_{\\rm V}$ in the V-band \\citep{Whittet01,AP07}. \\citet{AP10} showed that grain alignment is enhanced by the stellar radiation in the vicinity of a young star HD 97300 in the Chamaeleon I cloud. For stars in the Taurus dark cloud, \\citet{Whittet08} showed that the polarization efficiency $p_{\\rm K}/\\tau_{\\rm K}$ in the K-band, where $\\tau_{\\rm K}$ is optical depth, decreases smoothly with $A_{\\rm V}$ beyond the region where ice mantle feature was detected. This suggests that the alignment efficiency is not directly related to the state of grain surface, as is expected by the \"pinwheel\" alignment. Those observations suggest that the RATs alignment works in dark clouds and star forming regions. However, it is still not clear whether RATs alignment works or not in more diffuse clouds. The Pleiades cluster % is associated with the diffuse reflection nebula, where grain alignment may be enhanced by strong stellar radiation if the alignment by RATs works. This motivated us to observe polarization in the lines of sight to stars in the Pleiades cluster with the low-dispersion spectropolarimeter HBS \\citep{K99}. In this Letter, using our polarimetric data and those available from literature, we investigate correlations between polarization quantities and dust temperature, because such correlations may be expected from the RATs alignment theory. ", "conclusions": "\\subsection{Fractional Polarization and Position Angle} The observed fractional polarization $p_\\lambda$ and position angle $\\theta_\\lambda$ are shown in Figure \\ref{fig1}, except for HD 23985 and HD 24118 which show low polarization $\\lesssim0.1\\%$. We assume an empirical formula for $p_\\lambda$ by \\citet{Ser75}: \\begin{equation} p_\\lambda = p_{\\rm max}\\exp(-K\\ln^2(\\lambda/\\lambda_{\\rm max})), \\label{eq_ser} \\end{equation} where $p_{\\rm max}$ is maximum polarization, $K$ is a parameter that determines width of the curve, and $\\lambda_{\\rm max}$ is the wavelength at $p=p_{\\rm max}$. Those derived values of $p_{\\rm max}$, $\\lambda_{\\rm max}$, and $K$ are tabulated in Table \\ref{tab1}. Figure \\ref{fig1}a shows that $p_\\lambda$ is well expressed with equation (\\ref{eq_ser}). The position angle $\\theta_\\lambda$ is almost constant, though $\\theta_\\lambda$ of 19 Tau varies with $\\lambda$ in short wavelength $1/\\lambda \\gtrsim 2.3\\micron^{-1}$ (Figure \\ref{fig1}b). \\begin{table*} \\caption{Polarization of Stars in the Pleiades Cluster}\\label{tab1} \\begin{center} \\begin{tabular}{lccrrrcccccccc} \\hline Name & $p_{\\rm V}$\\footnotemark[$*$] & $p_{\\rm max}$\\footnotemark[$*$] & $\\lambda_{\\rm max}$\\footnotemark[$\\dagger$] & $K$\\footnotemark[$\\dagger$] & $\\theta_{\\rm V}$\\footnotemark[$\\dagger$] & Method & $R_{\\rm V}$ & $A_{\\rm V}$ & $\\alpha$ & $\\beta$ & $T_{\\rm dust}$ \\\\ \\ (Sp.Type) & [\\%] & [\\%] & $[\\micron]$ & & [deg] & & & [mag] & [\\%/mag] & $[\\micron]$ & [K] \\\\ \\hline 19 Tau & 0.26 & 0.28 & 0.36 & 0.26 & 140.2 & 1 & 2.74$\\pm.57$ & 0.13 & 2.1$\\pm$0.7 & 0.72$\\pm$.01 & 20.1 \\\\ \\ \\ \\ \\ \\ (B5IV) & - & - & $\\pm.11$ & $\\pm.18$ & $\\pm$4.4 & 2 & 3.20$\\pm.17$ & 0.13 & 2.1$\\pm$0.7 & 0.64$\\pm$.01 & - \\\\ 27 Tau & 0.34 & 0.35 & 0.61 & 0.91 & 112.6 & 2 & 3.20$\\pm.17$ & 0.09\\footnotemark[$\\ddagger$] & 4.1$\\pm$0.8 & 0.94$\\pm$.08 & 19.7 \\\\ \\ \\ \\ \\ \\ (B8III) & - & - & $\\pm.03$ & $\\pm.45$ & $\\pm$3.4 & - & - & - & - & - & - \\\\ HD23512 & 2.34 & 2.38 & 0.61 & 1.06 & 27.2 & 1 & 3.48$\\pm.11$ & 1.15 & 2.09$\\pm$.07 & 0.82$\\pm$.02 & 20.1 \\\\ \\ \\ \\ \\ \\ (A0V) & - & - & $\\pm.00$ & $\\pm.07$ & $\\pm$0.5 & 2 & 3.20$\\pm.17$ & 1.15 & 2.12$\\pm$.07 & 0.94$\\pm$.06 & - \\\\ HD23753 & 0.27 & 0.27 & 0.51 & 0.57 & 104.8 & 1 & 3.17$\\pm.76$ & 0.09 & 3.1$\\pm$1.6 & 0.65$\\pm$.14 & 19.3 \\\\ \\ \\ \\ \\ \\ (B8V) & - & - & $\\pm.01$ & $\\pm.16$ & $\\pm$4.2 & 2 & 3.20$\\pm.17$ & 0.09 & 3.1$\\pm$1.6 & 0.67$\\pm$.04 & - \\\\ HD24178 & 0.50 & 0.50 & 0.58 & 1.38 & 128.2 & 1 & 2.57$\\pm.17$ & 0.39 & 1.33$\\pm$.13 & 1.00$\\pm$.05 & 17.5 \\\\ \\ \\ \\ \\ \\ (A0) & - & - & $\\pm.00$ & $\\pm.13$ & $\\pm$2.3 & 2 & 3.20$\\pm.17$ & 0.39 & 1.33$\\pm$.13 & 0.88$\\pm$.06 & - \\\\ HD24368 & 0.61 & 0.61 & 0.52 & 0.92 & 95.1 & 1 & 5.16$\\pm1.43$ & 0.28 & 2.22$\\pm$.31 & 0.45${<-.03 \\atop +.15}$ & 17.5 \\\\ \\ \\ \\ \\ \\ (A2V) & - & - & $\\pm.01$ & $\\pm.15$ & $\\pm$1.9 & 2 & 3.20$\\pm.17$ & 0.28 & 2.28$\\pm$.32 & 0.74$\\pm$.05 & - \\\\ \\hline \\multicolumn{12}{@{}l@{}}{\\hbox to 0pt{\\parbox{180mm}{\\footnotesize \\footnotemark[$*$] Errors of $p_{\\rm V}$ and $p_{\\rm max}$ are estimated to be 0.04\\%. See text for details. \\par\\noindent \\footnotemark[$\\dagger$] Errors of $\\lambda_{\\rm max}$, $K$, and $\\theta_{\\rm V}$ are written in the 2nd line for each object. \\par\\noindent \\footnotemark[$\\ddagger$] Deduced from the color excess by \\citet{Crawford76}, and its error is 0.02 mag. Other values of $A_{\\rm V}$ are derived from 2MASS data, having errors of 0.04 mag. \\par\\noindent }\\hss}} \\end{tabular} \\end{center} \\end{table*} \\begin{figure} \\begin{center} \\FigureFile(80mm,135mm){fig1.eps} \\end{center} \\caption{Wavelength dependence of (a) fractional polarization $p_\\lambda$ and (b) position angle $\\theta_\\lambda$ for the Pleiades stars. Solid lines in (a) show results of fitting by equation (\\ref{eq_ser}), and dashed lines in (b) the values of $\\theta_{\\rm V}$. Data of each object are moved vertically in steps of 0.25 in (a) and $10^\\circ$ in (b). }\\label{fig1} \\end{figure} Since the scattered light is often strongly polarized, it may affect the polarimetric results of nebulous objects, if it is not subtracted properly, and/or if the objects are surrounded by optically thick cloud, e.g. young stellar objects such as R Mon \\citep{MSK99}. However, the brightness of nebulosity around the stars in the Pleiades is not intensive compared with the stellar light, typically $\\sim 20$ mag/arcsec$^2$ in the B or V-band. We thus expect that the effect of the nebulosity is subtracted with the sky background in the reduction process. In addition, the spectral types of the stars are normal. Thus the observed polarization is expected to be mainly foreground interstellar in origin. Nevertheless, 19 Tau may have a component of non-interstellar origin, because the position angle is variable: $\\theta_{\\rm V}$ was $114\\degree$ in \\citet{Mark77} and in \\citet{Breger86}, but it is $140\\pm4\\degree$ in our observation (Table \\ref{tab1}). We thus exclude 19 Tau in the following discussion. \\subsection{Polarization Efficiency} The $A_{\\rm V}$-dependence of $\\lambda_{\\rm max}$ in dark clouds suggests that the alignment of grains is induced by RATs (see Section 1). % However, \\citet{AP07} noticed that the extrapolated value of $\\lambda_{\\rm max}$ at $A_{\\rm V}=0$ in each cloud is correlated with the mean value of the ratio of total to selective extinction $R_{\\rm V}$, where $R_{\\rm V} \\equiv A_{\\rm V}/E_{\\rm B-V}$ and $E_{\\rm B-V}$ is the color excess for $B-V$. Since $R_{\\rm V}$ characterizes the extinction, this correlation means that $\\lambda_{\\rm max}$ depends not only on the alignment, but also on the size of total, i.e. aligned and nonaligned, grains. We thus use another quantity less affected by the variation of grain size. Compared with $\\lambda_{\\rm max}$, the polarization efficiency $p_{\\lambda}/A_{\\lambda}$ should be less dependent on the variation of grain size, because such variation will be canceled in $p_{\\lambda}/A_{\\lambda}$. We thus explore the observed properties of $p_{\\lambda}/A_{\\lambda}$, expecting to obtain information on alignment. It should be noted, however, that \\citet{Vosh08} and \\citet{Das10} showed that $p_{\\lambda}/A_{\\lambda}$ depends on grain size, shape, material, and other parameters. It would be possible to examine the properties of $p_{\\lambda}/A_{\\lambda}$ in more detail with using light scattering calculations (e.g. Matsumura \\& Seki 1991, 1996; \\cite{MB09}), but it is beyond the scope of this Letter. To evaluate extinction $A_\\lambda$, we have used two methods: {\\it Method 1}: On the assumption that the $\\lambda$-dependence of $A_\\lambda$ is determined by $R_{\\rm V}$ and scaled by $A_{\\rm V}$ \\citep{CCM89}, we evaluate $A_\\lambda$ with interpolating the data of $A_\\lambda/E_{\\rm B-V}$ for $R_{\\rm V}=2.1\\sim5.5$ tabulated in \\citet{F04}. The values of $A_{\\rm V}$ and $R_{\\rm V}$ are calculated by the formulae $A_{\\rm V}=1.1E_{\\rm V-K}$ and $R_{\\rm V}=1.1E_{\\rm V-K}/E_{\\rm B-V}$ \\citep{WvB80}, respectively, where $E_{\\rm V-K}$ is the color excess for $V-K$. We use $B$ and $V$ magnitudes in the Simbad database, while for the $K$ band, we transform $K_{\\rm S}$ magnitude in the Two Micron All Sky Survey (2MASS) into the $K$ magnitude in the system of \\citet{Koo83} with a formula by \\citet{Carpenter01}. We refer to \\citet{Fitz70} and \\citet{Koo83} for the intrinsic colors of $B-V$ and $V-K$, respectively. Since the errors of $R_{\\rm V}$ derived with Method 1 are large for some stars (Table \\ref{tab1}), we also use another method as below (Method 2). The error of $K_{\\rm S}$ for 27 Tau was particularly large, $\\sim0.3$ mag., we could not obtain reliable results, and excluded 27 Tau in the following discussion. {\\it Method 2}: We assume that the extinction properties are not variable within the Pleiades cluster, and apply the extinction curve for HD 23512 to other stars, scaling it by the value of $A_{\\rm V}$ deduced with Method 1. The extinction curve for HD 23152 is reduced by \\citet{FM07}, and most reliable among the Pleiades stars. We explore not only the Pleiades stars, but also the stars for which polarization and extinction data are available from literature. We have used the polarimetric data of various stars by Weitenbeck (1999, 2004). The data of HD 29647 \\citep{Whittet01} and HD 38087 \\citep{Ser75} are used in addition to those from \\citet{Weiten99}. Extinction data for those stars are cited from \\citet{FM07}. Also used are the data for high latitude clouds MBM 30 and MBM 20 (LDN 1642) by \\citet{SM96}. Figure \\ref{fig2} shows the $\\lambda$-dependence of $p_\\lambda/A_\\lambda$, which is derived with Method 1. The values of log($p_\\lambda/A_\\lambda$) decrease linearly with inverse wavelength $1/\\lambda$, though slight deviations from the linear relation are found. We thus make linear fitting in $\\lambda=0.4\\sim0.8\\micron$ with the equation: \\begin{equation} \\ln(p_\\lambda/A_\\lambda) = \\ln\\alpha - \\beta(1/\\lambda - 1/0.55\\micron), \\label{eq_fit} \\end{equation} where $\\lambda$ is in $\\micron$, and $\\alpha$ and $\\beta$ are parameters and tabulated in Tables \\ref{tab1} and \\ref{tab2}. \\begin{table} \\caption{Polarization Properties deduced from Literature.}\\label{tab2} \\begin{center} \\begin{tabular}{lccccc} \\hline HD or SAO & $\\alpha$ & $\\beta$ & $T_{\\rm dust}$ & Ref.\\footnotemark[$*$] \\\\ (Sp.Type) & [\\%/mag] & $[\\micron]$ & [K] & \\\\ \\hline 14889 (K0)\\footnotemark[$\\dagger$] & 2.12$\\pm$.21 & 0.46$\\pm$.05 & 18.0 & (1) \\\\ 29647 (B8III) & 0.65$\\pm$.02 & 1.06$\\pm$.05 & 15.0 & (5) \\\\ --- & 0.62$\\pm$.03 & 1.25$\\pm$.08 & --- & (3) \\\\ 30675 (B3V) & 2.70$\\pm$.13 & 0.83$\\pm$.07 & 16.0 & (5) \\\\ 37367 (B2IV-V) & 0.74$\\pm$.04 & 0.76$\\pm$.04 & 15.8 & (4) \\\\ 38087 (B5V) & 1.47$\\pm$.07 & 0.40$\\pm$.06 & 21.0 & (3) \\\\ --- & 1.51$\\pm$.07 & 0.51$\\pm$.03 & --- & (2) \\\\ 192001 (O9.5IV) & 0.84$\\pm$.04 & 0.92$\\pm$.05 & 19.8 & (4) \\\\ 193322 (O9V) & 1.57$\\pm$.10 & 0.91$\\pm$.08 & 20.4 & (4) \\\\ 210121 (B3V) & 1.76$\\pm$.16 & 0.79$\\pm$.11 & 17.4 & (4) \\\\ 216532 (O8.5V) & 0.87$\\pm$.02 & 0.74$\\pm$.04 & 18.9 & (4) \\\\ S149760 (K5)\\footnotemark[$\\dagger$] & 2.06$\\pm$.21 & 1.03$\\pm$.13 & 16.4 & (1) \\\\ \\hline \\multicolumn{5}{@{}l@{}}{\\hbox to 0pt{\\parbox{85mm}{\\footnotesize \\footnotemark[$*$] References: (1): \\citet{SM96}, (2): \\citet{Ser75}, (3): \\citet{Weiten99}, (4): \\citet{Weiten04}, (5): \\citet{Whittet01}. The data in the R-band of HD 38087 in \\citet{Ser75} is not used for fitting. \\par\\noindent \\footnotemark[$\\dagger$] Luminosity class is assumed as V, and the values of ($R_{\\rm V}$, $A_{\\rm V}$) are estimated to be (3.08$\\pm$.20, 1.17$\\pm$.11) for HD 14889, and (2.74$\\pm$.54, 1.15$\\pm$.12) for SAO 149760. }\\hss}} \\end{tabular} \\end{center} \\end{table} \\begin{figure} \\begin{center} \\FigureFile(80mm,135mm){fig2.eps} \\end{center} \\caption{ Wavelength dependence of polarization efficiency $p_\\lambda/A_\\lambda$ deduced with Method 1 (See text). Asterisks show the results of HD 38087 \\citep{Ser75} and filled triangles HD 29647 \\citep{Whittet01}, while other symbols are for the Pleiades stars as in Figure \\ref{fig1}. Long dashed lines are the results by \\citet{Weiten99}. The data are normalized by $\\alpha$ (see equation (\\ref{eq_fit})), and moved vertically in steps of factor 1.5. }\\label{fig2} \\end{figure} \\subsection{Polarization Efficiency and Dust Temperature} To discuss the correlations between the polarization properties and dust temperature, we use the temperature $T_{\\rm dust}$ by \\citet{SFD98}. They deduced $T_{\\rm dust}$ from COBE and IRAS data, on the assumption of $\\lambda^{-2}$ emissivity of grains in the infrared. Their data is homogeneous all over the sky, and thus suitable for our study that contains not only the lines of sight to Pleiades stars, but also those to other objects. Between $\\lambda_{\\rm max}$ and $T_{\\rm dust}$, we find a weak correlation in Figure \\ref{fig3}a, and the correlation coefficient $r$ is $-0.30$. The correlations between $\\beta$ and $T_{\\rm dust}$ are much better, i.e., $r=-0.54$ (Method 1, Figure \\ref{fig3}b) and $r=-0.57$ (Method 2, Figure \\ref{fig3}c). It is remarkable that the relative positions of HD 210121 and HD 38087 are different between in Figure \\ref{fig3}a and in Figure \\ref{fig3}bc. This is caused by the different values of $R_{\\rm V}$, i.e., $R_{\\rm V}$ of HD 210121 is smaller ($=2.0$, \\cite{FM07}), and that of HD 38087 is larger ($=5.8$, \\cite{FM07}) than other objects ($R_{\\rm V}\\sim3$). \\begin{figure} \\begin{center} \\FigureFile(80mm,110mm){fig3.eps} \\end{center} \\caption{(a) Correlation between $T_{\\rm dust}$ and $\\lambda_{\\rm max}$. Filled circles are the Pleiades stars, open circles the results by Weitenbeck (1999, 2004), filled squares those by \\citet{Whittet01}, and filled triangles those by \\citet{SM96}. Double circle is the average of interstellar medium. Result of linear fitting for 14 stars is drawn as solid line. Short dashed line presents a model prediction for $a'=a_{\\rm lower}$, while long dashed line $a'=2a_{\\rm lower}$. See text for details. (b) Same as (a) but for $T_{\\rm dust}$ and $\\beta$ with Method 1. (c) Same as (b) but with Method 2. }\\label{fig3} \\end{figure} We finally discuss the above mentioned correlations on the basis of the RATs alignment theory. Using their equation (5) in \\citet{ChoLaz07} and equation (67) in \\citet{DW96}, with typical values for physical quantities in interstellar space tabulated in Table 2 of \\citet{LH07}, we can express the smallest size $a_{\\rm lower}$ of aligned grains as \\begin{equation} (a_{\\rm lower}/1\\micron) = 0.089 \\times (T_{\\rm dust}/18{\\rm K})^{-2}, \\label{eq_ChoLaz} \\end{equation} where we assume that the efficiency $Q_\\Gamma$ for RATs is 0.1, and that the emissivity of grains is $\\propto \\lambda^{-2}$. On the size distribution of aligned grains, \\citet{Mathis86} showed that the observed polarization $p_\\lambda$ can be reproduced if the fraction $(1-\\exp(-(a/a')^3))$ of grains with radius $a$ are aligned, where $a'$ is a parameter for typical size of smallest aligned grains. \\citet{Mathis86} then obtained the equation: \\begin{equation} (a'/1\\micron) = 0.327 \\times (\\lambda_{\\rm max}/1\\micron)^{2.17}. \\label{wmax_a'} \\end{equation} If we assume $a'=a_{\\rm lower}$ (or $a'=2a_{\\rm lower}$), we can relate $\\lambda_{\\rm max}$ and $T_{\\rm dust}$ with equations (\\ref{eq_ChoLaz}) and (\\ref{wmax_a'}), and draw the short dashed (or long dashed) line in Figure \\ref{fig3}a. For the relation between $\\beta$ and $\\lambda_{\\rm max}$, we obtain \\begin{equation} (\\beta/1\\micron)=1.79 \\times (\\lambda_{\\rm max}/1\\micron)^{1.39}, \\end{equation} with using equation (\\ref{eq_ser}) and extinction curve for $R_{\\rm V}=3.1$ \\citep{F04}. We then write $\\beta$ as a function of $T_{\\rm dust}$, and draw short and long dashed lines, for $a'=a_{\\rm lower}$ and $a'=2a_{\\rm lower}$, respectively, in Figures \\ref{fig3}b and \\ref{fig3}c. Those lines in Figure \\ref{fig3} seem to follow the observations well, i.e., the correlation between $\\lambda_{\\rm max}$ and $T_{\\rm dust}$, and that between $\\beta$ and $T_{\\rm dust}$ can be explained by the RATs alignment theory. Since those data contain regions of various temperature, i.e., the Taurus dark cloud, reflection nebulae, etc., our results suggest that the alignment by RATs is ubiquitous in the interstellar space. \\bigskip This work was supported by the Thesis Supporting Program at Okayama Astrophysical Observatory of NAOJ, NINS (No.08A-S01, P.I. was Youko Kameura), and by the Kagawa University Specially Promoted Research Fund (FY2008). We are grateful to the staff members at Okayama Astrophysical Observatory for their support during the observations. This work has made use of the SIMBAD database, operated at CDS, Strasbourg, France." }, "1101/1101.3651_arXiv.txt": { "abstract": "Solid-state Slit Camera (SSC) is an X-ray camera onboard the MAXI mission of the International Space Station. Two sets of SSC sensors view X-ray sky using charge-coupled devices (CCDs) in 0.5--12\\,keV band. The total area for the X-ray detection is about 200\\,cm$\\rm ^2$ which is the largest among the missions of X-ray astronomy. The energy resolution at the CCD temperature of $-$70~\\degc is 145\\,eV in full width at the half maximum (FWHM) at 5.9\\,keV, and the field of view is 1\\deg .5 (FWHM) $\\times$ 90\\deg for each sensor. The SSC could make a whole-sky image with the energy resolution good enough to resolve line emissions, and monitor the whole-sky at the energy band of $<$ 2\\,keV for the first time in these decades. ", "introduction": "\\label{sec:intro} The Monitor of All-sky X-ray Image (MAXI) mission (\\cite{matsuoka2009}) is an all-sky monitor (ASM) launched by Space Shuttle Endeavor on 2009 Jul 16. After the successful installation onto the Japanese Experiment Module - Exposed Facility (JEM-EF or Kibo EF) on the International Space Station (ISS) by astronauts using the remote manipulator system, MAXI was activated on 2009 Aug 3. The initial checkout was also successfully completed, and MAXI is in the normal observation phase from Mar 2010. In the 1-year observation, MAXI has discovered many transient phenomena in X-ray band such as X-ray bursts, Gamma-ray bursts, X-ray flashes, star flares, outbursts of recurrent nova. MAXI now occupies the essential position in monitoring astronomy. The SIS (Solid-state Imaging Spectrometer) onboard the ASCA satellite was the first X-ray photon counting CCD (charge-coupled device) camera in space (\\cite{burke1991}). Its energy resolution was good enough to resolve characteristic K-shell X-rays from heavy elements such as oxygen, silicon, and iron. The ASCA/SIS opened a new window of the X-ray astronomy. Since then, CCDs have been standard focal plane detectors of X-ray missions, such as Chandra (\\cite{weisskopf2002}), XMM-Newton (\\cite{Lumb2000}), Swift (\\cite{gehrels2004}), and Suzaku (\\cite{mitsuda2007}). For the X-ray monitoring with a large field of view, the SXC of HETE2 employed X-ray CCDs for the first time (\\cite{Villasenor2003}). MAXI has two types of X-ray camera, the Gas Slit Camera (GSC) % and the Solid-state Slit Camera (SSC). The SSC is a CCD camera covering 0.5$-$12\\,keV energy range. The energy band below 2\\,keV has not been covered by ASMs in these decades, which is achieved with MAXI/SSC. The moderate energy resolution enables us to make all-sky maps with resolved line emissions. The X-ray detection area of the SSC is unprecedentedly large to obtain the enough photon statistics. This paper describes the design and the results of the ground performance test for the SSC. We overview the SSC system in section\\,\\ref{sec:SSC_system}. Details of the sensor unit and the onboard data processing are described in section\\,\\ref{sec:SSC_units} and \\ref{sec:processing}, respectively. The results of the CCD screening and the calibration on the ground are summarized in section\\,\\ref{sec:pre_test}. In-orbit performance of the SSC is described in Tsunemi et al. (2010). ", "conclusions": "We have prepared the SSC consisting of two SSCUs. Each sensor contains the slit, collimators, and X-ray CCDs. The slit and collimators of each SSCU define the size of field of view to be 1$^{\\circ}$.5(FWHM)$\\times$90\\arcdeg . SSC has 32 CCDs of FFTCCD-4673 fabricated by Hamamatsu Photonics K.K. The SSC has total X-ray detection area of about 200\\,cm$^2$ with a sensitive X-ray energy range of 0.5$-$12keV. SSC operation has been optimized as a slit camera onboard MAXI. CCDs are operated in parallel-sum mode. The binning number of 64 is standard, but it can be changed by commands. CCDs are readout one by one in each SSCU. The charge injection is implemented to compensate the performance degradation by charged particle. The working temperature of CCDs is $<-$60\\degc in orbit by using the peltier devices, LHP, and the radiators. The energy resolution (FWHM) of the SSC for 5.9\\,keV X-rays is 145\\,eV at the CCD temperature of $-$70\\degc and 149\\,eV at $-$60\\degc. This work is partly supported by a Grant-in-Aid for Scientific Research by the Ministry of Education, Culture, Sports, Science and Technology (16002004 and 19047001). M. K. is supported by JSPS Research Fellowship for Young Scientists (22-1677)." }, "1101/1101.3184_arXiv.txt": { "abstract": "We review X-ray plasma diagnostics based on the line ratios of He-like ions. Triplet/singlet line intensities can be used to determine electronic temperature and density, and were first developed for the study of the solar corona. Since the launches of the X-ray satellites Chandra and XMM-Newton, these diagnostics have been extended and used (from \\ion{C}{v} to \\ion{Si}{xiii}) for a wide variety of astrophysical plasmas such as stellar coronae, supernova remnants, solar system objects, active galactic nuclei, and X-ray binaries. Moreover, the intensities of He-like ions can be used to determine the ionization process(es) at work, as well as the distance between the X-ray plasma and the UV emission source for example in hot stars. In the near future thanks to the next generation of X-ray satellites (e.g., Astro-H and IXO), higher-Z He-like lines (e.g., iron) will be resolved, allowing plasmas with higher temperatures and densities to be probed. Moreover, the so-called satellite lines that are formed closed to parent He-like lines, will provide additional valuable diagnostics to determine electronic temperature, ionic fraction, departure from ionization equilibrium and/or from Maxwellian electron distribution. ", "introduction": "Spectral lines of H-like and He-like ions are among the most prominent features in X-ray spectra from a large variety of astrophysical sources. Compared to other ionic iso-electronic sequences, He-like ions are abundant over the widest temperature range in collisional plasmas due to their closed shell ground state. The most intense He-like lines correspond to transitions between the $n$ = 2 shell and the $n$ = 1 ground state shell (see Fig.~\\ref{fig:grotrian}): \\begin{itemize} \\item The resonance line, named in the literature either $R$ or $w$ is an electric dipole transition (E1; 1s$^2$ $^{1}$S$_{0}$ -- 1s\\,2p $^{1}$P$_{1}$); \\item The intercombination line $I$ is composed of two lines $x$ (M2: magnetic quadrupole transition; 1s$^{2}$ $^{1}$S$_{0}$ -1s\\,2p $^{3}$P$_{2}$) and $y$ (E1; 1s$^{2}$ $^{1}$S$_{0}$ - 1s\\,2p $^{3}$P$_{1}$). The quadrupole line $x$ only becomes intense for He-like ions heavier than \\ion{S}{xv}, and with the same intensity as the line $y$ for \\ion{Ca}{xix}. The transition from $^{3}$P$_{0}$ cannot decay to the ground level since this transition is strictly forbidden but decays to the $^{3}$S$_{1}$ metastable level. \\item The forbidden line $F$ or $z$ is a relativistic magnetic dipole transition (M1; 1s$^{2}$ $^{1}$S$_{0}$ - 1s\\,2s $^{3}$S$_{1}$). \\end{itemize} \\begin{figure}[!t] \\begin{center} \\includegraphics[height=0.5\\columnwidth]{dporquet_fig1.ps} \\end{center} \\caption{Simplified level scheme for He-like ions. $w$ (or $R$), $x,y$ (or $I$), and $z$ (or $F$) are the resonance, intercombination, and forbidden lines, respectively. {\\it Upward arrows} correspond to the electron collisional excitation transitions (solid) and to the photo-excitation from 2\\,$^{3}$S$_{1}$ to 2\\,$^{3}$P$_{0,1,2}$ levels (dashed). {\\it Downward arrows} correspond to the radiative transitions (including 2-photon continuum from 2\\,$^{1}$S$_{0}$ to the ground level). {\\it Thick downward arrows}: recombination (radiative and dielectronic) plus cascade processes. Figure from \\cite{PD00}. Courtesy of Astronomy \\& Astrophysics.} \\label{fig:grotrian} \\end{figure} \\noindent The 1s\\,2s $^{1}$S$_{0}$ level decays to the ground level by a two-photon process (see section~\\ref{sec:atomic}). The energies and wavelengths of these lines for \\ion{C}{v} (Z=6), \\ion{N}{vi}, \\ion{O}{vii}, \\ion{Ne}{ix}, \\ion{Mg}{xi}, \\ion{Si}{xiii}, \\ion{S}{xv}, \\ion{Ca}{xix}, and \\ion{Fe}{xxv} (Z=26) are given in Table 1. \\\\ These He-like lines were first observed in laboratory for C, F, Mg, Al \\cite[see][]{Edlen47} and later in solar plasmas thanks to the Orbiting Solar Observatory ({\\sl OSO}), and rocket experiments \\citep[e.g.,][]{Fritz67,Doschek70,Acton72,Grineva73}. However the detection of a significant line coincident with the wavelength of a single photon transition from the metastable $^{3}$S$_{1}$ level to the ground level could not be understood from theory, because the metastable level was expected to decay only by two-photon emission. \\cite{Gabriel69a} strongly suggested that this line might correspond to some unknown photon transition. This was confirmed by a quantum-relativistic calculation of \\cite{Griem69} who proved that indeed the $^{3}$S$_{1}$ level has a significant single-photon decay rate, so a line can be observed. Then \\cite{Gabriel69b} proposed that the relative intensities of these lines can be used for temperature and density diagnostics (see section~\\ref{sec:Tne}) for solar plasma: \\\\ \\begin{equation}\\label{eq:R} {\\cal R}~(n_{\\rm e})~\\equiv~\\frac{z}{x+y} ~~~~ \\left({\\rm or} \\equiv~\\frac{F}{I}\\right) \\end{equation} \\begin{equation}\\label{eq:G} {\\cal G}~(T_{\\rm e})~\\equiv~\\frac{z+(x+y)}{w} ~~~~ \\left({\\rm or} \\equiv~\\frac{F+I}{R}\\right) \\end{equation} These diagnostics have first been widely used for solar spectra \\citep[e.g.,][]{Doschek70,Acton72,McKenzie80,Wolfson83,Keenan84,Doyle86} and for X-ray spectra of tokamak plasmas \\citep[e.g.,][]{Kallne83,Keenan89}. \\\\ \\begin{table}[!t] \\begin{center} \\begin{tabular}{cccccccccc} \\hline \\hline & \\ion{C}{v}&\\ion{N}{vi}&\\ion{O}{vii} &\\ion{Ne}{ix} &\\ion{Mg}{xi} &\\ion{Si}{xiii}&\\ion{S}{xv}&\\ion{Ca}{xix} &\\ion{Fe}{xxv} \\\\ \\hline w & {\\tiny 40.2674} & {\\tiny 28.7870} & {\\tiny 21.6015} & {\\tiny 13.4473} & {\\tiny 9.1688} & {\\tiny 6.6479} & {\\tiny 5.0387} & {\\tiny 3.1772} & {\\tiny 1.8504} \\\\ & {\\tiny (0.3079)} & {\\tiny (0.4307)} & {\\tiny (0.5740)} & {\\tiny (0.9220)} & {\\tiny (1.3522)} & {\\tiny (1.8650)} & {\\tiny (2.4606)} & {\\tiny (3.9024)} & {\\tiny (6.7004)} \\\\ x & {\\tiny 40.7280} & {\\tiny 29.0819} & {\\tiny 21.8010} & {\\tiny 13.5503} &{\\tiny 9.2282} &{\\tiny 6.6850} & {\\tiny 5.0631} & {\\tiny 3.1891} & {\\tiny 1.8554} \\\\ & {\\tiny (0.3044)} & {\\tiny (0.4263)} & {\\tiny (0.5687)} & {\\tiny (0.9150)} & {\\tiny (1.3431)} &{\\tiny (1.8547)} & {\\tiny (2.4488)} & {\\tiny (3.8878)}& {\\tiny (6.6823)} \\\\ y & {\\tiny 40.7302} & {\\tiny 29.0843} & {\\tiny 21.8036} & {\\tiny 13.5531} &{\\tiny 9.2312} &{\\tiny 6.6882} & {\\tiny 5.0665} & {\\tiny 3.1927} & {\\tiny 1.8595} \\\\ & {\\tiny (0.3044)} & {\\tiny (0.4263)} & {\\tiny (0.5686)} & {\\tiny (0.9148)} & {\\tiny (1.3431)} &{\\tiny (1.8538)} & {\\tiny (2.4471)} & {\\tiny (3.8833)}& {\\tiny (6.6676)} \\\\ z & {\\tiny 41.4715} & {\\tiny 29.5347} & {\\tiny 22.0977} & {\\tiny 13.6990} & {\\tiny 9.3143} & {\\tiny 6.7403} & {\\tiny 5.1015} & {\\tiny 3.2110} & {\\tiny 1.8682} \\\\ & {\\tiny (0.2990)} & {\\tiny (0.4198)} & {\\tiny (0.5611)} & {\\tiny (0.9051)} & {\\tiny (1.3311)} &{\\tiny (1.8394)} & {\\tiny (2.4303)}& {\\tiny (3.8612)} & {\\tiny (6.6366)} \\\\ \\hline \\hline \\end{tabular} \\end{center} \\vspace*{-0.2cm} \\caption{Wavelengths (in \\AA), and in parentheses energies (in keV) of the resonance ($w$), intercombination ($x+y$) et forbidden ($z$) lines for several He-like ions.} \\label{lambda} \\end{table} It is now possible, thanks to the spectral resolution and the sensitivity of the current generation of X-ray satellites {\\sl Chandra} and {\\sl XMM-Newton}, to resolve the He-like ion lines and to use these diagnostics for extra-solar objects. Indeed, the He-like ion line ratios are valuable tools in the analysis of high-resolution spectra of a variety of plasmas such as: \\begin{itemize} \\item Collisional Ionization Equilibrium (CIE) plasmas or also called coronal plasmas:\\\\ in such plasmas, ionization is due to electron-ion collisions and the atomic levels are populated mainly by electron impact. It is commonly assumed that CIE plasmas are optically thin to their own radiation, and that there is no external radiation field that affects the ionization balance. However, in some cases, these assumptions are not fulfilled as discussed in sections~\\ref{sec:processdiagn} and \\ref{sec:Tne}.\\\\ {\\it E.g., solar and stellar coronae (OB stars, late type stars, active stars, T Tauri, ...)}, {\\it cluster of galaxies}, {\\it the hot intra-cluster medium}, {\\it Galactic ridge} and {\\it Galactic center X-ray emission, ...} \\\\ \\item Recombination-dominated or Photo-Ionization Equilibrium (PIE) plasmas:\\\\ in such plasmas, ionization is due to photons (ionizing radiation) and the atomic levels are populated mainly by radiative recombination of H-like ions to He-like ions directly or by cascade from upper levels. These plasma are generally overionized relative to the local electronic temperature and have a much smaller electronic temperature compared to CIE plasmas. That is why collisional excitations out of the ground state are inefficient and excited levels are populated via radiative recombination. However, as we will see in section~\\ref{sec:processdiagn}, photo-excitation can be a non-negligible process. \\\\ {\\it E.g., Active galactic nuclei, X-ray binaries, ...}\\\\ \\item Out of equilibrium plasmas, or non-ionization equilibrium (NIE) plasmas: Some astrophysical plasmas depart from ionization equilibrium . This occurs when one or several physical conditions of the plasma suddenly change, such as the temperature (section~\\ref{sec:NIE}), the density, or the photo-ionization radiation field.\\\\ {\\it E.g., Supernova remnants, solar and stellar flares, colliding winds in star clusters and X-ray binaries, cluster of galaxies, intra-cluster medium in merging galaxy clusters, ...}. \\end{itemize} Since the pioneering work of \\cite{Gabriel69b}, several works have been dedicated to the improvements of these diagnostics based on He-like line ions and their extension to other types of plasmas (photo-ionization equilibrium and non-ionization equilibrium): e.g., \\cite{Mewe72}, \\cite{Blumenthal72}, \\cite{Gabriel73}, \\cite{Mewe75,Mewe78a,Mewe78b,Mewe78c}, \\cite{Acton78}, \\cite{Pradhan81}, \\cite{Bely-Dubau82}, \\cite{Pradhan82,Pradhan85}, \\cite{Keenan84}, \\cite{Swartz93}, \\cite{Liedahl99}, \\cite{PD00,P01}, \\cite{Bautista00}, \\cite{Porter07}, and \\cite{Smith09}.\\\\ In the present paper, we review why and how the relative intensities of the He-like ion lines can be used for plasma diagnostics. We also present several observational results based on these diagnostics for different types of plasmas. Since a large number of papers dealing with He-like ions has been published, this review cannot be exhaustive.\\\\ The outline of this paper is as follows. First, we give a brief overview of atomic structure and a few basic processes that play an important role for the population of the upper level of the $n$=2 shell in He-like ions (sect.~\\ref{sec:atomic}). Section~\\ref{sec:processdiagn} concerns diagnostics of the different ionization processes: CIE and PIE plasmas (sect.~\\ref{sec:CIEPIE}), NIE plasmas (sect.~\\ref{sec:NIE}), and charge-transfer (sect.~\\ref{sec:CE}). Section~\\ref{sec:Tne} is dedicated to electronic temperature (sect.~\\ref{sec:Te}) and electronic density (sect.~\\ref{sec:ne}) diagnostics. Section~\\ref{sec:satlines} presents the possible diagnostics based on the satellite lines. In the last section, we conclude and bring some possible perspectives for the future of plasma diagnostics based on He-like ions and their satellite lines thanks to the next generation of X-ray satellites, such as {\\sl Astro-H} and {\\sl IXO}. \\newpage \\section {Atomic processes and atomic data}\\label{sec:atomic} Here we briefly introduce the main atomic processes that lead to the formation of He-like ion lines, as well as the impacts of atomic data, atomic model and spectral resolution on the calculation of the He-like line ratios. For more details about X-ray spectroscopy and atomic processes, see e.g., the very nice reviews from \\cite{Liedahl99}, \\cite{Mewe99}, \\cite{Paerels03}, \\cite{Kahn05} and \\cite{Kaastra08}. \\subsection{Main atomic processes} As illustrated on the Grotrian diagram reported in Fig.~\\ref{fig:grotrian} (see also Fig.~1 in \\citealt{Mewe78a}), the atomic levels of the He-like ions can be populated and depopulated by several atomic processes. \\\\ \\noindent{\\it Collisional excitation inside He-like ions}\\\\ In most plasmas, collisional excitations are mainly due to (projectile) electrons, however collisional excitation by protons and $\\alpha$-particles can be important in some cases (see below). Electron collisional excitations from the 1s$^2$ $^{1}S_{0}$ (ground) level to excited levels ($n$=2 and higher) require a large projectile energy to open the 1s$^2$ shell. They become efficient as temperature increases and favors the population of singlet levels, such as $^{1}$P$_{1}$ level (hence the resonance $w$ line). The excitation process of singlet excited levels from the singlet ground state does not require a change of the target spin. On the contrary, excitation of triplet levels is only possible by exchange of the projectile electron with one of the target electrons. As projectile energy increases, the exchange process becomes less efficient than the direct process, i.e., the non-exchange process. For highly ionized He-like ions, the spin-orbit interaction becomes important : spin-orbit interaction between singlet and triplet levels, for example 1s\\,2p $^{3}$P$_{1}$ mixed with 1s\\,2p $^{1}$P$_{1}$. It is responsible for the similar temperature behavior for $^{1}$P$_{1}$ and $^{3}$P$_{1}$ excitations. At high temperature, radiative cascade contributions from $n >2$ levels populate the $n\\ge$ 2 levels, cascades remaining inside singlet levels or triplet levels respectively. Due to the small radiative probabilities from the $n$=2 triplet level to the singlet ground level, this favors the populations of triplet levels, namely $^{3}$P$_{0,1,2}$ and $^{3}$S$_{1}$, compared to the singlet levels which can decay more directly to the ground level. At low temperature, the contribution of the auto-ionizing resonances to the electron scattering cross-sections enhances the forbidden $z$ (and in a smaller part the intercombination ones, $x$ and $y$) transition far more than the resonance transition \\citep{Pradhan81}. Excitations inside the $n$ = 2 shell should be taken into account even for low temperature plasmas, i.e., for photo-ionized plasmas. First, the metastable $^{3}S_{1}$ level can be depopulated to the $^{3}P_{0,1,2}$ levels when the density is high enough, i.e., above the so-called critical density (that depends on the He-like ions, see section~\\ref{sec:ne}). At much higher density, the 1s\\,2s $^{1}$S$_{0}$ level (upper level of the 2-photon transition) can be depopulated to 1s\\,2p $^{1}$P$_{1}$, thereby increasing the intensity of the resonance line \\citep{Gabriel72}. The calculations of proton impact excitation rates by \\cite{Blaha71} show that their contributions (and in a smaller part those of $\\alpha$-particles) can be non-negligible for high-$Z$ (i.e., $>$14) ions at very high temperature \\citep{Mewe78a}. However, new calculations of these proton excitation rates are required (Dubau et al., in preparation). \\\\ \\noindent {\\it Recombinations from H-like ions to He-like ions}\\\\ Recombinations from H-like ions to He-like ions can be due to radiative recombination or dielectronic recombination. Radiative recombination is highly efficient at low temperature (few eV) such as in photo-ionized plasmas, and favors the populations of the triplet levels, due to their higher statistical weight compared to singlet level. On the contrary, dielectronic recombination is efficient for high temperature plasmas, but it also favors the triplet levels. It is negligible in the low temperature range (i.e., photo-ionized plasmas). Hence, both recombination processes lead to an intense forbidden or intercombination lines (depending on the density, see section~\\ref{sec:ne}) compared to the resonance line.\\\\ \\noindent {\\it Inner-shell ionization of Li-like ions}\\\\ Inner-shell ionization of Li-like ions can significantly contribute to the formation of the forbidden line of He-like ions, hence increasing the value of the \\calR\\ ratio at the low density limit and the value of the \\calG\\ ratio \\citep[e.g.,][]{Doschek70,Gabriel72,Mewe75,Mewe78a,Oelgoetz04}. For this process to have an impact on the intensity of the forbidden He-like ions, both the relative abundance of Li-like to He-like, N(Li-like)/N(He-like), and the ionization coefficient rate must be large. This latter condition is reached at high temperature. In collisional ionization equilibrium plasmas and close to the temperature of maximum formation of He-like ions, both conditions are not fulfilled since the relative abundance of Li-like ions is very small. A high abundance of Li-like ions and a high electronic temperature can occur in transient plasmas such as a in supernova remnants and during solar/stellar flares (section~\\ref{sec:NIE}), and this process is important especially for high-$Z$ ions. \\\\ \\noindent{\\it Other atomic and physical processes}\\\\ Other atomic processes should be considered in some cases such as photo-excitation (section~\\ref{sec:ne}) or charge exchange (section~\\ref{sec:CE}). For very high densities, not considered here, several atomic processes have to be taken into account \\citep[e.g.,][]{Bautista00}. In addition, resonance line scattering and optical depth might have an impact on the line ratios (section~\\ref{sec:CIEPIE}). \\subsection{Importance of the accuracy of the atomic model and atomic data} As shown by several authors, to perform line ratio calculations it is of importance to use a good atomic model with accurate atomic data \\cite[e.g.,][]{Mewe78a,Pradhan81,PD00,P01,Bautista00, Smith01,Smith09,Porter07}. \\\\ It is not possible to use a $LS$ model of He-like states to simulate the intensities of the resonance, intercombination and forbidden lines even for low charge He-like ions. Nevertheless some fine-structure data can be converted from $LS$ data, particularly collisional data due to the non-relativistic nature of the main electron-electron interaction, $1/r_{ij}$. Be aware, however it can not be used as a general rule because relativistic effects increase rapidly as the nuclear charge increases. A second point concerns the number of He-like levels included in the model, and maybe H-like and Li-like levels as well. The $n=2$ triplet levels, $^3$S$_1$, $^3$P$_{0,1,2}$ being very sensitive to recombination from H-like ions, the model must include high $n>$2 levels cascading to $n=2$ levels, even at low temperature. As temperature increases, radiative cascades contribution from collisional excitation of high $n>$2 levels can have also a significant influence on the populations of the $n=2$ levels \\citep[e.g.,][]{PD00,Smith09}. Electron ionization from Li-like ions is also possible and has a direct influence on the forbidden line intensity. Besides, proton and $\\alpha$-particle excitations might play a role on the density diagnostic for high-$Z$ He-like ions. A reliable simulation therefore requires collecting first a huge amount of accurate atomic data related to $n=2$ and also to $n >2$, up to $n=5$, or may be $n$=10. Papers giving all these data with great accuracy do not exist. Some papers contain apparently very accurate data for electron excitation but for only very few transitions. Further papers assert that the preceding calculations are not complete because they do not contain some important effects, such as radiation damping of resonances, which invalidate their accuracy. On the other hand, some authors give apparently less accurate data but including these effects. To illustrate this last point, we mention the impressive work of \\cite{Sampson83}, for electron collisional excitation of He-like excited levels for $n=1$ and $n=2$ up to $n=5$, for $4\\le Z\\le 74$. It is a Coulomb-Born-Exchange (CBE) calculation between fine-structure levels, apparently not very accurate. In particular, auto-ionizing resonances are not included. But they were inserted, including also radiation damping, in a following work by \\cite{Zhang87}. Indeed, auto-ionizing resonances have to be taken into account for a good calculation of collisional excitation rates \\citep[e.g.,][]{Pradhan81,Smith09}. But how to include them correctly ? There exists two different approaches either implicitly or explicitly. CBE, already mentioned, or Distorted Wave (DW) data do not include at all auto-ionizing resonances but it is possible to include them explicitly afterward. Whereas in Close-Coupling approximations, such as R-matrix (non-relativistic, Breit-Pauli or Dirac formalisms), they are implicitly included. These later approximations are therefore apparently better. But these resonance effects can be strong and, sometimes they are strongly overestimated because auto-ionizing resonances can also decay by radiative transitions not included in the approximations, the so-called radiation damping effect. The radiation damping is also responsible for Dielectronic Recombination, included in the model but as a recombination process (see a later section). One must therefore takes care not to include twice the same process, i.e., to separate the two decays of resonance either as excitation or as recombination. Many different methods have been proposed to overcome this problem of radiation damping in close-coupling calculations, \\cite[e.g.,][]{Zhang95}. They can give different results. What are the best ? The comparison between all these atomic data obtained by different methods is very interesting but beyond the scope of this review. Experimental measurements of the atomic data and/or line ratio of He-like ions in laboratory devices, such as tokamak, EBIT \\citep[e.g.,][]{Kallne83,Keenan89,Beiersdorfer92,Wargelin08,Beiersdorfer09,Brown09} could be of great importance to resolve some discrepancies between theoretical calculations and observations \\citep[e.g.,][]{Ness03b,Testa04a,Chen06,Smith09}. As a summary on this point, the most important is first to have a good atomic model containing the best data. \\subsection{Impact of the spectral resolution} Blending of the He-like lines with satellite lines (defined in section~\\ref{sec:satlines}) depends on the resolution of the observed spectra. Therefore, the calculations of the line ratio have to take into account contributions from unresolved satellite lines, especially for high-$Z$ ions \\cite[see e.g.,][]{P01,Sylwester08}. At low and moderate spectral resolutions all satellite lines are unresolved. In the near future higher spectral resolution will be obtained thanks to X-ray calorimeters and gratings (aboard e.g., Astro-H and IXO) at high energies and some $n$=2-3 satellites lines will be resolved for high-$Z$ ions, and can be used to probe plasma properties (see section~\\ref{sec:satlines}). However $n>3$ satellite lines will not be resolved and the \\calG\\ and \\calR\\ ratio calculations must account for their contributions. Additionally, possible contamination could be due to other elements such as \\ion{Fe}{xix} lines with the intercombination line of \\ion{Ne}{ix} lines for high enough temperature \\citep{McKenzie80,McKenzie85,Wolfson83, Ness03b}. ", "conclusions": "\\label{sec:conc} As reviewed here, diagnostics based on the main He-like ion lines can be very useful to determine the properties of astrophysical plasmas, such as the ionization processes (collisional excitation, recombination, charge transfer), departure from ionization equilibrium, electronic temperature, electronic density, and distance between the plasma where the He-like ions are formed and the UV radiation source (e.g., for stellar photosphere). The main advantages of the use of these close lines is that, for a given He-like ions, the line ratios are emitted in the same emitting volume, and are moreover insensitive to instrumental calibration, Galactic column density effects, and elemental abundances. These diagnostics were first applied to the solar corona for both active regions and flares. Thanks to the current generation of X-ray satellites, {\\sl Chandra} and {\\sl XMM-Newton}, these diagnostics are successfully used to probe the physical properties of extra-solar plasmas such as stellar coronae (from cool to hot stars), active galactic nuclei, X-ray binaries, and supernova remnants. \\\\ With the next generation of X-ray satellites, namely {\\sl Astro-H} and {\\sl IXO}, unprecedented spectral resolution combined with higher sensitivity will be reached. For {\\sl Astro-H} (JAXA, launch planned in 2014) with the Soft X-ray Spectrometer (SXS) aboard, a FWHM spectral resolution of at least 7\\,eV (with a goal to 4\\,eV) over the 0.3--12\\,keV energy range will be attained. While for {\\sl IXO} (ESA, NASA, JAXA), the planned resolving power of the grating (XGS) will be at least 3\\,000 for the soft 0.3--1\\,keV energy band, and a spectral resolution of 2.5\\,eV for the X-ray Microcalorimeter Spectrometer (XMS) over the broad 0.3--7\\,keV energy range. Therefore the resonance, intercombination and forbidden lines of He-like ions will be resolved up to \\ion{Ni}{xxvii}, including \\ion{Fe}{xxv}. Therefore higher density and/or higher temperature plasmas will be probed such as, for example, in magnetic cataclysmic variables, X-ray binaries, active galactic nuclei, and galaxy clusters. The resolution of several satellite lines will permit useful diagnostics for the determination of electronic temperature, ionic fraction, departure from ionization equilibrium and from Maxwellian electron distribution. Moreover, the combination of higher spectral resolution and sensitivity will be used to determine various physical properties of numerous astrophysical objects, including weak and high-redshift ones. Then statistics of the physical properties of a certain classes of objects (active galactic nuclei, X-ray binaries, galaxy clusters, stellar corona, supernova remnants) will be performed over luminosities, types, accretion rates, and distances. \\\\ We would like to conclude this review with the following advice from \\cite{Liedahl99}: ``In general, we need to bring to bear as many diagnostics as are available in order to make an internally consistent interpretation of an X-ray spectrum''." }, "1101/1101.1979_arXiv.txt": { "abstract": "We explore the use of the bispectrum for understanding quasiperiodic oscillations. The bispectrum is a statistic which probes the relations between the relative phases of the Fourier spectrum at different frequencies. The use of the bispectrum allows us to break the degeneracies between different models for time series which produce identical power spectra. We look at data from several observations of GRS~1915+105 when the source shows strong quasi-periodic oscillations and strong broadband noise components in its power spectrum. We show that, despite strong similarities in the power spectrum, the bispectra can differ strongly. In all cases, there are frequency ranges where the bicoherence, a measure of nonlinearity, is strong for frequencies involving the frequency of the quasi-periodic oscillations, indicating that the quasi-periodic oscillations are coupled to the noise components, rather than being generated independently. We compare the bicoherences from the data to simple models, finding some qualitative similarities. ", "introduction": "Recent studies of the variability properties of accreting black holes and neutron stars have shown that these systems' Fourier power spectra in certain spectral states are well-described by a sum of many Lorentzian components (see e.g. Olive et al. 1998; Nowak 2000; Belloni, Psaltis \\& van der Klis 2002; van Straaten et al. 2002; Pottschmidt et al. 2003). Essentially the same components seem to appear in nearly all sources, and at nearly all luminosities, and their frequencies tend to be well correlated with one another and with the source spectral states (Wijnands \\& van der Klis 1999; Psaltis, Belloni \\& van der Klis 1999). This phenomenology suggests that there may be a single origin for most of the variability features in the power spectra of accreting black holes and neutron stars. Similar correlations in frequencies between the different components seem to apply even to accreting white dwarfs (e.g. Mauche 2002; Warner, Woudt \\& Pretorius 2003). Coupling between different components in the power spectrum has been suggested in many theoretical contexts. In a shot noise model (e.g. Terrell 1972), variability components on the timescale of the shot will be correlated with one another, resulting in, e.g. fast rise exponential decay or exponential rise, fast decay profiles, but there should be no coupling on timescales longer than those of the shots. More sophisticated models of variability, for example, self-organized criticality (e.g. Takeuchi, Mineshige \\& Negoro 1995; Takeuchi \\& Mineshige 1997), predict correlations between the arrival times and/or intensities of the shots. Other related ``reservoir'' models (e.g. Merloni \\& Fabian 2001; Maccarone \\& Coppi 2002; Malzac, Merloni \\& Fabian 2004), where the accretion disk and/or the relativistic jet taps an energy supply effectively enough to reduce the available energy for future emission also predict variability correlated over many frequencies. Resonance models for producing quasi-periodic oscillations should also clearly produce non-linear variability (e.g. Psaltis \\& Norman 2001; Abramowicz \\& Kluzniak 2001; Schnittman \\& Bertschinger 2003; Maccarone \\& Schnittman 2005), while non-resonant models for producing the same QPOs (e.g. Chen \\& Taam 1992,1995; Rezzolla et al. 2003; Giannios \\& Spruit 2004) could, but need not show coupling between the different frequencies. Propagation models, where disturbances move through an accretion disc, also should produce non-linear coupling of variability components, as the variability properties are modified at each annulus (Lyubarskii 1997). A key to verifying and understanding this possible unified origin for variability is to go beyond the simple power spectrum and begin studying their non-linear variability. The first attempts at this failed to detect signatures of non-linearity, with the methods including the time skewness (Priedhorsky et al. 1978) and searches for a low dimensional chaotic attractor (Lochner, Swank \\& Szymkowiak 1989). On the other hand, more recent work with better light curves did establish that the light curves of Cygnus X-1 are not time reversible (Timmer et al. 2000; Maccarone \\& Coppi 2002), that there exists a correlation between rms amplitude and flux of a source that is inconsistent with pure shot noise models (Uttley \\& McHardy 2001), that there is coupling between variability components on all observable time scales (Maccarone \\& Coppi 2002), and that in some cases, there is a low dimensional chaotic attractor after all (Misra et al. 2004), and the observed light curves may be described by a Lorenz system (Misra et al. 2006). With this in mind, we now approach looking at the properties of the coupling between quasi-periodic components and noise components in GRS~1915+105, a bright Galactic X-ray binary with an accreting black hole which has been the subject of several long observations with the {\\it Rossi X-ray Timing Explorer \\rm(\\it{RXTE}\\rm)}. We treat this work as a pilot study -- the first attempt to apply the bispectrum to astronomical data with strong quasi-periodic oscillations. As such, we aim to illustrate the power of the technique by showing data and simple models that can give very similar power spectra, and very different bicoherences, but consider it beyond the scope of the work to try to fit models to the data precisely. We present computations of the bicoherence for several observations of GRS~1915+105, discuss the meaning of the bicoherence, and present a few toy models for the bicoherence which we compare with the data. ", "conclusions": "We have shown clearly that the variability in the noise components and quasi-periodic oscillations of GRS~1915+105 are correlated with one another. The bicoherence provides a good discriminator among variability models which produce similar power spectra through quite different physical processes. We have found that it is plausible that in the ``plateau'' state of GRS 1915+105, the variability is caused by a reservoir of energy being drained by a noise component (which could be the radio jet) and a quasi-periodic component, while in the brighter part of the $\\chi$ state, the variability is consistent with a white noise input spectrum driving a damped harmonic oscillator with a non-linear restoring force. While the models presented here are almost certainly not unique solutions for what is occuring physically in these systems, it is quite clear that the bicoherence will provide excellent constraints on more physically motivated models for the variability, and that we can definitely identify cases where the properties of the power spectra are quite similar, but the properties of the light curves are quite different." }, "1101/1101.1641.txt": { "abstract": "Causes of the unsatisfactory condition of the gravitational-wave experi\\-ments are discussed and a new outlook at the detection of gravitational waves of astrophysical origin is proposed. It is shown that there are strong grounds for identifying the so-called giant pulses in the pulsar NP 0532 radiation with gravimagnetic shock waves (GMSW) excited in the neutron star magnetosphere by sporadic gravitational radiation of this pulsar. ", "introduction": "In the history of physics of the 20th century, I suppose, there is no such a grave experimental problem (except the controlled thermonuclear fusion problem) that, being solved for over thirty years by different research groups as the gravitational wave (GW) detection problem. Although much means are used to solve it, no sufficiently convincing possitive results have been obtained. What are the reasons for this situation? An error in the gravitational theory? The experimentalists\u0092 incapability? Are there realizable opportunities to detect gravi\\-ta\\-ti\\-o\\-nal radiation in the visible future? Is the GW detection problem worth studing? We will try to answer these questions in the present paper. As any other radiation detection problem, this one also splits into two independent problems: (1) \u0093GW sources\u0094 and (2) \u0093GW detectors\u0094. It is important not to forget to join both branches by solving a concrete experimental problem. ", "conclusions": "% Besides NP 0532, among all known pulsars only PSR 0833 seems to be able to emit (but more seldom) giant pulses. Other pulsars are too old for it. Therefore it is necessary to concentrate the main effort on observations of these two pulsars. It should be stressed that there is no other mechanism able to accelerate a shock wave to subluminal velocities. Therefore an investigation of the giant pulse spectrum in the X-ray range, aimed at discovering a violet shift in the radiation spectrum, is of utmost importance. A comprehensive study of the giant pulses (their shapes and instantaneous spectrum) will allow one not only to verify the existence of gravitational radiation, but also to get additional information on the neutron stars structure and the processes in their interior. In turn it is necessary to study the GMSW pulse formation in detail theoretically. \\subsection*{Acknowledgement} The author is thankful to S.V. Sushkov for help in carring out the calculations and to N.A. Zvereva for translating the article into English." }, "1101/1101.1238_arXiv.txt": { "abstract": "The prototype of accreting, pulsating white dwarfs (GW Lib) underwent a large amplitude dwarf nova outburst in 2007. We used ultraviolet data from {\\it GALEX} and ground-based optical photometry and spectroscopy to follow GW Lib for three years following this outburst. Several variations are apparent during this interval. The optical shows a superhump modulation in the months following outburst while a 19 min quasi-periodic modulation lasting for several months is apparent in the year after outburst. A long timescale (about 4 hr) modulation first appears in the UV a year after outburst and increases in amplitude in the following years. This variation also appears in the optical 2 years after outburst but is not in phase with the UV. The pre-outburst pulsations are not yet visible after 3 years, likely indicating the white dwarf has not returned to its quiescent state. ", "introduction": "GW Librae was first thought to be a nova (Duerbeck 1987) because of the large amplitude of an outburst during its discovery in 1983, but a later spectrum at quiescence revealed it to be a dwarf nova. The lack of outbursts for the next two decades, combined with the large outburst amplitude argued for a classification as a WZ Sge type of dwarf nova (Howell, Szkody \\& Cannizzo 1995). Time-resolved spectroscopy (Szkody, Desai \\& Hoard 2000; Thorstensen et al. 2002) revealed a very short orbital period of 76.78 min, consistent with this classification. However, the unique nature of GW Lib became apparent when it was discovered to have non-radial pulsations (Warner \\& van Zyl 1998) and it became the prototype of the small class of accreting, pulsating white dwarfs which currently contains 13 members. Exploration of its amplitude spectrum by van Zyl et al. (2000, 2004) identified three prominent pulsation periods at 648, 376 and 236 s along with fine structure of closely spaced frequencies around the 648 and 376 s periods. The amplitudes of the pulsations changed on monthly timescales, typical of most hydrogen atmosphere white dwarf pulsators (ZZ Ceti stars). Ultraviolet observations of GW Lib (Szkody et al. 2002) revealed a hot white dwarf with a temperature consistent with the theoretical instability strip at 15,000K due to \\ion{He}{2} ionization (Arras et al. 2006). The best fit was achieved with a high mass (0.8 M$_{\\odot}$) white dwarf with 0.1 solar metal abundance and a dual temperature (63\\% of the surface at 13,300K and 37\\% at 17,100K). The UV data showed higher amplitude pulsations than at optical wavelengths with the same periods as evident in the optical. All stellar pulsators, including variable white dwarfs, reveal wavelength dependent amplitudes whenever their atmospheres suffer from limb-darkening effects. Recent studies of the temperatures of nine of the other accreting pulsators (Szkody et al. 2010) showed that their instability strip is wider than for ZZ Ceti stars and some stars appeared to stop pulsating during the Hubble Space Telescope observations. One possible explanation for at least one of the systems that did not show pulsations was that accretion heating from an outburst could have caused the white dwarf to move out of its instability strip. Outbursts are known to cause major heating of the white dwarfs (e.g. WZ Sge; Godon et al. 2006) which then take several years to cool to their quiescent values, even though the optical light reaches quiescence in only a few months. On 12 April 2007, GW Lib went into outburst (Templeton et al. 2007), reaching 8th magnitude from its quiescent brightness near 17th magnitude. The outburst was followed spectroscopically (Nogami et al. 2009; Hiroi et al. 2009; van Spaandonk et al. 2010a) and photometrically (Copperwheat et al. 2009, Schwieterman et al. 2010) as well as with X-ray telescopes (Byckling et al. 2009). The spectroscopy revealed the usual prominent optically thick disk near outburst with a slow return toward quiescence over 3 months. However, the outburst was unusual compared to other dwarf novae and WZ Sge objects in several ways. The X-ray flux of GW Lib increased by 3 orders of magnitude during the outburst and remained one order of magnitude higher than pre-outburst even at 2 years past outburst. As Bykling et al. (2009) point out, GW Lib is the only other dwarf nova other than U Gem that shows a larger X-ray flux during outburst than quiescence. The optical flux of GW Lib also remained about a half-magnitude higher than quiescence 2-3 years past outburst, whereas most WZ Sge stars return to optical quiescence in a few months (e.g. AL Com and WZ Sge; Szkody et al. 1998). Finally, there was no evidence of the echo outbursts that many WZ Sge stars show on the decline from outburst (e.g. EG Cnc; Patterson et al. 2002). We observed GW Lib for three years following its outburst, using ultraviolet data from {\\it GALEX}, and ground-based optical photometry and spectroscopy. Our aim was to understand the heating and cooling of GW Lib as it moved out of its instability strip and then re-entered. As of the current time, the white dwarf has not yet resumed its pre-outburst character. Yet, the photometry has revealed some interesting facets of GW Lib during its decline from superoutburst. ", "conclusions": "Our UV and optical data on GW Lib during the three years following its outburst has shown the following unique traits: \\begin{itemize} \\item The accretion disk shows large optical variability following the steep decline part of its post-outburst light curve. \\item A superhump modulation evident only in optical, not ultraviolet, follows this transition. \\item One year post-outburst, a 19 min quasi-period is visible in the optical for 2-4 months. \\item The ultraviolet begins to show variability on $\\sim$4 hour timescales a year after the outburst and the amplitude of this variability increases in the second and third years after outburst. \\item The optical begins to show a 4 hr variation 2 years after outburst which is similar to the ultraviolet but not in phase and this variation can disappear from one night to the next. \\item At one year past outburst, the white dwarf likely remained hotter than at quiescence, although a range of temperatures and disk components is possible. \\item There is no evidence of the return of the pre-outburst pulsations for the 3 years following outburst. \\end{itemize} The very large outburst amplitude, the long-lasting X-ray emission, the hot white dwarf, the lack of post-outburst rebrightenings, the presence of quasi-periodicities and the long period modulation all point to a large accretion episode, resulting in a hot white dwarf and a massive accretion disk. While the structure of the disk that results in all the periodicities observed is not known, following the detailed transitions in systems like this provide some clues to the types of phenomena that the disk undergoes as it transitions from its dominant optically thick emission state to its minor quiescent contribution to the system light. Obtaining the pulsation periods and amplitudes when the white dwarf cools enough to resume its modes will reveal the effect of the accretion event on the interior layers of the white dwarf. It is interesting to compare the behavior of GW Lib after its outburst to that of two other similar systems that also had recent outbursts and contain pulsating white dwarfs. SDSS0745+45 had an outburst in October 2006, showed no visible pulsations one year post outburst (Szkody et al. 2010) but pulsations at the same pre-outburst periods had returned by February 2010 (Mukadam et al. 2010b). V455 And had an outburst in September 2007 and its pulsation return is currently being monitored. A comparison of the outburst amplitudes and return timescales and characteristics of the pre-and post outburst pulsations of these three systems should provide some clues to the depth of heating and interaction of the accretion event with the driving zone of the pulsations." }, "1101/1101.5307_arXiv.txt": { "abstract": "In this article we present a model of formation of a galaxy with a black hole in the center. It is based on the Lema\\^{\\i}tre--Tolman solution and is a refinement of an earlier model. The most important improvement is the choice of the interior geometry of the black hole allowing for the formation of Gyrs old black holes. Other refinements are the use of an arbitrary Friedmann model as the background (unperturbed) initial state and the adaptation of the model to an arbitrary density profile of the galaxy. Our main interest was the M87 galaxy (NGC 4486), which hosts a supermassive black hole of mass $3.2\\cdot 10^{9}M_{\\odot}$. It is shown that for this particular galaxy, within the framework of our model and for the initial state being a perturbation of the $\\Lambda$CDM model, the age of the black hole can be up to $12.7$ Gyrs. The dependence of the model on the chosen parameters at the time of last scattering was also studied. The maximal age of the black hole as a function of the $\\Omega_m$ and $\\Omega_\\Lambda$ parameters for the M87 galaxy can be $3.717$ or $12.708$~Gyr. ", "introduction": "Recent years brought increasing evidence that large galaxies host massive black holes at their centers. The best evidence was obtained for the Milky Way, due to the proximity of the galactic center located at the distance $R_0 = 8.33\\pm 0.35$ kpc. Analysis of orbits of the S-stars cluster near the galactic core, especially the S2 star, allowed to conclude that a mass of $4.31\\pm 0.36\\cdot 10^{6}M_\\odot$ is responsible for the apparent movement of the stars \\cite{Gillessen:2009a,Gillessen:2009b,Gualandris:2010,Ghez:2008}. Such mass enclosed in a small volume, of radius less than $0.01$~pc, can only be a black hole. The existence of black holes at the centers of galaxies calls for construction of models describing their creation and growth. This work presents improvements to one of such models \\cite{Krasinski:2004a}. It is thought that the currently existing structures in the Universe (like galaxies, clusters of galaxies and voids) evolved out of small inhomogeneities that are observed as directional variations ($\\Delta T / T \\approx 10^{-5}$) of the temperature of the cosmic microwave background (CMB) radiation. The generation of these inhomogeneities is a problem only when one insists on using models of the Universe that are born spatially homogeneous. In an inhomogeneous model, such as the Lema\\^{\\i}tre--Tolman (LT) model used here, the Universe emerges from the Big Bang being already inhomogeneous. The inhomogeneities are generated within the Big Bang and their amplitudes are arbitrary parameters that can be adapted to observational constraints. (But it must be stressed that the LT model, in which the matter is assumed to be dust, cannot be applied to epochs earlier than last scattering. A {\\em still more general} inhomogeneous model that includes pressure gradients must be used for the pre-last-scattering epoch.) It has already been proven in an earlier paper \\cite{Krasinski:2004a}, using an LT model, that with a suitable choice of the velocity profile at last scattering, an object can be created that has the same density profile at present as the M87 galaxy, and contains a black hole around the center that has the same mass as the black hole in M87. (For an exposition of the method used in Ref. \\cite{Krasinski:2004a} see Refs. \\cite{Krasinski:2001,Krasinski:2004b} and \\cite{Plebanski:2006}.) The black hole is created because matter around the center expands more slowly than farther away, what leads, at a certain moment, to the creation of the apparent horizon surrounding a trapped region, and soon after to the creation of a Big Crunch singularity at the center. In Ref. \\cite{Krasinski:2004a} two particular configurations of the LT model were considered. In one, the black hole formed around a pre-existing wormhole within a fraction of a second after the Big Bang. In the other, there was no wormhole or black hole initially, and it formed during the evolution of the proto-galaxy. The first configuration is somewhat exotic and we will not pursue it here. The problem with the second configuration was that the implied age of the black hole was only $4 \\times 10^8$ years, what is inconsistent with astrophysical implications \\cite{2008MNRAS.391..481S,2006ApJ...643..641H,2004ApJ...614L..25Y}. The reason for the implied definite age was that the LT model used there contained too few free parameters. This paper is an improvement over Ref. \\cite{Krasinski:2004a}. The method of constructing the model is the same as before. We define a certain density or velocity profile at last scattering, taking care that the amplitude remains within the limits implied by the measurements of the temperature distribution of the CMB radiation. We define another density profile at the present time that agrees with the observationally determined density profile of the M87 galaxy \\cite{Fabricant:1980} and contains a black hole at the center, of mass $3.2\\pm 0.9\\cdot 10^{9}M_\\odot$, equal to the mass of the black hole observed in M87 \\cite{Macchetto:1997}. The improvement consists in the LT model having more free parameters. Thanks to this, the age of the black hole is no longer determined and can be up to $\\sim 12.7$ Gyr, in agreement with what astrophysics tells us \\cite{2008MNRAS.391..481S,2006ApJ...643..641H,2004ApJ...614L..25Y}.\\footnote{We do not wish to enter any discussion of the method by which the age of the black hole in M87 is inferred. We only wish to demonstrate that, whatever that inference is, our model can be made consistent with it.} Our model has certain weaknesses, of which we are aware, but, just as was stated in Ref. \\cite{Krasinski:2004a}, we treat it as an exploratory step into a new territory: we intend to test a new method that we hope will be improved in the future. The deficiencies to be removed in the future are the following: \\begin{enumerate} \\item No real galaxy is spherically symmetric. The model can be said to apply approximately to elliptic galaxies (M87 is one), but spherical symmetry leads to the next problem listed below. \\item Real galaxies rotate. In a spherically symmetric model rotation is necessarily zero. The presence of rotation influences the time scale of evolution, for example, by slowing down the collapse it may significantly delay the formation of the black hole. However, no exact solutions of Einstein's equations are known that would describe matter that is expanding and rotating at the same time. All known expanding solutions have zero rotation, all known rotating solutions are either stationary or unrealistic for other reasons (for more on this see Ref. \\cite{Plebanski:2006}). So if one wishes to have an evolving configuration that can be described by the exact formalism of general relativity, not much is left beyond spherically symmetric models. \\item As was stated in Ref. \\cite{Krasinski:2004a}, the perturbation that would evolve into a single galaxy would have the diameter of approx. 0.004$^{\\circ}$ in the CMB sky. The current best angular resolution of measurements of the temperature fluctuations of the CMB radiation is 0.2$^{\\circ}$ \\cite{WMAP7:2011a,WMAP7:2011b}. Consequently, the amplitude of fluctuation at this large angular scale does not give us information that we need to constrain our model. Lacking any better possibility, just as in the earlier paper \\cite{Krasinski:2004a}, we took care that the amplitudes of our initial density and velocity profiles do not exceed the limits set at 0.2$^{\\circ}$. \\end{enumerate} The remarks above show that our toy model cannot be literally taken as the actual model of an existing galaxy. However, it avoids a few other deficiencies that could be contemplated: \\begin{enumerate} \\item Once the black hole is formed in an LT model, where pressure is zero and all motions are radial, it keeps accreting matter until it swallows up the whole mass contained in the region where $E < 0$. This happens in a finite time, which is arbitrarily long in the neighbourhood of $E = 0$. However, if the function $E(M)$ has such a profile that at a certain $M = M_g$ $E(M_g) = 0$ and $E(M) > 0$ for $M > M_g$, then the mass from the region $M \\geq M_g$ is not accreted onto the galaxy. Thus, $M_g$, which is independent of time, can be interpreted as the mass of the galaxy. This is outside the region considered in our paper. The geometrical radius of the $M = M_g$ surface will be expanding as dictated by the $E = 0$ evolution equation of the LT model. For the surface of a galaxy\\footnote{Data for the M87 galaxy taken from \\cite{Wu:2006}, values of the constants $c$ and $G$ from http://physics.nist.gov/cuu/Constants/index.html, and the relations between distance units from http://www.asknumbers.com/LengthConversion.aspx. All values rounded off.} of radius 32 kpc $\\approx 10^{21}$ m and mass $2.4 \\times 10^{12} M_{\\odot} = 4.8 \\times 10^{42}$ kg, if it were to evolve by the $E = 0$ Lema\\^{\\i}tre -- Tolman equation ${R,_t}^2 = 2M/R$, the current velocity of expansion would be approx. 800 km/s $\\approx 8.2 \\times 10^{-4}$ pc/yr, and constantly decreasing because $\\Lambda = 0$. This is negligible compared to the error in determining the edge of a galaxy. \\item Each real galaxy is surrounded by vacuum. The cosmological model takes over at a considerably larger scale. If one wants to illustrate this situation in a toy model like ours, it is enough to match our LT galaxy model to the Schwarzschild spacetime at a certain mass $M = M_g$, and then to match the Schwarzschild region on the outside to another LT or Friedmann region modelling the Universe. Both matchings would take place outside the region we consider and would have no influence on what happens inside the galaxy. An explicit example of such matchings is given in Ref. \\cite{Matravers:2001}. \\end{enumerate} The structure of this paper is as follows. Sec.~\\ref{Sec:LT_model} presents the basic properties of the Lema\\^{\\i}tre--Tolman model. Sec.~\\ref{Sec:basic_model} briefly outlines the method introduced in Refs.~\\cite{Krasinski:2001,Krasinski:2004b} and emphasises the key elements of the model. In Sec.~\\ref{Sec:development} we describe the improvements to the basic model: the more general forms of the free functions in the LT model and the more general (and more realistic) density profile of the galaxy at the present time. This section also gives the necessary equations for the use of an arbitrary FLRW model as the background at the initial instant (in Ref. \\cite{Krasinski:2004a} the background was assumed to be spatially flat). Sec.~\\ref{Sec:M87} presents the results of application of the model to the M87 galaxy. Evolution from spatially homogeneous initial profiles of velocity and density to a galaxy with a black hole is described. A spatially homogeneous (flat) density profile is obviously within the observational limits on density perturbations, but may lead to the following problem. After the LT model is uniquely determined by the initial and final density profiles, the initial velocity profile that caused the condensation can be {\\em calculated} from the model and may turn out to have a too large amplitude. The same may happen with the initial velocity profile being flat -- the calculated amplitude of the initial density profile may turn out to be too large. Graphs shown in Sec.~\\ref{Sec:M87} prove that this did not happen; all implied profiles are consistent with the observational constraints. Sec.~\\ref{Sec:M87} also shows how the age of the black hole and the arbitrary functions in the LT model change with the initial FLRW model. The last section contains summary and conclusions. ", "conclusions": "We have demonstrated that within the inhomogeneous LT model a perturbation of the velocity or density profile at the recombination can evolve into a spherically symmetric galaxy--like object with a black hole at the center of the age of Gyr. This work is a refinement to the model presented in \\cite{Krasinski:2004a}, which allows for such an early creation of black holes, together with the usage of an arbitrary density profile describing the present day galaxy and an arbitrary FLRW model as the background matter reservoir. The evolution of such a perturbation leading to the present mass distribution of a galaxy progresses without any shell crossing singularities. In the model, the mass at the center increases in consequence of different expansion rates in the central region and farther away. The crucial point for the time of the black hole creation is the interior geometry of the black hole, that is the bang time function $t_B$ and the crunch time function $t_C(M)$ in the region $M